{
    "9803/astro-ph9803186_arXiv.txt": {
        "abstract": " ",
        "introduction": "Observations of  elemental abundances in metal-poor halo stars provide important evidence regarding the early history, evolution and age of the Galaxy. Spectroscopic studies over a number of years have indicated the presence of neutron-capture, specifically rapid neutron-capture ({\\it i.e.} {\\it r}-process), elements in a number of these metal-poor halo stars (see {\\it e.g.} Spite and Spite 1978, Sneden and Parthasarathy 1983, Sneden and Pilachowski 1985, Gilroy et al. 1988, Gratton and Sneden 1991, 1994, Sneden et al. 1994, McWilliam et al. 1995a, 1995b, Cowan et al. 1996, Sneden et al. 1996, Burris et al. 1998). In addition, abundance comparisons of the {\\it r}-process and {\\it s}-process ({\\it i.e.} slow neutron-capture) elements between the oldest metal-poor halo stars and more metal-rich halo and disk stars provide direct evidence about the nature of the chemical evolution of the Galaxy. ",
        "conclusions": "Ground-based observations of the halo stars have indicated the presence of a number of neutron-capture elements. The availability of the HST  has allowed for other spectral regions to be studied, and as a result more elements have been detected in these stars. For example, Figures 2 and 3 illustrate (see also Sneden et al. 1998) that we have now detected the element Ge in (three) halo stars. In addition, the 3$^{rd}$ {\\it r}-process peak elements, Os-Pt, as well as Pb have also now been detected in two stars using the HST (Cowan et al. 1996, Sneden et al. 1998). Including the ground-based detections of Th in CS~22892--052 (see Figure 1) and in HD 115444 (see Cowan et al. 1998), {\\it r}-process elements from  proton numbers of Z = 32 to 90 have now been observed in metal-poor stars. This is a much wider range in proton (and mass) number than ever seen before and now includes the important 3$^{rd}$ {\\it r}-process peak. These detections further demonstrate that the {\\it r}-process, ranging up to the formation of the element Th, was in operation early in the history of the Galaxy.  The observations also provide important information about the nature of the progenitors of the halo stars. The {\\it r}-process elements cannot be internally synthesized in the halo stars. Instead they must be produced in a previous generation (or generations) in an {\\it r}-process site. Since the metal-poor halo stars were formed early in the history of the Galaxy, presumably shortly after formation, the presence of {\\it r}-process elements in the halo stars requires very short evolutionary timescales for their progenitors. This further implies massive stars. While there is some uncertainty about the exact nature of the astrophysical site for the {\\it r}-process, it has long been suspected to be in supernovae, particularly core collapse supernovae from massive stars (see Cowan et al. 1991a). The abundance observations of the metal-poor halo stars appear to support that suspicion.  It has also become clear with the accumulating data that the neutron-capture elements in the metal-poor stars have an abundance pattern that appears to be the same as the solar {\\it r}-process distribution. This is illustrated vividly in Figure 1, where all of the elements from Ba to Os in CS~22892--052 have relative solar abundances. While this correlation has been noted in the past (see {\\it e.g.} Gilroy et al. 1988),  the high-resolution data in this star, covering a wide range of elements including Os in the 3$^{rd}$ {\\it r}-process peak, makes the argument much stronger. This same solar {\\it r}-process pattern also appears in other stars, as shown in Figure 2. Both ground-based and HST observations of HD~115444 ([Fe/H] = --2.7) show that the  elemental abundance from Ba to Pt are  consistent with a scaled solar {\\it r}-process curve. Further evidence of this is seen in the metal-poor halo star HD~122563 (Sneden et al. 1998). Elements such as Ba and La, which today are made predominantly in the {\\it s}-process, appear to have been made exclusively in the {\\it r}-process early in the history of the Galaxy. While this has been suggested previously (see {\\it e.g.} Truran 1981), the new observational data strongly support the contention that most (or all) of the elements were made in the {\\it r}-process at the earliest times in the Galaxy. The observations indicate the same relative {\\it r}-process abundance pattern in the oldest galactic stars and in the solar material, at least for elements with Z $\\ge$ 56. Therefore, the data indicate that the {\\it solar system {\\it r}-process abundances are not the result of global averages over different types of stars and epochs}. Instead, the stellar data suggest that the conditions that produced the {\\it r}-process elements are narrowly confined, perhaps both in terms of temperature and density of the nucleosynthesis and in terms of the mass range of the astrophysical sites (see Wheeler et al. 1998, Freiburghaus et al. 1998). The apparent lack of mixing early in the history of the Galaxy, when the relative abundance pattern is already apparent, demonstrated by Figure 6 also makes it less likely that the solar system {\\it r}-process distribution is the result of averages.  We note in Figure 1, however, that while the elements in CS~22892--052 from Ba and above (Z $\\ge$ 56) are well-fit by the solar {\\it r}-process distribution, the extrapolation to the lower mass elements does not entirely fit the abundance data. In particular, while Sr and Zr do appear to be solar, the Y abundance is far below the solar curve. It is difficult to explain why two but not three of these neighboring elements in this star are solar. We note, further, that the abundance data for HD~115444, shown in Figure 2, show a similar separation between the lighter and heavier n-capture elemental abundances. In this star, again the abundances from Ba and above appear solar, but Sr-Zr do not. There may be several possible explanations. The weak {\\it s}-process, expected to occur during helium core burning in massive stars, is expected to contribute to the abundances of the elements from Sr-Zr. The data may be showing such a contribution and Cowan et al. (1995) even suggested some combination of the weak {\\it s}-process and the {\\it r}-process might be needed to explain the abundances of Sr-Zr in CS~22892--052. We note, however, one problem in this scenario is  the apparent difficulty of producing {\\it s}-process elements in stars of extremely low metallicity. An alternative explanation may be that the more massive {\\it r}-process elements are synthesized in one site and the lower mass elements in another. Based upon meteorite data, Wasserburg et al. (1996) have suggested the existence of two {\\it r}-process sites with the separation in production occurring near mass number 140, {\\it i.e.} near Ba. Possible alternative {\\it r}-process sites have been discussed by Wheeler et al. (1998) and Baron et al. (1998). Further observations and analyses will be needed to understand the formation history of the lower mass {\\it r}-process elements.  The most metal-rich of the halo stars studied here is HD~126238, with a metallicity of [Fe/H] = --1.7. We see the same basic trends for this star in Figure 3 that are  seen in the more metal-poor stars CS~22892--052 and HD~115444. We note, however, that the abundance of Ba seems to lie above the solar {\\it r}-process curve. It was suggested by Cowan et al. (1996) that this might indicate some {\\it s}-process contribution to the Ba abundance. In other words, by the time that HD~126238 formed at a metallicity of --1.7, presumably more recently than the other two previously mentioned stars, some galactic {\\it s}-processing had occurred. It is seen in Figure 3 that the stellar Ba abundance is still below the total solar Ba abundance leading Cowan et al. to suggest that only the most massive stellar contributors to the {\\it s}-process had evolved at that point in time. Some support for their contention is given by Figure 4, which indicates the galactic chemical evolutionary trends for Ba as a function of metallicity. In this case metallicity is indicated by $\\alpha$, which may be a more reliable metallicity indicator than Fe, which is formed in both Type II and Type~I supernovae. These data spanning a wide range in metallicty might be explained by an evolutionary delay in the production of {\\it s}-process material. As demonstrated  earlier in this paper, at early times in the Galaxy Ba apparently is produced from the {\\it r}-process. We note the clear change in slope in Figure 4 at a metallicity [$\\alpha$/H] $\\simeq$ --2 that appears to indicate the onset of the main {\\it s}-process nucleosynthesis production for Ba (and presumably other n-capture elements) in the Galaxy. Further evidence of this change in production mechanism, as a function of metallicity (and presumably time), for Ba is given in Figure 5. At very low metallicities (and early times) the {\\it s}-process element Ba and the {\\it r}-process element Eu appear to be synthesized solely in the {\\it r}-process. While there is scatter in the available data, we see that at the lowest values of [$\\alpha$/H] the [Ba/Eu] value in most of the stars is consistent with a pure {\\it r}-process origin.  The long-lived radioactive nuclei (known as chronometers) in the uranium-thorium region are formed exclusively by the {\\it r}-process and can be used to determine the ages of stars and the Galaxy. One such chronometer, Th, has been detected in CS~22892--052 (Sneden et al. 1996, Cowan et al. 1997) (see Figure 1). Comparison between the initial abundance value produced in an {\\it r}-process site (often known as the production value) and the observed abundance value leads to a direct estimate of stellar ages. The abundances of the stable elements in the 3$^{rd}$ {\\it r}-process peak, a nuclear region nearby to the U-Th region, can be used to help constrain the predictions of the  initial values of the long-lived radioactive chronometers, independent of knowing the site for the {\\it r}-process.  Comparing the solar  and the observed  Th/Eu ratio in CS 22892--052, Cowan et al. (1997) found an age estimate of 15 $\\pm$ 4 Gyr for this star. They noted that consideration of galactic chemical evolution could lead to an older age of 17 $\\pm$ 4 Gyr. Pfeiffer et al. (1997) employed  newer and more accurate nuclear data in the context of a waiting point approximation {\\it r}-process model. Using the stable stellar and solar data to constrain the predicted abundances of the radioactive {\\it r}-process nuclei, they compared the initial (as opposed to solar) value of Th/Eu with the stellar value and found a best estimate for this star of 13.5 Gyr. Their result for CS~22892--052 was not only consistent with Cowan et al. (1997), but is also consistent with recent globular cluster age determinations based upon Hipparcos data (see Pont et al. 1998). (See also Cowan et al. 1991a,b for a discussion of galactic and cosmological age determinations.)  This technique, based upon predicted and observed radioactive chronometers, has been extended to an additional star with an age result approximately the same as that for CS~22892--052 (Cowan et al. 1998). We caution, however, that there are still many uncertainties, and to improve the accuracy of the chronometric estimates will require more observational and theoretical studies. It is encouraging to note, though, that the detection of thorium in CS 22892--052, and other halo stars, offers promise as an independent technique for determining stellar ages, and thus putting limits on galactic and cosmological age estimates."
    },
    "9803/gr-qc9803087_arXiv.txt": {
        "abstract": "Despite the fact that the Schwarzschild and Kerr solutions for the Einstein equations, when written in standard Schwarzschild and Boyer-Lindquist coordinates, present coordinate singularities, all numerical studies of accretion flows onto collapsed objects have been widely using them over the years. This approach introduces conceptual and practical complications in places where a smooth solution should be guaranteed, i.e., at the gravitational radius.  In the present paper, we propose an alternative way of solving the general relativistic hydrodynamic equations in background (fixed) black hole spacetimes. We identify classes of coordinates in which the (possibly rotating) black hole metric is free of coordinate singularities at the horizon, independent of time, and admits a spacelike decomposition. In the spherically symmetric, non-rotating case, we re-derive exact solutions for dust and perfect fluid accretion in Eddington-Finkelstein coordinates, and compare with numerical hydrodynamic integrations. We perform representative axisymmetric computations. These demonstrations suggest that the use of those coordinate systems carries significant improvements over the standard approach, especially for higher dimensional studies. ",
        "introduction": "\\label{sec:introduction}  It is well known that the black hole solutions of the field equations of general relativity, with the Schwarzschild and Kerr solutions being astrophysically the most relevant, exhibit coordinate singularities when written in coordinates adapted to the exterior region. The very notion of a black hole was greatly clarified with the discovery, in the early sixties, of coordinate systems that remove those singularities, and indeed cover the whole spacetime~\\cite{Kruskal60} (for a general overview see~\\cite{MTW73},\\cite{Hawking}). The (rotating) black hole solution for the metric tensor, expressed in standard Boyer-Lindquist coordinates $(t,r,\\theta,\\phi)$ is singular at the roots of the equation $\\Delta=r^2-2Mr+a^2=0$ (where $M$ is the mass and $a$ the angular momentum aspect of the hole~\\cite{units}). In the spherically symmetric Schwarzschild case, the singularity at $r=2M$ is removable with the use of appropriate transformations of the radial and time coordinates. Slightly more complicated transformations, involving also the azimuthal coordinate, can remove the singularities at $r=r_{\\pm}=M \\pm (M^2-a^2)^{1/2}$ of a rotating black hole.  The location of the coordinate singularity coincides with the event horizon of the black hole. Asymptotic observers lose causal contact with events near the black hole at precisely this location. Hence, despite the singular appearance of the metric, in simulations of matter flows in the background gravitational field of a black hole, one is in principle allowed to consider the open interval extending from the event horizon to some ``far zone\" at a large, but finite, distance from the hole.  This approach has been widely used over the years in the numerical simulations of flows around black holes~\\cite{Wi72} -\\cite{Font98a}.  The blow-up of the metric components at the horizon has implications on the behavior of hydrodynamical quantities. The coordinate flow velocity becomes ultra-relativistic and reaches the speed of light at the horizon. As a consequence, the Lorentz factor becomes infinite causing any numerical code to crash. Placing the inner boundary close to the horizon, required for capturing the effects of the relativistic potential, introduces large gradients in all hydrodynamical variables. The steep radial behavior makes the task of numerical evolution with a reasonable degree of accuracy challenging.  Besides practical considerations, the approach of evolving only the exterior domain lacks of a well defined location for the inner boundary: the computation cannot include the horizon surface, but must commence at a location ``sufficiently close'' to it. Physically, the influence of the horizon region on the solution will progressively red-shift away the closer one gets to the horizon. The validity of a certain choice though, must be continuously reasserted, as new flows or effects are being investigated, with careful tests of convergence, as the inner boundary is progressively moved inwards. Such tests are complicated by the fact that, as mentioned in the previous paragraph, the solutions appear singular at the horizon and hence demand increasingly more resolution.  Ameliorating those problems has motivated the use of a logarithmic radial coordinate (the so-called {\\it tortoise} coordinate~\\cite{RW57}). This technique relegates the event horizon to a infinite, negative, coordinate distance.  An equidistant grid in the tortoise $r_{*}$ coordinate maps into an increasingly dense grid in the Schwarzschild $r$ coordinate, with infinite density at the horizon. This approach has proven successful in the extensive semi-analytic studies of black hole perturbation theory, and recently also for the axisymmetric integration of curvature perturbations as an initial value problem~\\cite{Krivanetal}. In wave systems, the ambiguity of the location of the inner boundary is addressed by the simple limiting form of the governing equations near the horizon. This is the well known fact that black holes act as finite potential barriers to electromagnetic and gravitational perturbations~\\cite{chandra}.  The use of a tortoise coordinate does not bring similarly extensive benefits to the study of the hydrodynamical equations. The issue of the inner boundary location is less transparent for those equations. The steepness of the solution and the ``artificially'' high coordinate velocities persist, although they are now more treatable due to the substantial increase in resolution. In three dimensional simulations using Cartesian coordinates, the tortoise technique is not possible at all. It has been shown recently~\\cite{Papadopoulos98a} that {\\em adaptive mesh refinement} can indeed provide, even for 3D systems, the required resolution close to the event horizon. However, as alluded to above, these undesired pathologies can be eliminated in the case of fixed given black hole backgrounds with rather simple coordinate transformations. This reserves the power of adaptive mesh refinement for the more physically interesting features of the solution.  There is considerable freedom in choosing coordinates regular at the horizon, which can be productively reduced by imposing criteria that can enhance their suitability for numerical applications. The obvious first criterion is of course the {\\em regularity of the metric}, in particular at the horizon. A second condition is that the constant time surfaces are {\\em everywhere spacelike}, as this is, currently, a pre-requisite for the implementation of modern numerical methods for relativistic hydrodynamic flows. An important third criterion is the {\\em time-independence} of the metric components. This leads to constant in time coefficients in the equations and simplifies disentangling the true hydrodynamical evolution from coordinate effects in the black hole background. Interestingly, we will see that those conditions still do not fix the coordinate system uniquely. We show that the number of available choices greatly reduces, but is still infinite, in the rotating black hole case. We give several examples of such systems, which we collectively call {\\em horizon adapted coordinate systems}.  Such coordinate systems address the issues raised above in a straight-forward way: any radius {\\em inside} the horizon is equally appropriate (in the idealized continuum limit) for the imposition of a boundary condition, as the domain is causally disconnected from the exterior. Importantly, the irrelevance of the inner boundary location will persist even after the inclusion of other possible local physical processes that may be considered in conjuction with the hydrodynamical flow, e.g., radiative processes.  The coordinate velocities of the flow will be bounded at the horizon, as they represent projections of the (finite) fluid four velocity onto a regular coordinate system. Hence the hydrodynamical nature of the flow becomes considerably less demanding on the integration algorithm.  Gradients in the solution for {\\em scalar variables} such as pressure and enthalpy will of course persist. Those are physical and due to the curvature of the black hole, which requires a significant dynamic range for its resolution.  The organization of this paper is as follows. In section II we introduce a class of horizon adapted coordinate systems for a non-rotating black hole. The rotating black hole and the more restricted class of coordinates available in that case are discussed in the Appendix.  We outline the numerical hydrodynamical framework of our computations in section III. Exact solutions for spherical (Bondi) accretion are presented in section IV for both dust and perfect fluids. Section V describes the numerical results.  In our numerical simulations we focus mainly on the spherically symmetric case, which captures the essential nature of the problem. Some axisymmetric computations are also briefly considered. The coordinate system on which we base our computations is the celebrated Eddington-Finkelstein coordinate system~\\cite{Eddington24},\\cite{Finkel58}.  Our main aim in this report is to show the functionality of this class of coordinate systems as frameworks for the integration of the equations of relativistic hydrodynamics in black hole spacetimes.  In three-dimensions (and rotating holes) computations of accretion onto black holes are likely to benefit significantly by the adoption of a horizon adapted coordinate system. ",
        "conclusions": "We have presented a family of {\\it horizon adapted coordinate systems} for the numerical study of accretion flows around black holes. In the rotating case we identify a discrete but infinite family, of which the first simple members are well known stationary coordinate systems. In the non-rotating case the freedom in building the regular stationary foliation is quantified by (at least) the space of bounded positive functions of one variable.  We have shown how these systems allow for a better numerical treatment of accretion scenarios. Existing numerical studies of accretion flows onto black holes have been performed in the original, singular systems, i.e., Schwarzschild coordinates for the Schwarzschild solution and Boyer-Lindquist coordinates for the rotating (Kerr) solution. Although it is possible to solve the problem in these pathological coordinates, one is introducing artificial complexity, being forced to use very high resolution to deal with the unphysically large gradients that develop in the vicinity of the horizon. This may prevent the accurate solution of three dimensional problems.  At the same time, the ambiguity regarding the position of the inner boundary of the domain (which should be the horizon) introduces a convergence criterion that must be enforced at all times if the solutions are to be trusted.  We have focussed on the particular case of the Eddington-Finkelstein form of the Schwarzschild metric. The general relativistic hydrodynamic equations are now regular at the horizon, which permits an accurate description of accretion flows. We have demonstrated the feasibility of this approach with the numerical study of the spherical accretion (Bondi accretion) of dust and perfect fluid, and the comparison with the exact solutions which we re-derived in this coordinate system. We have also shown the functionality of the new coordinates in axisymmetric computations of relativistic Bondi-Hoyle accretion flows.  In a forthcoming paper we plan to extend this approach to the rotating case, considering horizon adapted coordinate systems to study accretion flows in stationary Kerr spacetimes. Three dimensional accretion flows onto black holes are interesting both from an astrophysical and geometrical point of view, as they are thought to correspond to observable electromagnetic emission, and hence may help map the relativistic rotating black hole potential. The framework proposed here will help detailed numerical studies of such systems in the near future."
    },
    "9803/astro-ph9803073_arXiv.txt": {
        "abstract": "The depletion of lithium during the pre-main sequence and main sequence phases of stellar evolution plays a crucial role in the comparison of the predictions of big bang nucleosynthesis with the abundances observed in halo stars.  Previous work has indicated a wide range of possible depletion factors, ranging from minimal in standard (non-rotating) stellar models to as much as an order of magnitude in models which include rotational mixing. Recent progress in the study of the angular momentum evolution of low mass stars (Krishnamurthi \\etal 1997a) permits the construction of theoretical models which reproduce the angular momentum evolution of low mass open cluster stars.  The distribution of initial angular momenta can be inferred from stellar rotation data in young open clusters.  In this paper we report on the application of these models to the study of lithium depletion in main sequence halo stars.   A range of initial angular momenta produces a range of lithium depletion factors on the main sequence.  Using the distribution of initial conditions inferred from young open clusters leads to a well-defined halo lithium plateau with modest scatter and a small population of outliers.  The mass dependent angular momentum loss law inferred from open cluster studies produces a nearly flat plateau, unlike previous models which exhibited a downwards curvature for hotter temperatures in the \\7li - T$_{\\rm eff}$ plane.    The overall depletion factor for the plateau stars is sensitive primarily to the solar initial angular momentum used in the calibration for the mixing diffusion coefficients.  The \\6li/\\7li depletion ratio is also examined. We find that the dispersion in the plateau and the \\6li/\\7li depletion ratio scale with the absolute \\7li depletion in the plateau and we use observational data  to set bounds on the \\7li depletion in  main sequence halo stars. A maximum of 0.4 dex  depletion is set by the observed dispersion and \\6li/\\7li depletion ratio and a minimum of 0.2 dex  depletion is required by both the presence of highly overdepleted halo stars and consistency with the solar and open cluster \\7li data.  The cosmological implications of these bounds on the primordial abundance of \\7li are discussed. ",
        "introduction": "The status of big bang nucleosynthesis (BBN) as a cornerstone of the hot big bang cosmology rests on the agreement between the theoretical predictions and the primordial abundances of the light elements deuterium (D), helium-3 (\\3he), helium-4 (\\4he), and lithium-7 (\\7li) inferred from observational data (\\cite{YTSSO}; \\cite{WSSOK}; \\cite{CSTI}).  Comparisons of this type often rely on models of galactic chemical and/or stellar evolution in order to associate the observed abundances of these light elements at the present epoch with their predicted primordial values.  Recently the confrontation between prediction and observation has come under stricter scrutiny as it appeared that the primordial abundance of deuterium inferred from ISM/Solar data was smaller than its predicted abundance - the value of which follows from requiring that the standard big bang model predictions for \\4he are in good agreement with the abundance inferred from observations of metal-poor extragalactic \\hii regions (\\cite{Hata95}; \\cite{CSTII}; \\cite{Hata97}).  Very recent measurements of the deuterium abundance along the lines-of-sight to high-redshift QSOs(\\cite{BT97}), along with a reanalysis of the \\4he data in light of new observations(\\cite{OSS96}), increase the tension between prediction and observation.  The role of \\7li in these comparisons has been minor due to the large uncertainty in the estimate of the amount of lithium destruction during the lifetimes of \\popii (\\ie, metal poor) halo stars which could accomodate primordial lithium abundances ranging from the observed plateau value\\footnote{The lithium abundances of \\popii stars with $T_{\\rm eff} > 5800$K and [Fe/H] $< -1.3$ is nearly independent of metallicity ($[Li] = 2.25\\pm 0.1$, where $[X] = 12 + \\log y_X$ with $y_X$ the number ratio of X to hydrogen) and, hence, is referred to as a ``plateau\".} (no depletion) up to a factor of ten larger. This range in possible depletion factors is equally compatible with primordial lithium abundances which correspond to either ``low deuterium'' (which favors a primordial lithium abundance a factor of 3 higher than the plateau value), or the observed \\4he (which favors the plateau value).  The arguments in favor of minimal lithium depletion are the flatness of the \\popii lithium abundance plateau (at low metallicity and high temperature) with respect to both metallicity and temperature and the low dispersion in the lithium abundance at fixed T$_{\\rm eff}$.   This was generally consistent with ``standard'' stellar models (\\ie, models without rotation which burn lithium via convection during pre-main sequence evolution) which predicted little depletion (\\cite{DDKKR}).  There are however some specific areas of disagreement in the comparison of the halo star data with standard models. The observed dispersion may be greater than that predicted (\\cite{DPD93}; \\cite{Thorburn94}; \\cite{Ryan96}) or not (\\cite{MPB95}; \\cite{Spite96}).  There is a small population of highly overdepleted stars which appear normal except for \\7li (\\cite{Thorburn94}; \\cite{NRBD}), and the trends with metal abundance appear to conflict with expectations from the models (Thorburn 1994, \\cite{CD94}).  However, the overall agreement between standard stellar evolution theory and observations in halo stars is good enough that nonstandard models would probably not be invoked to explain this data set. Therefore many investigators have adopted the reasonable assumption that the observed \\popii lithium abundances are close to the primordial value.  The overall properties of \\7li depletion in standard models have been extensively studied (for a review see \\cite{MP97}) and there are some model independent predictions:  There is partial \\7li depletion in the pre-main sequence (pre-MS), little or no main sequence (MS) depletion for stars hotter than about 5500 K, and there should be little or no dispersion in the \\7li abundance at a given T$_{\\rm eff}$ in clusters.  Unfortunately, the observational data obtained from halo stars does not serve as a good diagnostic for these models since the initial abundance is not known and we have no information on the history of the \\7li abundance. Instead, one can look to  the Sun and open clusters,  systems with a nearly uniform initial lithium abundance ([Li] = 3.2 to 3.4) where detailed abundance data is available as a function of mass, age, and composition. These \\popi data provide stringent tests of theoretical models and it has been known for quite some time (\\eg, see \\cite{WS65} and \\cite{Z72}) that standard models fail these tests.   The disagreement between the data and standard models has increased as both the observational data and the standard theoretical models have improved. The open cluster data provide clear evidence that the \\7li abundance decreases with increasing age on the MS, contrary to the standard model prediction that the convective zones of these hot ($\\geq$ 5500K) stars are not deep enough to destroy \\7li on the MS.  In addition, there is strong evidence for a dispersion in abundance at fixed T$_{\\rm eff}$ and unexpected mass-dependent effects, such as the strong depletion seen in mid-F stars relative to both hotter and cooler stars (the ``Li dip'' of \\cite{BT86} - for extensive reviews of this issue, see \\cite{PH88}, \\cite{MC91}, \\cite{S95}, \\cite{Bal95}, and \\cite{MP97}).  For standard models, these features cannot be explained by variations in the input physics (\\eg, opacities); rather, they indicate the operation of physical processes not usually included in standard models. By extension, these processes could also be operating in the \\popii stars.  Another potentially useful diagnostic of lithium depletion can be found in globular clusters.  Clusters provide samples which are homogeneous in age and composition, so one would expect a smaller dispersion in a cluster sample than in a field star sample.  This is certainly the case for \\popi stars:  \\popi field stars do show a larger dispersion than open cluster stars (\\cite{LHE91}; \\cite{FMS96}).  However, recent observations of Li in globular cluster subgiant and turnoff stars show the opposite trend and create some serious difficulties for the minimal \\7li depletion scenario for metal poor stars (\\cite{DBK95}, \\cite{BDSK97}; \\cite{TDRRO97}). There are clear star-to-star differences in M92;  out of 3 stars observed with very similar color, 2 have lithium abundances well below the plateau value.  In NGC 6397, 20 stars near the turnoff were observed.  For 7 with identical B-V color, there is a scatter of a factor of 2-3 in lithium abundance.  Therefore a standard model treatment of the halo field stars must explain why the lithium abundances in globular clusters, but not in field stars, are anomalous.  The disagreement between the \\popi lithium abundances and the predictions of standard stellar models has stimulated investigation of nonstandard stellar models.  The most prominent explanations are mixing (either from rotation or waves), microscopic diffusion, and mass loss (\\cite{MP97}; \\cite{MC91}).  For \\popi G and K stars, microscopic diffusion and mass loss are not  likely explanations (\\cite{MC91}, \\cite{SF92}), which leaves mixing as a logical candidate.  Mixing induced by rotation can explain many of the overall properties of the \\popi data.  Rotational mixing can, on relatively long timescales, mix material from the base of the convective zone to interior regions of the star where lithium can be burned; both the rotation rate and the \\7li depletion rate decrease with increased age; and a distribution of initial rotation rates can produce a distribution of \\7li depletions and thus a dispersion in abundance at fixed T$_{\\rm eff}$.  Although models with rotational mixing have the qualitative properties needed to explain the Pop I and Pop II data, two classes of objections have been raised: 1)  the uniqueness of the solutions and adequacy of the physical model; and 2) discrepancies in the quantitative comparison of observation and theory. Rotational models require an understanding of the angular momentum as a function of radius and as a function of time.  There have been persistent difficulties in reproducing the surface rotation rates as a function of mass and time in open clusters (\\cite{CDP95a},b; \\cite{KMC95}; \\cite{Bou95}). In addition, models with internal angular momentum transport from hydrodynamic mechanisms, such as the ones used in this paper, predict more rapid core rotation in the Sun than is compatible with helioseismology data (\\cite{CDP95a}, Krishnamurthi et al. 1997a (KPBS), \\cite{TST95}).  Either of these difficulties could potentially affect the degree of mixing in the models.  Models with rotational mixing predict that a range of initial angular momenta will generate a range of lithium depletion rates and therefore a dispersion in abundance among stars of the same mass, composition, and age.  However, the difficulty in reproducing the angular momentum evolution of low mass stars has made quantifying the predicted dispersion difficult, and therefore only qualitative estimates have been made (\\eg, see \\cite{PKD}(PKD), \\cite{CD94}, \\cite{CDP95b}, and \\cite{CVZ92}).  There have also been mismatches between the mean trend inferred from the data and from all classes of theoretical models. Chaboyer \\& Demarque (1994), for example, concluded that the observed mean trend of lithium abundance with T$_{\\rm eff}$ in metal-poor halo stars was not reproduced in {\\it any} class of models, be they standard, rotational, or with diffusion.  There are also cases, such as the cool stars in young open clusters, where the observed dispersion is large but the rotational models predict little, if any, dispersion.  The case for or against standard model lithium depletion is not as clear cut when we consider halo stars.  It is argued that models with significant rotationally induced depletion could not produce a flat plateau with limited dispersion about the mean plateau abundance (\\eg, \\cite{BM97}).  Furthermore, it was expected that such models would destroy far too much \\6li (\\cite{SFOSW}) in conflict with the observation of \\6li in HD 84937 (see section 4.2).  On the other hand, standard (convective burning) models show lithium depletion trends contrary to the \\popi data.  At minimum the mechanisms responsible for the \\popi \\7li pattern need to be identified and shown to not affect \\popii stars.  It is difficult to construct a model where the nonstandard effects completely cancel for \\popii stars while still retaining consistency with the \\popi data.  Models with microscopic diffusion predict modest \\7li depletion in plateau halo stars.  In general these same models predict that the timescale for changes in the surface \\7li abundance decreases with increased mass; this produces a downward curvature in the \\7li-T$_{\\rm eff}$ relationship at high temperatures which is not observed. Mass loss can counteract this trend (\\cite{Sw95}; \\cite{VC95}).  In either case, the net effect is modest depletion at the 0.2 dex level.  Rotational mixing can also suppress diffusion (Chaboyer \\& Demarque 1994, \\cite{VC95}).   The cancellation of different effects still requires at least some mean \\7li depletion in the halo stars.  In a recent preprint \\cite{VC98} note that although the surface \\7li abundance varies strongly with mass in diffusive models, the peak subsurface abundance does not (note that the height of the peak is not preserved in models which include mixing).  They then use the height of that peak to constrain the absolute \\7li abundance, arguing that the appropriate mass loss strips each star down to the region that contains this peak \\7li abundance.    The survival of \\6li is also problematic in unmixed models with sufficient mass loss to expose the peak abundance; in a model of ours which can be compared to Figure 1 in \\cite{VC98}, \\6li drops to half its surface value in the outer 0.01 solar masses, well below the comparable mass content in the \\7li preserving region of 0.02 solar masses and also the peak in the \\cite{VC98} model.  Further, it is disturbing to use a class of models which do not include a fundamental characteristic of stars, namely, rotation, when that property has been shown to be capable of affecting the issue being studied.  In our view the most serious source of uncertainty in models with rotational mixing has been the understanding of the angular momentum evolution of low mass stars.  We will show that many of the difficulties in reconciling observation and theory are resolved when models that are consistent with the rotation data in open clusters are used. In order to make definitive predictions of the amount of rotational mixing, it is necessary to follow the stellar angular momentum histories which requires knowledge of the initial distribution of rotation velocities along with an angular momentum loss law.  There has been significant recent progress in constructing theoretical models consistent with the rotation data in young open clusters (KPBS; \\cite{CL94}; \\cite{KMC95}; \\cite{BFA97}; \\cite{A97}).  With the latest generation of models, we can both infer the distribution of initial conditions and place strong constraints on the surface rotation as a function of time.  This enables us to quantify the expected dispersion in lithium abundance and greatly reduces the sensitivity of the model predictions to uncertainties in the input physics. Those models which accurately reproduce the lithium abundances observed in open clusters can then be used to predict lithium depletion in halo stars. The general trend is that these ``open cluster normalized'' rotation models for halo stars predict more lithium depletion than do the standard models but less depletion than that predicted by earlier studies of rotational mixing with a less sophisticated treatment of angular momentum evolution.  Our goal in this paper is to constrain the amount of lithium depletion in \\popii stars using several observables: (1) an estimate of the halo star initial rotation rate distribution as derived from open clusters, (2) the absence of large dispersion in the observed lithium abundances of the \\popii ``plateau'' stars, (3) the \\6li abundance and/or the \\6li/\\7li abundance ratio in HD 84937.  We argue that, individually and in combination, these observables point to \\popii halo star lithium depletion of at least 0.2 dex (following from the open cluster data) but no more than 0.4 dex (a consequence of the narrowness of the plateau and of the \\6li considerations).  Based on the observational data we can then infer the primordial lithium abundance and compare and contrast it with that predicted by standard BBN for consistency with the inferred primordial abundances of D and/or \\4he.  \\section {Method and Comparison with Previous Models}  \\subsection { Angular Momentum Evolution and Rotational Mixing}  In standard stellar models, lithium depletion is a strong function of mass and composition; it also depends on the input physics, particularly the opacities and model atmospheres used to specify the surface boundary condition. In our models the standard model physics is the same as described in KPBS (\\cite{KPBS}): namely, interior opacities from OPAL (\\cite{RI92}), low-T opacities and model atmospheres at solar abundance from \\cite{Ku91a},\\cite{Ku91b}); nuclear reaction rates from \\cite{BPW95}, including the \\cite{CF88} \\7li(p, $\\alpha$) cross-section; Yale EOS; a solar calibrated mixing length, and $Y=0.235$. Chaboyer \\& Demarque (1994) noted that unusual Li-T$_{\\rm eff}$ trends for cool metal-poor stars occurred in models using the Kurucz atmospheres and could be lessened by using a grey atmosphere.  We therefore ran halo star models with grey atmospheres and a mixing length of 1.25 as obtained for solar models constructed under the same assumptions.   Stellar models which include rotational mixing require additional input physics beyond standard stellar models.  The important new ingredients include:  \\begin{enumerate}  \\item  a distribution of initial angular momenta  \\item a prescription for angular momentum loss  \\item a prescription for the internal transport of angular momentum and the associated mixing in radiative regions  \\item the impact of rotation on the structure of the model.  \\end{enumerate} There have been important changes in the treatment of the first three ingredients since the study of rotational mixing in halo dwarfs by \\cite{PDD}; in this section we discuss the ingredients of the current set of models and compare them with previous work.  The treatment of angular momentum evolution in this set of models is the same as KPBS.  The structural effects of rotation have been computed using the method of \\cite{KT70}, and are small for low mass stars which experience angular momentum loss.  \\subsubsection {Initial Angular Momentum}  The studies of lithium depletion in the presence of rotationally induced mixing by \\cite{PKSD} (the Sun), PKD (open cluster stars), and PDD (halo stars) all used a range of initial rotation rates in the pre-MS from the measured rotation velocities of young T Tauri stars (10-60 km/s) at a typical reference age of 1 Myr.  No interaction between accretion disks and the protostar was included; the entire range of initial conditions was generated from the range in the initial rotation velocity.  This caused some difficulty in reproducing the observed range in rotation (a factor of 20 in young open clusters), especially since the higher angular momentum loss rate in rapid rotators acted to reduce the range in initial rotation rates.  This approach was also used by Chaboyer \\etal (1995a,b) for the Sun and open cluster stars respectively and by Chaboyer \\& Demarque (1994) for halo stars, although these papers used a different angular momentum loss law than the earlier work.  As protostars evolve their moment of inertia decreases; this would produce an increase in rotation rate with decreased luminosity. In contrast, the rotation periods of classical T Tauri stars appear to be nearly uniform - they do not scale with luminosity as they would if the stars were conserving angular momentum as they contracted towards the main sequence (\\cite{B93}, \\cite{E93}).  Weak-lined T Tauri stars, which are thought to be protostars without significant accretion disks, have much shorter rotation periods. This behavior would be expected if there were an accretion disk regulating the rotation of the central object and spinup only occurred when the disk is no longer magnetically coupled to the protostar (\\cite{K91}; \\cite{CCQ95}; \\cite{KMC95}).  In this revised picture there is a constant surface rotation rate while there is a sufficiently massive disk around the central object; stars which detach from their disks earlier experience a larger change in angular velocity as they contract to the main sequence than those which detach from their disk later (and are only free to begin to spin up when their moment of inertia is smaller).  This disk-locking hypothesis leads to a larger range in rotation rates for models of young open cluster stars, in better agreement with the data, and explains the very slow rotation of the majority of young stars (\\cite{BFA97}; \\cite{CCQ95}; \\cite{KMC95}; KPBS). These differences in the initial conditions lead to interesting consequences for the lithium depletion pattern.  Although the initial rotation periods are comparable to those used in the earlier studies, the inclusion of star-disk coupling implies that the MS angular momentum is much lower for models with long disk lifetimes than it would be if the disks were absent.  Disk lifetimes of 3 Myr are needed to match the rotation rates of the slow rotators in young clusters.  The mean depletion for the majority of stars is therefore significantly lower than in previous studies because most young stars are slow rotators and the degree of mixing decreases for models with lower MS angular momentum.  \\subsubsection {Angular momentum loss.}  PKSD, PKD, and PDD all used a loss law of the  form ${dJ}/{dt} \\propto \\omega^3$, where $\\omega$ is the angular velocity (\\cite{K88}). This implies greater angular momentum loss for rapid rotators than for slow rotators.  However, models of this type do not reproduce the rapid rotators seen in young open clusters; the spin down of high angular momentum objects is too severe (PKD).  On both theoretical and observational grounds, a more realistic functional form for the loss law is  \\begin {equation} \\frac{dJ}{dt} \\propto \\omega^3\\ \\ \\ \\ \\ \\ (\\omega < \\omega_{\\rm crit}), \\end {equation} and \\begin {equation} \\frac{dJ}{dt} \\propto \\omega \\omega_{\\rm crit}^2 \\ \\ \\ \\ \\ \\ (\\omega > \\omega_{\\rm crit}) \\end {equation} (\\cite{MB91}; \\cite{CCQ95}; \\cite{M84}; \\cite{BS96}). Here $\\omega_{\\rm crit}$ is the angular velocity at which the angular momentum loss rate saturates; $\\omega_{\\rm crit}$ can be estimated from several observable quantities, thought to be correlated with surface magnetic field strength, which saturate at 5-20 times the solar rotation period (see \\cite{PS96}, \\cite{Kr97b} for reviews and discussions of recent work).  A loss law which saturates at high rotation rates allows rapid rotation to survive into the early main sequence phase; a loss law of this form was adopted by Chaboyer \\etal (1995a,b) and Chaboyer \\& Demarque (1994).  The observational data indicate that the duration of the rapid rotator phase is a function of mass in the sense that rapid rotation survives longer in lower mass stars.  There is also indirect observational evidence for a mass-dependent saturation threshold (\\cite{PS96}; \\cite{Kr97b}).  Models where the saturation threshold scales with the convective overturn time scale can reproduce the observed mass dependent spin down pattern.  We adopt the mass-dependent saturation threshold of KPBS,  \\begin {equation} \\omega_{\\rm crit} = \\omega_{\\rm crit}(\\odot) \\frac {\\tau_{\\rm conv}(\\odot)} {\\tau_{\\rm conv} (*)}, \\end {equation} where the convective overturn time scale $\\tau_{\\rm conv}$ as a function of ZAMS T$_{\\rm eff}$ was taken from the 200 Myr isochrones of Kim \\& Demarque (1996; KD).  \\cite{KD96} only considered models at solar abundance, and we would need to recalibrate the KPBS models if we used a time and abundance dependent overturn timescale for the models.  A comparison of the convection zone depth (as a function of T$_{\\rm eff}$, Z, and age) in the tables of PKD and PDD indicates that at young ages the convection zone depth at a given T$_{\\rm eff}$, and therefore the overturn time scale, depends only weakly on metal abundance.  The value of $\\omega_{\\rm crit}$ is important primarily in the early main sequence.  We therefore evolved standard halo star models to an age of 200 Myr and used the same $\\tau_{\\rm conv}$ as a function of T$_{\\rm eff}$ for the metal-poor models as reported by KD for the solar abundance models.  We will consider models with a mass, composition, and time dependent overturn time scale in a future paper in preparation (\\cite{N98}). A loss law with a mass-dependent saturation threshold acts preferentially to suppress rapid rotation in the more massive stars, unlike the constant saturation threshold adopted by Chaboyer \\& Demarque (1994).  This reduces both the absolute depletion and the dispersion of lithium abundances in hotter halo plateau stars relative to cooler plateau stars.  \\subsubsection {Angular momentum transport.}  The treatment of angular momentum transport is the same as in KPBS; a detailed review of the time scale estimates can be found in Pinsonneault (1997).  We consider internal angular momentum transport by hydrodynamic mechanisms alone, and do not include potentially important mechanisms such as gravity waves (\\cite{KQ97}; \\cite{ZTM97}) and magnetic fields (\\cite{CM92}, \\cite{CM93}; \\cite{BCM98}).  In previous work the degree of mixing was found to be insensitive to the assumptions about internal angular momentum transport (PKD; PDD; \\cite{CDP95a},b).  \\cite{CVZ92} and \\cite{R96} obtained similar depletion patterns in models with solid body rotation, and \\cite{BCM98} reported similar results in models with magnetic fields and hydrodynamic mechanisms.  This can be understood because the diffusion coefficients for angular momentum transport are actually largely determined by the balance between the surface boundary condition (the applied torque) and the flux of angular momentum from below a given shell.  Theory gives an estimate of the ratio of the diffusion coefficients for mixing to those for angular momentum, which is related to the existence of anisotropic turbulence in stellar interiors (\\cite{CZ92}). The diffusion coefficients for mixing are calibrated on the Sun, but our lack of information on the rotational history of the Sun requires an exploration of models with varying solar initial conditions. ",
        "conclusions": "The lithium abundances of stars have interesting implications for cosmology and for the theory of stellar structure and evolution.  The conclusions of this paper impact both of these areas.  Standard stellar models implicitly assume that a variety of physical processes known to occur in real stars can be neglected; for many purposes this is a good approximation.  There is now extensive evidence in \\popi stars that the predictions of standard models are not in agreement with all the data.  Furthermore, work from a variety of investigators on physical effects not included in the standard models has concluded that there are physically well-motivated processes which can affect the surface lithium abundances of stars.  Standard models and the previous generation of nonstandard models predicted initial \\7li abundances that differ by up to a factor of ten from the same set of current \\popii data, with very different implications for cosmology.  We believe that mild envelope mixing driven by rotation is the most promising candidate for explaining the {\\it complete} observational picture. An improved treatment of angular momentum evolution in low mass stars now allows us to generate stellar models with rotation which are consistent with the rotation rates observed as a function of mass and age.  We have inferred the distribution of initial angular momenta in young clusters and computed the distribution of lithium depletion factors as a function of mass, composition, and age.  The resulting lithium depletion pattern is in good agreement with both \\popi and \\popii data, and the distribution of theoretical depletion factors is consistent with the distribution of abundances.  For \\popii stars in particular, the observed slope of the [Li]-T$_{\\rm eff}$ relationship, the observed dispersion in abundance at fixed T$_{\\rm eff}$, the existence of a small population of overdepleted stars, and the simultaneous detection of \\6li and \\7li in one halo star are all consistent with the predictions of the theoretical models.  The primary uncertainty in the absolute depletion of \\7li is the initial angular momentum of the Sun, which is used to calibrate the mixing diffusion coefficients.  We have used the observed properties mentioned above to set bounds on \\7li depletion for \\popii stars.  Those models which fit the \\popi data best predict a small but real depletion of lithium in the warm (T$_{\\rm eff} \\geq 5800$K), metal-poor ([Fe/H] $\\leq -1.3$) \\popii stars: 0.2 $\\leq$ logD$_{7} \\leq$ 0.4.  These same models are consistent with the very small observed dispersion in abundances about the plateau value and with the survival of some \\6li.  These depletion factors are significantly less than in previous models with rotational mixing (PDD, Chaboyer \\& Demarque 1994), and this difference can be directly attributed to the adoption of an improved treatment of angular momentum evolution. The uncertainties in the \\7li depletion factor can be reduced by a combination of new data and improved stellar evolution models.  A large set of accurate lithium abundances in old open clusters would enable us to calibrate the diffusion coefficients for mixing without relying on the solar calibration.  Microscopic diffusion and rotational mixing should be considered together, although the work of Chaboyer \\& Demarque (1994) indicated that the predictions of such models are similar to those which include rotation alone.  The observed \\7li depletion pattern is relatively insensitive to the treatment of internal angular momentum transport, with a factor of ten in the time scale corresponding to less than a 10\\% change in the logarithmic lithium depletion factor in the work of PDD.  However, the solar rotation curve does provide evidence for angular momentum transport from mechanisms not considered in the current paper, such as magnetic fields and/or internal gravity waves.  Lithium depletion in models which include these effects and hydrodynamic mechanisms should be explored.  Intermediate metallicity stars could also provide a valuable test of the dependence of lithium depletion on metal abundance.  When our constraints on lithium-depletion are combined with the abundances inferred from observations of the Spite plateau halo stars ([Li]$_{\\rm OBS} = 2.25 \\pm$ 0.10) we find for the  range in the abundance of primordial lithium: 2.35 $\\leq$ [Li]$_{\\rm P} \\leq$ 2.75 (2.2 $\\leq 10^{10}$(Li/H)$_{\\rm P} \\leq$ 5.6).  Although lithium is not the ideal baryometer, a comparison with the predictions of standard BBN identifies two options: low-$\\eta$ and high-$\\eta$.  The low-$\\eta$ branch suggests a very low baryon density, open Universe which may be in conflict with the baryon density inferred from observations of the Lyman-alpha forest.  Although the helium abundance predicted for this low-$\\eta$ option agrees with that inferred from \\hii region data, the very high predicted deuterium abundance may not be consistent with observations.  In contrast, the high-$\\eta$ branch corresponds to a baryon density consistent with the Ly-$\\alpha$ data and permits a higher density Universe.  In this case the predicted primordial deuterium abundance is in excellent agreement with the low deuterium QSO data  but a considerably higher primordial helium abundance is predicted than is inferred from some of the observational data."
    },
    "9803/astro-ph9803245_arXiv.txt": {
        "abstract": "The low-redshift evolution of the intergalactic medium is investigated using hydrodynamic cosmological simulations. The assumed cosmological model is a critical density cold dark matter universe. The imposed uniform background of ionizing radiation has the amplitude, shape and redshift evolution as computed from the observed quasar luminosity function by Haardt \\& Madau. We have analysed simulated \\lya spectra using Voigt-profile fitting, mimicking the procedure with which quasar spectra are analysed. Our simulations reproduce the observed evolution of the number of \\lya absorption lines over the whole observed interval of $z=0.5$ to $z=4$. In particular, our simulations show that the decrease in the rate of evolution of \\lya absorption lines at $z\\le 2$, as observed by {\\em Hubble Space Telescope}, can be explained by the steep decline in the photo-ionizing background resulting from the rapid decline in quasar numbers at low redshift. ",
        "introduction": "Neutral hydrogen in the intergalactic medium produces a forest of \\lya absorption lines blueward of the \\lya emission line in quasar spectra (Lynds 1971). Observations of quasars at redshifts $z>2$ show that there is strong cosmological evolution in the number of \\lya lines, which can be characterised by a power law $dN/dz \\propto (1+z)^\\gamma$ (Sargent \\etal 1980), where $N$ is the number of lines above a threshold rest-frame equivalent width $W$ (typically $W>0.32\\AA$). Studies at high resolution using the Keck telescope find $\\gamma=2.78\\pm 0.71$ for $2<z<3.5$ (\\eg Kim \\etal 1997). At even higher $z$ the evolution is still stronger: Williger \\etal (1994) find $\\gamma>4$ for $z>4$ using the CTIO 4m telescope. In contrast at low redshifts, observations using the {\\em Hubble Space Telescope} find much less evolution, $\\gamma=0.48\\pm0.62$ for $z<1$ (Morris \\etal 1991, Bahcall \\etal 1991, 1993, Impey \\etal 1996).  Recently, hydrodynamic simulations of hierarchical structure formation in a cold dark matter (CDM) dominated universe have been shown to be remarkably successful in reproducing this \\lya forest in the redshift range $z=4\\rightarrow 2$ (Cen \\etal 1994, Zhang, Anninos \\& Norman 1995, Miralda-Escud\\'e \\etal 1996, Hernquist \\etal 1996, Wadsley \\& Bond 1996, Zhang \\etal 1997, Theuns \\etal 1998). These simulations show that the weaker \\lya lines (neutral hydrogen column density $N_\\H\\le 10^{14}$ cm$^{-2}$) are predominantly produced in the filamentary and sheet-like structures that form naturally in this structure formation scenario. Velocity structure in these lines is often due to residual Hubble flow since many of the absorbing structures are expanding. In contrast, the stronger lines ($N_\\H\\ge 10^{16}$ cm$^{-2}$) tend to occur when the line of sight passes near a dense virialised halo.  Over the redshift range investigated in these simulations a photo-ionizing background close to that inferred from quasars (Haardt \\& Madau 1996) is required to explain the properties of the \\lya forest. In fact, although most simulations have assumed a critical density, scale-invariant CDM universe, other variants of the CDM model provide acceptable fits with relatively small changes to the ionizing background (Cen \\etal 1994, Miralda-Escud\\'e \\etal 1996). The general success of CDM-like models in explaining the high redshift ($z\\ge 2$) properties of the \\lya forest is impressive.  In this {\\em Letter} we investigate using hydrodynamic simulations whether a CDM universe with a photo-ionizing background dominated by quasars can explain the observed transition in the cosmological evolution of the number of \\lya lines at $z\\le 2$. ",
        "conclusions": "\\begin{figure*} \\setlength{\\unitlength}{1cm} \\centering \\begin{picture}(17,12) \\put(1, -4){\\special{psfile=fig2.ps hscale=65 vscale=65}} \\end{picture} \\caption{Evolution of the number of lines with given range in column density with redshift from simulations compared against observed evolution. Column density cut $10^{13.1}$cm$^{-2}\\le N_\\H\\le 10^{14}$ cm$^{-2}$-- simulations: circles connected with dotted line; data: open triangles (Kim \\etal 1997). Column density cut $10^{13.77}$ cm$^{-2}\\le N_\\H\\le 10^{16}$ cm$^{-2}$-- simulations: circles connected with solid line; data: filled squares (Bahcall \\etal 1993), open squares (Impey \\etal 1996), filled triangles (Kim \\etal 1997), filled pentagon (Lu \\etal 1996). Column density cut $10^{14.5}$ cm$^{-2}\\le N_\\H\\le 10^{16}$ cm$^{-2}$ -- simulation: circles connected with long dashed line; data: long dashed line shows evolution from Williger \\etal (1994). Large filled circles are simulation results for the low resolution, large box size, run, open circles are for a higher resolution, smaller box size, run. Large open pentagon: re-analysis of simulation at $z=0.5$, but imposing the ionizing background appropriate to $z=2$. Data were taken from figure~2 in Kim \\etal (1997), except for the Wiliger \\etal (1994) data. } \\label{fig:fig2} \\end{figure*} We show in figure~\\ref{fig:spectra} examples of simulated spectra at $z=3$, 2, 1 \\& 0.5 for the $L=22.22$ Mpc lower resolution simulation. Fluctuations in neutral hydrogen density, caused by gas tracing dark matter potential wells, produce absorption features similar to those in observed spectra. At low redshifts, there are large regions of the spectrum with very low absorption. These regions are separated by prominent absorption features, most of which are just a single strong line. At higher redshifts, there is considerable absorption over most of the spectrum and many lines are blended. Many of the strong lines at $z=0.5$ can be traced back to higher redshifts. There is a clear decrease in the comoving number of lines with decreasing redshift.  The evolution of the column density distribution with redshift is illustrated in figure~\\ref{fig:fig1}. There is clear evolution in the simulated column density distribution with redshift. The rate of change depends on column density, with higher column density lines undergoing stronger evolution leading to steepening of the distribution. At $z=2$, the column density distribution is $\\propto N_\\H^{-1.6}$ whereas this has steepened to $\\propto N_\\H^{-2.1}$ at $z=0.5$ (see figure~\\ref{fig:fig1}). The rate of evolution also depends on redshift, with considerably stronger evolution at higher redshifts. The number of weak \\lya lines ($\\le 10^{13.1}$ cm$^{-2}$) remains approximately constant. At redshifts 3 and 2, there is good agreement between the simulated column density distribution for our higher resolution simulation ($L=5.5$~Mpc) and the observed one.  The drop in the number of lines with redshift can be quantified by counting the number of lines within a given column density range. The simulation results are compared to observations in figure~\\ref{fig:fig2}. The simulations reproduce well the number of lines at a given redshift for all three column density cuts. Crucially, they also match very well the observed number of lines at low redshift. Consequently, the hierarchical picture of galaxy formation in a critical density universe, coupled with the observed evolution in the quasar luminosity function, can explain the observed evolution of the number of \\lya lines over the entire observed redshift range.  We have re-analysed the $z=0.5$ output time after increasing the imposed ionization flux from its $z=0.5$ value to the value appropriate to $z=2$. The number of lines with $10^{13.77}$ cm$^{-2}\\le N_\\H\\le 10^{16}$ cm$^{-2}$ is shown by the open pentagon in figure~\\ref{fig:fig2}. This point falls onto the extrapolation for the number density evolution for $z\\ge 2$. Consequently, the dominant reason for the higher number of lines at low $z$ compared to what would be expected by extrapolation from high $z$, is the decrease in ionizing flux from $z=2$ to $z=0$, itself a consequence of the evolution of the quasar luminosity function.  An estimate of the reliability of these simulations can be obtained by comparing the two simulations run at different resolutions. The higher resolution simulation produces more lines at all column densities, but the difference between the two simulations is well within the error bars of the observational results. This gives us confidence that we can reliable predict the number of lines from these simulations.  In summary: our numerical simulations show that the properties of the \\lya forest are in excellent agreement with what is expected in a cold dark matter universe with a photo-ionizing background dominated by quasar light. In particular, our simulations show that the observed decrease in the rate of evolution of \\lya absorption lines at $z\\le 2$ can be explained by the steep decline in the photo-ionizing background resulting from the rapid decline in quasar numbers at low redshift."
    },
    "9803/astro-ph9803303_arXiv.txt": {
        "abstract": "We report on the March-April 1997 BeppoSAX observations of Aql X-1, the first to monitor the evolution of the spectral and time variability properties of a neutron star soft X--ray transient from the outburst decay to quiescence. We observed a fast X--ray flux decay, which brought the source luminosity from $\\sim 10^{36}$ to $\\sim 10^{33}\\ergs$ in less than 10 days. The X--ray spectrum showed a power law high energy tail with photon index $\\Gamma\\sim 2$ which hardened to $\\Gamma \\sim 1-1.5$ as the source reached quiescence. These observations, together with the detection by RossiXTE of a periodicity of a few milliseconds during an X--ray burst, likely indicate that the rapid flux decay is caused by the onset of the propeller effect arising from the very fast rotation of the neutron star magnetosphere. The X--ray luminosity and hard spectrum that characterise the quiescent emission can be consistently interpreted as shock emission by a turned-on rotation-powered pulsar. ",
        "introduction": "Soft X--Ray Transients (SXRTs), when in outburst, show properties similar to those of persistent Low Mass X--Ray Binaries containing a neutron star (LMXRBs; White et al. 1984; Tanaka \\& Shibaza\\-ki 1996; Campana et al. 1998). The large variations in the accretion rate that are characteristic of SXRTs allow the investigation of a variety of regimes for the neutron stars in these systems which are inaccessible to persistent LMXRBs. While it is clear that, when in outbursts, SXRTs are powered by accretion, the origin of the low luminosity X--ray emission that has been detected in the quiescent state of several SXRTs is still unclear. An interesting possibility is that a millisecond radio pulsar (MSP) turns on in the quiescent state of SXRTs (Stella et al. 1994). This would provide a ``missing link\" between persistent LMXRBs and recycled MSPs. \\begin{figure*}[!htb] \\psfig{figure=aql_fig.ps,width=4.3cm} \\caption{ Light curve of the Feb.-Mar. 1997 outburst of Aql X-1 (panel {\\it a}. Data before and after MJD 50514 were collected with the RossiXTE ASM (2--10 keV) and the BeppoSAX MECS (1.5--10 keV), respectively. RossiXTE ASM count rates are converted to (unabsorbed) luminosities using a conversion factor of $4\\times 10^{35}\\ergs$ (before MJD 50512) and $2\\times10^{35}\\ergs$ (after MJD 50512) as derived from RossiXTE spectral fits (Zhang, Yu \\& Zhang 1998). BeppoSAX luminosities are derived directly from the spectral data (see text). The evolution of the flux from MJD 50480 to MJD 50512 is well fit by a Gaussian centered on MJD 50483.2. This fit however does not provide an acceptable description for later times (see the dot-dashed line), not even if the accretion luminosity is calculated in the propeller regime (dashed line). The straight solid line represents the X--ray luminosity corresponding to the closure of the centrifugal barrier $L_{\\rm min}$ (for a magnetic field of $10^8$ G and a spin period of 1.8 ms) and the straight dashed line the luminosity gap due to the action of the centrifugal barrier, $L_{\\rm cor}$. The dotted line marks the minimum luminosity in the propeller regime ($L_{\\rm lc}$). Panel {\\it b} shows the BeppoSAX unfolded spectra of Aql X-1 during the early stages of the fast decline (1) and during the quiescent phase (3--6, summed). The best fit spectral model (black body plus power law) is superposed to the data.} \\label{tot} \\end{figure*} Aql X-1 is the most active SXRT known: more than 30 X--ray and/or optical outbursts have been detected so far. The companion star has been identified with the K1V variable star V1333 Aql and an orbital period of 19 hr has been measured (Chevalier and Ilovaisky 1991). The outbursts of Aql X-1 are generally characterised by a fast rise (5--10 d) followed by a slow decay, with an $e-$folding time of 30--70~d (see Tanaka \\& Shibazaki 1996 and Campana et al. 1998 and references therein). Type I X--ray bursts were discovered during the declining phase of an outburst (Koyama et al. 1981), testifying to the presence of a neutron star. Peak X--ray luminosities are in the $\\sim (1-4)\\times 10^{37}\\ergs$ range (for the $\\sim 2.5$ kpc distance inferred from its optical counterpart; Thorstensen et al. 1978). Close to the outburst maximum the X--ray spectrum is soft with an equivalent bremsstrahlung temperature of $k\\,T_{\\rm br} \\sim 4-5$~keV. Sporadic observations of Aql X-1 during the early part of the outburst decay (Czerny et al. 1987; Tanaka \\& Shibazaki 1996; Verbunt et al. 1994) showed that when the source luminosity drops below $\\sim 10^{36}\\ergs$ the spectrum changes to a power law with a photon index of $\\Gamma\\sim 2$, extending up to energies of $\\sim 100$ keV (Harmon et al. 1996). ROSAT PSPC observations revealed Aql X-1 in quiescence on three occasions at a level of $\\sim 10^{33}\\ergs$ (0.4--2.4 keV; Verbunt et al. 1994). In this lower energy band the spectrum is considerably softer  and consistent with a black body temperature of $k\\,T_{\\rm bb} \\sim 0.3$~keV. ",
        "conclusions": "\\begin{table*} \\begin{center} \\caption{Summary of SAX NFIs observations.} \\label{tab1} \\begin{tabular}{cccccc} Obs./Date               & LECS/MECS-PDS  & LECS Count Rate              & MECS Count Rate & PDS Count Rate\\\\ & Expos. (s)     & (c s$^{-1}$)                 & (c s$^{-1}$)    & (c s$^{-1}$)\\\\ \\hline 1/March  8$^{th}$, 1997 &\\, 5240/21342   & $0.84 \\pm 0.02$              & $2.2\\pm 0.01$                 & $0.87\\pm0.06$ \\\\ 2/March 12$^{th}$, 1997 &\\, 3247/21225   & $(2.5\\pm0.4) \\times 10^{-2}$ & $(7.4\\pm 0.2) \\times 10^{-2}$ & $\\lsim 0.19$ \\\\ 3/March 17$^{th}$, 1997 &\\, 5755/17258   & $(5.6\\pm1.7) \\times 10^{-3}$ & $(1.3\\pm 0.1) \\times 10^{-2}$ & $\\lsim 0.24$ \\\\ 4/March 20$^{th}$, 1997 &\\, 4287/22589   & $(5.8\\pm1.9) \\times 10^{-3}$ & $(1.6\\pm 0.1) \\times 10^{-2}$ & $\\lsim 0.17$\\\\ 5/April 2$^{nd}$, 1997 &\\, 8440/23576     & $(6.2\\pm1.3) \\times 10^{-3}$ & $(1.2\\pm 0.1) \\times 10^{-2}$& $\\lsim 0.17$ \\\\ 6$^{\\rm o}$/May 6$^{th}$, 1997 & 11789/21703 & $(6.7\\pm1.1) \\times 10^{-3}$ &   --- & $\\lsim 0.19$ \\\\ \\hline \\end{tabular} \\end{center} {\\noindent $^{\\rm o}$ No MECS data were obtained.} \\end{table*} \\begin{table*} \\begin{center} \\caption{Summary of spectral fits.} \\label{tab2} \\begin{tabular}{ccccccc} Obs.$^*$ & Black body  & Black body  & Power law    & PL/BB     & Mean Luminosity$^\\dag$  & Reduced\\\\ & $k\\,T_{\\rm bb}$ (keV) & $R_{\\rm bb}$ (km)& $\\Gamma$ & flux ratio$^{\\ddag}$&(erg s$^{-1}$)& $\\chi^2$\\\\ \\hline 1    & $0.42\\pm0.02$&$2.6\\pm0.3$&$1.9\\pm0.1$& 3.7 & $9\\times 10^{34}$ &1.0 \\\\ 2    & $0.3_{-0.2}^{+0.1}$&$0.7_{-0.3}^{+6.6}$&$1.8\\pm0.7$& 1.6 & $2\\times 10^{33}$ &0.9 \\\\ \\ 3--6 & $0.3\\pm0.1$&$0.8_{-0.1}^{+0.4}$&$1.0\\pm0.3$& 0.7 & $6\\times 10^{32}$ &1.3 \\\\ \\hline \\end{tabular} \\end{center} { \\noindent Errors are $1\\,\\sigma$.  \\noindent $^*$ Spectra from the LECS and MECS (and PDS for the first observation) detectors have been considered. The spectra corresponding to the quiescent state have been summed up, in order to increase the statistics and an upper limit from the summed PDS data was also used to better constrain the power law slope.  \\noindent $\\dag$ Unabsorbed X--ray luminosity in the energy range 0.5--10 keV. In the case of the first observation the PDS data were included in the fit (the unabsorbed 0.5--100 keV luminosity amounts to $2\\times 10^{35}\\ergs$).  \\noindent $\\ddag$ Power law to black body flux ratio in the 0.5--10 keV energy range.}  \\end{table*} The BeppoSAX observations enabled us to follow for the first time the evolution of a SXRT outburst down to quiescence. The sharp flux decay leading to the quiescent state of Aql X-1 is reminiscent of the final evolution of dwarf novae outbursts (e.g.  Ponman et al. 1995; Osaki 1996), although there are obvious differences with respect to the X--ray luminosities and spectra involved in the two cases, likely resulting from the different efficiencies in the gravitational energy release between white dwarfs and neutron stars.  Models of low mass X--ray transient outbursts hosting an old neutron star or a black hole are largely built in analogy with dwarf novae outbursts. In particular, van Paradijs (1996) showed that the different range of time-averaged mass accretion rates over which the dwarf nova and low mass X--ray transient outbursts were observed to take place is well explained by the higher level of disk irradiation caused by the higher accretion efficiency of neutron stars and black holes. However, the outburst evolution of low mass X--ray transients presents important differences. In particular, the steepening in the X--ray flux decrease of Aql X-1 has no clear parallel in low mass X--ray transients containing Black Hole Candidates (BHCs). The best sampled light curves of these sources show an exponential-like decay (sometimes with a superposed secondary outburst) with an $e-$folding time of $\\sim 30$ d and extending up to four decades in flux, with no indication of a sudden steepening (Chen et al. 1997). In addition, BHC transients display a larger luminosity range between outburst peak and quiescence than neutron star SXRTs (Garcia et al. 1998 and references therein). Being the mass donor stars and the binary parameters quite similar in the two cases, it appears natural to attribute these differences to the different nature of the underlying object: neutron stars possess a surface and, likely, a magnetosphere, while BHCs do not.  When in outburst accretion down to the neutron star surface takes place in SXRTs, as testified by the similarity of their properties with those of persistent LMXRBs, especially the occurrence of type I bursts and the X--ray spectra and luminosities. The mass inflow rate during the outburst decay decreases, causing the expansion of the magnetospheric radius, $r_{\\rm m}$. Thus, accretion onto the neutron star surface can continue as long as the centrifugal drag exerted by the corotating magnetosphere on the accreting material is weaker than gravity (Illarionov \\& Sunyaev 1975; Stella et al. 1986). This occurs above a limiting luminosity  $L_{\\rm min}=G\\,M\\,\\mdot_{\\rm min}/R \\sim 4\\times10^{36} \\,B_8^2\\,P_{-3}^{-7/3} ~{\\rm erg\\,s^{-1}}$, where $G$ is the gravitational constant; $M$, $R$, $B=B_8\\,10^8$ G and $P=P_{-3}\\,10^{-3}$ ms are the neutron star  mass, radius, magnetic field and spin period, respectively (here and in the following we assume $M=1.4\\msole$ and $R=10^6$ cm). As $r_{\\rm m}$ reaches the corotation radius, $r_{\\rm cor}$, accretion onto the surface is inhibited and a lower accretion luminosity ($<L_{\\rm {min}}$) of $L_{\\rm cor}=G\\,M\\,\\mdot_{\\rm min}/r_{\\rm cor}\\sim 2\\times 10^{36}\\,B_8^2\\,P_{-3}^{-3}\\ergs$ is released. After this luminosity gap the source enters the propeller regime. If the mass inflow rate decreases further, the expansion of $r_{\\rm m}$ can continue up to the light cylinder radius, $r_{\\rm lc}$, providing a lower limit to the accretion luminosity that can be emitted in the propeller regime $ L_{\\rm lc}=G\\,M\\,\\mdot_{\\rm lc}/r_{\\rm lc}\\sim 2\\times 10^{34}\\,B_8^2\\,P_{-3}^{-9/2}\\ergs$. Below $L_{\\rm {lc}}$ the radio pulsar mechanism may turn on and the pulsar relativistic wind interacts with the incoming matter pushing it outwards. Matter inflowing through the Roche lobe is stopped by the radio pulsar radiation pressure, giving rise to a shock front (Illarionov \\& Sunyaev 1975; Shaham \\& Tavani 1991). Clearly these regimes have no equivalent in the case of black hole accretion. \\subsection{The onset of the propeller} During the February-March 1997 outburst of Aql X-1, RossiXTE observations led to the discovery of a nearly coherent modulation at $\\sim 550$~Hz ($\\sim 1.8$ ms) during a type I X--ray burst. A single QPO peak, with a centroid frequency ranging from $\\nu_{QPO}\\sim 750$ to 830 Hz, was also observed at two different flux levels, when the persistent luminosity was $\\sim 1.2\\times 10^{36}\\ergs$ and $\\sim 1.7\\times 10^{36}\\ergs$ (Zhang et al. 1998). In the presence of a single QPO peak, the magnetospheric and sonic point beat frequency models (Alpar \\& Shaham 1985; Miller et al. 1997) interpretation is ambiguous in that the QPO peak could represent either the Keplerian frequency at the inner disk boundary or the beat frequency. Moreover, the burst periodicity at $\\sim 550$~Hz may represent the neutron star spin frequency, $\\nu_s$, or half its value (Zhang et al. 1998). In either cases, the possibility that accretion onto the neutron star surface takes place even in the quiescent state of Aql X-1 faces serious difficulties: for a quiescent luminosity of order $10^{33}\\ergs$ a magnetic field of only $\\lsim 5\\times 10^6$ G would be required, in order to fulfill the condition $r_{\\rm m}\\lsim r_{\\rm cor}$. For such a low magnetic field, Aql~X-1 and, by inference, LMXRBs with kHz QPOs can hardly be the progenitors of recycled MSPs. More crucially, the marked steepening in the outburst decay that takes place below  $\\sim 1\\times 10^{36}\\ergs$, is accompanied by a marked spectral hardening, resulting from a sudden decrease of the flux in the black body spectral component. This is clearly suggestive of a transition taking place deep in the gravitational well of the neutron star, where most of the X--rays are produced. The most appealing mechanism is a transition to the propeller regime, where most of the inflowing matter is stopped at the magnetospheric boundary (Zhang, Yu \\& Zhang 1998). In Fig. 1a, the luminosity at MJD 50512 is identified with $L_{\\rm {min}}$ and the lower horizontal lines indicate the luminosity interval during which Aql X-1 is likely in the propeller regime.  Additional information on the neutron star magnetic field (and spin) can be inferred as follows. The observed ratio of the luminosity, $L_{QPO}$, when the QPO at $\\sim 800$~Hz were detected and the luminosity $L_{\\rm min}$ when the centrifugal barrier closes is $L_{QPO}/L_{\\rm min} \\sim 1.2-1.7$. At $L_{\\rm min}$ the Keplerian frequency of matter at the magnetospheric boundary is, by definition, equal to the spin frequency, i.e.  $\\nu_s \\sim 550$ or $\\sim 275$~Hz for Aql X-1. Based on beat-frequency models, at $L_{QPO}$ the Keplerian frequency at the inner disk boundary can be either $\\nu_{K,QPO}\\sim 800$~Hz or $\\nu_{K,QPO}\\sim \\nu_s + 800$~Hz, depending on whether the single kHz QPOs observed corresponds to the Keplerian or the beat frequency. In the magnetospheric beat-frequency models, simple theory predicts that the Keplerian frequency at the magnetospheric boundary is $\\propto L^{3/7}$; in the radiation pressure-dominated regime relevant to the case at hand, the Ghosh and Lamb (1992) model predicts instead $\\propto L^{3/13}$. Therefore we expect $\\nu_{K,QPO}/\\nu_s \\sim (L_{QPO}/L_{\\rm min})^{3/7} \\sim 1.2$ and $\\nu_{K,QPO}/\\nu_s \\sim (L_{QPO}/L_{\\rm min})^{3/13} \\sim 1.1$, in the two models, respectively. Such a low ratio clearly favors the interpretation in which $\\nu_{K,QPO}\\sim 800$~Hz and $\\nu_s\\sim 550$~Hz. In the sonic point beat-frequency model (Miller et al. 1997), the innermost disk radius is well within the magnetosphere, implying that the Keplerian frequency at the magnetospheric boundary is $< \\nu_{K,QPO}$. In this case all possible combinations of $\\nu_s$ and $\\nu_{K,QPO}$ are allowed. By using the observed $L_{\\rm min}$, a neutron star magnetic field of $B\\sim 1-3 \\times 10^8$~G (depending on the adopted model of the disk-magnetosphere interaction) is obtained in the case $\\nu_s\\sim 550$~Hz and $B\\sim 2-6 \\times 10^8$~G in the case $\\nu_s\\sim 275$~Hz.  Once in the propeller regime, the accretion efficiency decreases further as the magnetosphere expands for decreasing mass inflow rates ($L_{\\rm acc}\\propto \\mdot^{9/7}$). The $\\sim 1$~d exponential-like luminosity decline observed with BeppoSAX is considerably faster than the propeller accretion luminosity extrapolated from the first part of the outburst (e.g. the Gaussian profile shown by the dashed line in Fig. 1a). We note here that the spectral transition accompanying the onset of the centrifugal barrier may also modify the irradiation properties of the accretion disk, contributing to X--ray luminosity turn off. Alternatively, an active contribution of the ``propeller'' mechanism or the neutron star spin-down energy dissipated into the inflowing matter cannot be excluded. \\subsection{A turned-on rotation-powered pulsar?} It is unlikely that the quiescent luminosity of Aql~X-1 is powered by magnetospheric accretion in the propeller regime. As shown in Fig. 1a, the quiescent X--ray luminosity is probably lower than the minimum magnetospheric accretion luminosity $L_{\\rm lc}$ allowed in the propeller phase (this remains true for $B\\gsim 6\\times 10^7$ G if $\\nu_s\\sim 550$ Hz, and for $B\\gsim 3\\times 10^8$ G if $\\nu_s\\sim 275$ Hz). Moreover the BeppoSAX X--ray spectrum shows a pronounced decrease in the power law to black body flux ratio together with a flattening of the power law component between the fast decay phase and quiescence, suggesting that a transition to shock emission from the interaction of a radio pulsar wind with the matter outflowing from the companion star has taken place. Note that an  X--ray spectrum with a slope of $\\Gamma \\sim 1.5$ has been observed from the radio pulsar PSR~B1259--63 immersed in the wind of its Be star companion. Models of this interaction predict that a power law X--ray spectrum with a slope around $\\Gamma \\sim 1.5$ should be produced for a wide range of parameters (Tavani \\& Arons 1997).  The additional soft X--ray component observed during the outburst decay (see Table 2) might be emitted at the polar caps as a result of the residual neutron star accretion in the propeller phase. Note that the equivalent black body radius decreases for decreasing X--ray luminosities, just as it would be expected if the magnetospheric boundary expanded. Alternatively, the black body-like spectral component observed in quiescence could be due to the stre\\-aming of energetic particles that hit the polar caps, in close analogy to the soft X--ray component observed, in MSPs, at the weaker level of $\\sim 10^{30}-10^{31}\\ergs$ (Becker \\& Tr\\\"umper 1997).  Assuming a magnetic field in the range derived in section 3.1 (i.e. $B\\sim 1-3\\times10^8$ G for $\\nu_s=550$ Hz and $B\\sim 2-6\\times10^8$ G for $\\nu_s=275$ Hz), we can consistently explain the $\\sim 10^{33}\\ergs$ quiescent X--ray luminosity as powered by a radio pulsar enshrouded by matter outflowing from the companion star, if the conversion efficiency of spin-down luminosity to X--ray is $\\sim 0.1-10$\\%. This is consistent with modeling and observations of enshrouded pulsars (Tavani 1991; Verbunt et al. 1996).  There are chances of observing a MSP (a simple scaling from MSPs implies a signal at 400 MHz of $\\sim 10$ mJy; see Kulkarni et al. 1990), even though the emission would probably be sporadic, like in the case of the pulsar PSR~1744--24A due to the large amount of circumstellar matter (see Lyne et al. 1991; Shaham \\& Tavani 1991).  In summary Aql~X-1 appears to provide the first example of an old fast rotating neutron stars undergoing a transition to the propeller regime at first, followed by a transition to the radio pulsar regime, as the transient X--ray emission approaches its quiescent level. Therefore, Aql X-1 (and possibly SXRTs in general) likely represents the long-sought ``missing link'' between LMXRBs and recycled MSPs."
    },
    "9803/hep-ph9803293_arXiv.txt": {
        "abstract": "The M-theory, the strong-coupling heterotic string theory, presents various interesting new phenomenologies. The M-theory bulk axion is one of these. The decay constant in this context is estimated as $F_a\\simeq 10^{16}$~GeV. Direct searches for the M-theory axion seem impossible because of the large decay constant. However, we point out that large isocurvature fluctuations of the M-theory axion are obtained in a hybrid inflation model, which will most likely be detectable in future satellite experiments on anisotropies of cosmic microwave background radiation. ",
        "introduction": "It is widely believed that nonperturbative effects play a central role in describing the real world in superstring theories. Horava and Witten~\\cite{Horava-Witten} have argued that the strong-coupling heterotic string theory is dual to an M-theory compactified on ${\\bf S}^1/{\\bf Z}_2$. The low-energy limit of this M-theory is well described by eleven-dimensional supergravity, where the fundamental energy scale is the eleven-dimensional Planck mass $M_{11}\\simeq 6\\times 10^{16}$~GeV, rather than the four-dimensional one $M_4\\simeq 2\\times 10^{18}$~GeV~\\cite{Witten}. The four-dimensional Planck mass $M_4$ is only an effective parameter appearing at low energies.  This M-theory description of strong-coupling heterotic string theory leads to various interesting new phenomenologies. Banks and Dine~\\cite{Banks-Dine-1} have pointed out that the bulk moduli fields provide axion candidates in the M-theory compactified further on a Calabi-Yau manifold, and some of the axions survive at low energies, since world-sheet instanton effects are suppressed owing to the large compactification radius in the string tension unit.  One of the string axions may acquire its mass dominantly from the QCD anomaly, and it plays the role of the Peccei-Quinn axion~\\cite{Peccei-Quinn} in solving the strong $CP$ problem.  The decay constant $F_a$ of the M-theory axion is estimated as~\\cite{Banks-Dine-1,Choi} \\begin{equation} F_a\\simeq 10^{16} {\\rm GeV}. \\end{equation} This value greatly violates the constraint $10^{10} {\\rm GeV}\\lesssim F_a\\lesssim 10^{12}$~GeV, derived in standard cosmology~\\cite{Kolb-Turner}. However, this problem may be solved by late-time entropy production through decays of moduli fields~\\cite{Kawasaki-Moroi-Yanagida-1,Banks-Dine-2} or a thermal inflaton field~\\cite{Lyth-Stewart}. Since a very low reheating temperature, such as $T_R\\simeq 1$ - $10$MeV, is required to solve the above problem, it is hard to imagine a consistent production of the lightest supersymmetric (SUSY) particles, and they cannot be the dark matter~\\cite{Kawasaki-Moroi-Yanagida-2}. Thus, the M-theory axion seems the most plausible candidate for the dark matter in our universe.  If the M-theory axion is indeed the dark matter, fluctuations of the axion density should consist of mixtures of adiabatic and isocurvature modes in general~\\cite{Turner,Lyth,ksy}. In this paper we show that a hybrid inflation model proposed by Linde and Riotto~\\cite{Linde-Riotto} naturally produces isocurvature fluctuations of the axion comparable to the adiabatic fluctuations that are in the accessible range of future satellite experiments on anisotropies of the cosmic microwave background radiation (CMB). The present analysis, therefore, confirms the previous proposal~\\cite{Kawasaki-Yanagida}. ",
        "conclusions": "The mixture of isocurvature and adiabatic fluctuations is astrophysically interesting. Since isocurvature fluctuations yield anisotropies of the CMB that are six times larger than those caused by adiabatic fluctuations, mixed fluctuations reduce the amplitude of the power spectrum if the amplitude is normalized by COBE.  It is well known that the standard cold dark matter scenario ($\\Omega_0=1, h=0.5$) with COBE-normalized pure adiabatic fluctuations predicts density fluctuations that are too large on scales of galaxies and clusters. This problem is avoided if the isocurvature fluctuations are mixed with adiabatic ones, as is pointed in Ref.~\\cite{ksy}. Furthermore, it can be shown that for the general flat universe ($\\lambda_0 +\\Omega_0=1$, with $\\lambda_0$ the cosmological constant), the shape and amplitude of the power spectrum are in good agreement with observations if $\\alpha \\sim 0.05$~\\cite{KKSY}.  The CMB anisotropies induced by the isocurvature fluctuations can be distinguished from those produced by pure adiabatic fluctuations~\\cite{ksy}, because the shapes of the angular power spectrum of CMB anisotropies are quite different from each other on small angular scales. The most significant effect of the mixture of the isocurvature fluctuation is that the acoustic peak in the angular power spectrum decreases.  In Fig.~\\ref{fig:mix_sample} we show the angular power spectrum for $\\alpha=0.05$, $\\Omega_0=0.4$ and $h=0.7$ as an example. It is seen that the height of the acoustic peak ($\\ell \\sim 200$) in the case of mixed fluctuations is greatly reduced compared with the pure adiabatic case.  Since the axion decay constant $F_a$ is much higher than $10^{12}$~GeV, a direct search for the M-theory axion is implausible. Therefore, observations of CMB anisotropies by future satellite experiments (MAP~\\cite{MAP}, PLANCK~\\cite{PLANCK}) are very crucial to test the M-theory axion hypothesis."
    },
    "9803/astro-ph9803317_arXiv.txt": {
        "abstract": "s{ We report the development of numerical tools for the topological analysis of sub--degree resolution, all--sky maps. Software to be released in the HEALFAST (V0.9) package defines neighbour relationships for the HEALPIX tessellation of the sphere. We apply this routine to a fast extrema search which scales strictly linearly in the number of pixels, $N_{p}$. We also present a highly efficient algorithm for simulating the gradient vector and curvature tensor fields ``on--the--fly'' with the temperature map, needing only of order $N_{p} \\log_2 N_{p}$ more operations.} ",
        "introduction": "\\label{sec:intro}  The Microwave Anisotropy Probe (MAP) and the Planck Surveyor satellite missions promise to produce high--resolution full sky maps of the Cosmic Microwave Background (CMB) within the next decade. These maps will contain a great deal of cosmologically relevant information and therefore present a tremendous opportunity for doing high--precision cosmology. Owing to the unprecedented wealth of data, this opportunity brings with it the challenge of developing tools which help extract this information in manageable processing time.  Much attention has been given to the problem of estimating the angular power--spectrum of CMB anisotropies. To illustrate the difficulties one is facing, we note that even choosing the fastest presently available algorithm and allowing for parallelisation and technological advance, the estimation of the angular power--spectrum $C_l$ of CMB anisotropies will take $\\gtrsim 6$ years \\cite{borrill} for future million pixel maps.  While the $C_l$ spectrum is the statistic of choice for a large class of cosmological models where the primordial fluctuations are Gaussian distributed, the possibility remains that another physical mechanism produced structure in the universe. In this case the CMB maps will almost certainly be non--Gaussian and contain more information than merely the power on different angular scales. Failing this, non--linear gravitational effects and foregrounds will produce detectable non--Gaussian signals.  Even for Gaussian maps, it has been shown that extrema statistics capture information about cosmological parameters and their theory is well developed \\cite{sazhin,BE,VJ,BSMCS}. The importance of studying these and other topological properties of CMB temperature and polarisation has been noted by several authors \\cite{topos}. Extrema have the advantage that they dominate the noise and are therefore very robust. As a consequence they have been among the first scientific results reported from recent CMB observations (one of them even at this conference \\cite{spot}!).  The aim of this talk is to present tools which allow the implementation of a wide range of data analysis and processing techniques on the sphere. \\newpage ",
        "conclusions": "In this talk we presented tools which are generally applicable for the analysis of topological properties of maps on the sphere such as the detection of extrema, saddle points and zero crossings of a function, even though they can be used more generally for any kind of local analysis of spherical data sets. The computational time required for all sky computations scales strictly as the number of pixels in the map, $N_{p}$. Further, we extend existing simulation tools for Gaussian fields to generate not only temperature maps but also the associated gradient vector and the curvature tensor fields, with the extra computational cost scaling as $N_{p} \\log_2 N_{p}$."
    },
    "9803/astro-ph9803251_arXiv.txt": {
        "abstract": "\\rightskip = 0pt \\noindent We suggest that the high--velocity clouds (HVCs) are large clouds, with typical diameters of 25 kpc and containing $5 \\times 10^7$ solar masses of neutral gas and $3 \\times 10^8$ solar masses of dark matter, falling onto the Local Group; altogether the HVCs contain 10$^{11}$ solar masses of neutral gas.  Our reexamination of the Local--Group hypothesis for the HVCs connects their properties to the hierarchical structure formation scenario and to the gas seen in absorbtion towards quasars.  We begin by showing that at least one HVC complex (besides the Magellanic Stream) must be extragalactic at a distance ${>} 40$ kpc from the Galactic center, with a diameter ${>} 20$ kpc and a mass ${>} 10^8$ solar masses. We then discuss a number of other clouds that are positionally associated with the Local Group galaxies. The kinematics of the entire ensemble of HVCs is inconsistent with a Galactic origin, but implies that the HVCs are falling towards the Local Group barycenter. The HVCs obey an angular--size/velocity relation consistent with the Local Group infall model.  We simulate the dynamical evolution of the Local Group. The simulated properties of material falling into the Local Group reproduce the location of two of the three most significant groupings of clouds and the kinematics of the entire cloud ensemble (excluding the Magellanic Stream). We interpret the third grouping (the A, C, and M complexes) as tidally unstable nearby material falling onto the Galactic disk. We interpret the more distant HVCs as dark matter ``mini--halos'' moving along filaments towards the Local Group.  Most poor galaxy groups should contain HI structures to large distances bound to the group.  We suggest that the HVCs are local analogues of the Lyman--limit absorbing clouds observed against distant quasars.  We argue that there is a Galactic fountain in the Milky Way, but that the fountain does not explain the origin of the HVCs. Our analysis of the HI data leads to the detection of a vertical infall of low--velocity gas towards the plane.  We suggest that the fountain is a local phenomenon involving neutral gas with characteristic velocities of 6 \\kmse rather than 100 \\kms. This implies that the chemical evolution of the Galactic disk is governed by episodic infall of metal-poor HVC gas that only slowly mixes with the rest of the interstellar medium.  The Local--Group infall hypothesis makes a number of testable predictions.  The HVCs should have sub-solar metallicities. Their H$\\alpha $ emission should be less than that seen from the Magellanic Stream. The clouds should not be seen in absorption to nearby stars. The clouds should be detectable in both emission and absorption around other groups. We show that current observations are consistent with these predictions and discuss future tests. ",
        "introduction": "\\label{sec:intro}  Since their discovery in 1963 by \\markcite{Muller63}Muller, Oort, \\& Raimond, the nature of the high--velocity hydrogen clouds (HVCs) has remained a mystery.  HVCs are clouds that deviate from Galactic circular rotation by as much as several hundred kilometers per second.  Although astronomers have speculated about the origin of HVCs since their detection, no single explanation has encompassed the vast quantity of data that has been collected on the clouds (see \\markcite{WvW97} Wakker \\& van Woerden 1997, hereafter WvW97, and references therein). Particularly important is the lack of agreement on a characteristic distance for the clouds, because most of the relevant physical parameters depend on distance to one order or another.  In the 1970s, a well--defined subset of the clouds was identified as a tidal stream associated with the Magellanic Clouds (\\markcite{Mathewson74}Mathewson, Cleary, \\& Murray 1974), but since then no consensus has arisen regarding the nature of the remaining clouds which constitute the majority of HVCs.  In this paper, we suggest that the HVCs represent infall of the intergalactic medium onto the Local Group. Previous authors have explored the possibility that the HVCs are infalling primordial gas (\\markcite{Oort66,Oort70}Oort 1966, 1970) and have associated the HVCs with the Local Group (\\markcite{Verschuur69}Verschuur 1969; \\markcite{Kerr69}Kerr \\& Sullivan 1969; \\markcite{Arp85}Arp 1985; \\markcite{Bajaja87}Bajaja, Morras, \\& P\\\"oppel 1987; \\markcite{Arp91}Arp \\& Sulentic 1991), but subsequent observations always produced fundamental difficulties for the particular models considered.  Here, we assemble evidence based on new general--purpose surveys of atomic hydrogen gas by \\markcite{Stark92}Stark et al. (1992) and by \\markcite{Hartmann97}Hartmann \\& Burton (1997), and on the HVC surveys by \\markcite{Hulsbosch88}Hulsbosch \\& Wakker (1988) and by \\markcite{Bajaja85}Bajaja et al. (1985), and consider theoretical arguments in the context of modern cosmology. We argue that the clouds are matter accreting onto the Local Group of galaxies. Their velocities would thus largely reflect the motion of the clouds in the gravitational potential of the Local Group and the motion of the LSR about the Galactic center.  We suggest that the clouds represent the building blocks from which the Local Group was assembled and that they continue to fuel star formation in the disk of the Milky Way.  The evidence is presented as follows. In Section 2, we review some of the observed properties of the high--velocity clouds, and indicate those which are not consistent with a Galactic origin for the HVCs. In Section 3, we detail the observations entering our analysis. In Section 4, we discuss the stability of the HVCs against gravitational collapse and against Galactic tidal forces, and suggest that these considerations imply  that the most of the clouds are extragalactic at distances typical of the Local Group. The stability criteria imply, however, that the largest clouds are nearby and possibly are interacting with the Milky Way.  In Section 5, we discuss three individual clouds, one of which must be beyond the disk of the Milky Way, and two others that appear to be associated with M31 and M33, suggesting that at least some of the HVCs may be extragalactic. We identify a subset of the HVCs centered near the barycenter of the Local Group, and show that its kinematics as well as that of the entire HVC ensemble are well described as being at rest with respect to the Local--Group Standard of Rest (LGSR); the kinematics are thus inconsistent with a Galactic origin.  The entire HVC ensemble is also shown to exhibit an angular--size/velocity relation consistent with membership in the Local Group.  In Section 6, we simulate the accretion history of the Local Group and show that the simulation accounts for the observed distribution and kinematics of the HVC ensemble. The agreement between the simulation and the observations supports inferences about similarities between the Local Group HVCs and the Ly--$\\alpha$ absorbing clouds observed toward quasars. We show that the velocity extrema observed for the HVCs are consistent with their membership in the Local Group.  In Section 7, we discuss the distances and abundances of the HVCs in the context of the Local Group HVC hypothesis, and show that the hypothesis is consistent with all of the observations made to date. We review extragalactic HI searches for HVCs which have revealed clouds with properties similar to those we derive, in about the expected numbers.  In Section 8, we discuss the implied mass accretion rate, and implications for the chemical evolution of the disk of the the Milky Way.  We also present evidence for the Galactic fountain in low--velocity HI which suggests that the HI disk of the Milky Way is not in hydrostatic equilibrium.  In Section 9, we conclude by discussing predictions and future tests of the model, and summarize the principal arguments made in this paper. ",
        "conclusions": "\\label{sec:finale}   Most cosmologists believe that galaxy formation is a hierarchical process:  galaxies grow by accreting small clouds of gas and dark matter. This process is a continuing one and we expect that galaxies and groups are currently accreting new clouds. We simulated this process for the Local Group and found that properties of the accreted clouds are similar to certain properties of the high--velocity--cloud phenomenon (excluding the Magellanic Stream HVCs):  $\\bullet $ Most of the HVCs are located either near the general direction of M31, towards the barycenter of the Local Group, or in the antibarycenter direction, some 180\\dege from the direction of M31 (see Figure~\\ref{fig:lb}).  $\\bullet $ HVCs have chemical abundances similar to that of intra--group gas, and different from the abundances characteristic of the inner Galaxy. If HVCs were ejected from the inner Galaxy as part of a Galactic fountain, then their metal abundance would exceed the solar value, and this is not observed.  $\\bullet $ HVCs have an angular--size/velocity relation that is consistent with the clouds being nearly self--gravitating, and at a distance of $\\sim 300$ kpc.  If the Local--Group HVC hypothesis discussed in this paper is correct, then studies of HVCs can directly probe the process of galaxy formation. The validity of this hypothesis can be tested by a number of future observations:  $\\bullet $ Most observations of nearby galaxies would not have detected the gas clouds that are equivalent to the HVCs. Moreover, many HI maps of external galaxies extend just beyond the Holmberg radius. Our discussion would have the typical HVC located nearly a Mpc from the galactic center. It will be interesting to test our hypothesis with deep HI observations of isolated groups and filaments, searching for HI clouds associated with groups, rather than with individual galaxies.  $\\bullet $ Lyman--limit clouds, which are seen in absorption towards distant quasars, have column densities similar to those of the HVCs. Observations of nearby Lyman--limit and Lyman--$\\alpha $ systems show that they are not all associated directly with individual galaxies, but rather with groups of galaxies (\\markcite{Oort81}Oort 1981; \\markcite{Stocke95}Stocke et al. 1995; \\markcite{vanGorkom96}van Gorkom et al. 1996; \\markcite{Rauch96}Rauch, Weymann, \\& Morris 1996). In the scenario outlined in this paper, we expect that these clouds would have properties similar to those of the local HVCs. Thus, it would be interesting to use STIS to look for lower--column--density high--velocity HI clouds, which would correspond to the Lyman--$\\alpha $ clouds.  $\\bullet $ The simulations predict large amounts of gas accreting onto M31 and the Local Group from the region of space beyond M31, under the gravitational attraction of both M31 and our own Galaxy. Because this gas is several Mpc away, the gas clouds are expected to have small angular sizes and relatively low column densities. Deep HI observations in the M31 direction should be able to detect this gas.  The hypothesis central to this paper, namely that HVCs are at distances of around 1 Mpc, would be falsified by the detection of absorbtion in an HVC seen against stars in the Milky Way halo in the direction of M31 or in the anti--M31 direction.  On the other hand, further measurements of low levels of H$\\alpha$ emission towards these HVCs will strengthen the case for their extragalactic nature."
    },
    "9803/astro-ph9803067_arXiv.txt": {
        "abstract": "A solitary millisecond pulsar, if near the mass limit, and undergoing a phase transition, either first or second order, provided the transition is to a substantially more compressible phase, will emit a blatantly obvious signal---spontaneous spin-up.  Normally a pulsar spins down by angular momentum loss to radiation. The signal is trivial to detect and  is estimated to be ``on'' for 1/50 of the spin-down era of millisecond pulsars. Presently about 25 solitary millisecond pulsars are known. The phenomenon is analogous to ``backbending'' observed in high spin nuclei in the 1970's. ",
        "introduction": "The formation of a new phase of matter, a softer one than nuclear matter,  may cause a rapidly rotating pulsar to produce a prolonged signal that is dramatic,  easy to detect and easy to understand \\cite{glen97:a}.   The most plausible high density phase  transition is deconfinement as predicted by QCD \\cite{asymptotic}.  The signal I will describe will occur for  either a first or second order transition so long as  it is accompanied by a sufficient softening of the \\eosp.  (Cf. Fig.\\ \\ref{eos}.)  Strictly speaking we do not even know that quarks can be deconfined under  extreme conditions or otherwise. It is an `expectation' based on the QCD property of asymptotic freedom \\cite{asymptotic}. We would like to prove that this phase is a possible  phase of matter. If so, it would have pervaded the very early  universe, but  quark confinement in hadrons occurred at an early  time   and the thermal equilibrium that existed then leaves no signal today.   \\begin{figure}[htb] \\vspace{-.25in} \\begin{minipage}[t]{80mm} \\makebox[79mm] {\\psfig{figure=ps.qm1,width=3in,height=3.6in} } \\caption {The \\eos (labeled `Hybrid') of neutron star matter with a first order deconfinement phase transition. The normal phase contains nucleons, hyperons and leptons in equilibrium. The mixed phase contains as well, the three light flavor quarks. Comparison is made with a case in which deconfinement is not taken account of (labeled `n+p+H'). (Nuclear properties include the observed binding, saturation density, symmetry energy and  $K=300$ MeV, $m^\\star_{{\\rm sat}}/m=0.7$) \\label{eos_k300_y_h} \\label{eos} } \\end{minipage} \\hspace{\\fill} \\begin{minipage}[t]{75mm} \\makebox[74mm] {\\psfig{figure=ps.qm2,width=3in,height=3.6in}} \\caption { Density profiles of two stars of the same mass $M=1.42 \\msun$ but differing composition; (1) Hyperon star (neutron-proton-hyperon-lepton), (2) Hybrid (a pure quark matter core surrounded by mixed phase and outer pure hadronic confined phase).  \\Eoss as in Fig.\\ \\protect\\ref{eos_k300_y_h}. Interior differences are dramatic but not directly measureable. (For a description of neutron star matter and relativistic stars see Ref.\\ \\protect\\cite{book}. \\label{prof_k240} } \\end{minipage} \\end{figure}  From the balance of gravitational and centrifugal forces on a particle at the surface of rapidly rotating stars such as the millisecond pulsars, we know that the central density is a few times nuclear, the same range of energy densities as are expected to be produced in relativistic nuclear collisions. Let us assume that the critical deconfinement density occurs in the density range spanned by spherical stars and therefore in the population of slow pulsars to which the Crab belongs and of which there are about 800 presently known. In this case  newly born neutron stars have a quark core essentially from birth. But we have no way to tell if this is indeed the case. Models of neutron stars having different composition  generally differ in the range of masses and radii permitted, and their density profiles may be very different (cf. Fig.\\ \\ref{prof_k240}). However as far as measureable properties are concerned, ordinary neutron stars and hybrid stars (neutron stars with a quark core) are practically indistinguishable. Cooling rates are presumeably sensitive to the internal composition, but theoretical estimates are very uncertain.   I will sketch the generally accepted evolutionary  life of pulsars \\cite{heuvel91:a}. There are two distinct populations of pulsars (see Figs.\\ \\ref{period} and \\ref{pulsar}), the  canonical pulsars (about 800 of them now known) with periods between  1/50 sec and 8 sec, and the millisecond pulsars (about 50), which are believed  to be more evolved than the former. Stars are born into the first of these populations, and may, on a very long time-scale, evolve into the second. (See Fig.\\ \\ref{pulsar}.)  \\begin{figure}[htb] \\vspace{-.25in} \\begin{center} \\begin{minipage}[t]{130mm}\\hspace{.8in} \\makebox[79mm] {\\psfig{figure=ps.qm3,width=4.2in,height=3.5in} } \\caption { Distribution of pulsar periods. The lower group consists of `recycled' millisecond pulsars, the higher to the canonical pulsars in the first stage of their evolution (see Fig.\\ \\protect\\ref{pulsar}). \\label{period} } \\end{minipage} \\end{center} \\vspace{-.35in} \\end{figure} As the stellar core of a luminous star collapses to form a neutron star, it is spun up by conservation of angular momentum and acquires an enormous magnetic field of $10^{12}{\\rm~to~}10^{13}$ gauss because of flux conservation. The star is born as a rotating magnetic dipole. It has a tremendous store of rotational energy that will keep it spinning for 10 million years. The electromagnetic radiation beamed along the spinning dipole is what we see as pulsed radio emission once each rotation, if as observers, we lie on the cone swept out by the beam.  \\begin{figure}[htb] \\vspace{-.25in} \\begin{center} \\begin{minipage}[t]{130mm}\\hspace{.8in} \\makebox[79mm] {\\psfig{figure=ps.qm4,width=4.2in,height=3.5in} } \\vspace{-.35in} \\caption { Pulsars are born with field $B\\sim 10^{12} {\\rm~to~} 10^{13}$ gauss and evolve toward the right. Periods change with time as $\\sim 1/P$ and so pulsars accumulate at large $P$. They become radio silent in about $10^7$ years and remain stagnant until they capture a companion star, or unless they had one all along. Accretion from the less dense companion spins them up along a line like `silent evolution'.  A combination of the now weaker field but higher frequency turns them on again as `recycled' millisecond pulsars. They then again evolve toward shorter period, but now on a very long time-scale because of the weaker field. \\label{pulsar} } \\end{minipage} \\end{center} \\vspace{-.35in} \\end{figure}  In a plot of magnetic field $B$ vs period $P$ (Fig.\\ \\ref{pulsar}), stars move from top left to right because of loss of angular momentum to radiation. They disappear as active pulsars when a combination of angular velocity  and field strength is insufficient to produce radiation. It takes about $10^7$ years to complete this first phase. However, either from birth, or afterward, the star may have  or capture  a lower density companion as is often evidenced by the presence of an orbiting white dwarf. During an accretion era, the compact star is spun up by infalling matter that it tears off from its less dense companion. It looses some magnetic field during accretion, perhaps by ohmic resistance during the long radio silent  era.  In this `silent' era the the star moves diagonally from top right to bottom left in the $B-P$ diagram. The neutron star becomes centrifugally flattened as it approaches millisecond periods. The central density falls. The core of quark matter shrinks as quarks recombine to form baryons. The quark core may disappear altogether.  The silent neutron star, having completed a part of its life cycle, turns on again as a millisecond pulsar of low magnetic field when the lower field but higher angular velocity can once again produce radiation. Presently 50 such pulsars have been discovered, half of which still have a binary companion. The number of millisecond pulsars presently known is believed to be a fraction of the total population  because of search selection effects.  Millisecond  pulsars, which have    weaker fields ($10^8{\\rm~to~}10^9$ gauss), spin down very slowly since the deceleration is proportional to $B^2$. Their characteristic age is $P/2\\dot{P}\\sim 10^9$ years. The central density  is initially centrifugally diluted but as it spins down, the central density will rise again and the critical density will be reached, first at the center, and then in an expanding region. The growth of the central region of deconfined matter is paced by the slow spin-down, slow because of the coupling of rotation of the  stellar magnetic dipole to electromagnetic processes.  Stiff  nuclear matter is being replaced in the core  by highly compressible quark matter. The weight of the overlaying layers of nuclear matter weigh down on the core and compress it. Its density rises. The star shrinks---mass is redistributed with growing concentration at the center. The by-now more massive central region gravitationally compresses the outer nuclear matter even further, amplifying the effect. The density profile for a star at three angular velocities,  (1) the limiting Kepler velocity which is stretched in the equatorial plane and centrally diluted,  (2) an intermediate angular velocity, and  (3) a non-rotating star,  are shown in Fig.\\  \\ref{prof_k300b180}. We see that the central density rises with decreasing angular velocity by a factor of three and the equatorial radius decreases by 30 percent. In contrast, for a model for which the phase transition did not take place, the central density would change by only a few percent \\cite{weber90:d}. The phase boundaries are shown in Fig.\\ \\ref{omega_r_k300B180} from the highest rotational frequency to zero rotation. \\begin{figure}[htb] \\vspace{-.5in} \\begin{minipage}[t]{80mm} \\makebox[79mm] {\\psfig{figure=ps.qm5,width=3in,height=3.6in} } \\caption { Mass profiles as a function of equatorial radius of a star rotating at three different frequencies, as marked. At low frequency the star is very dense in its core, having a 4 km central region of highly compressible pure quark matter. At intermediate frequency, the pure quark matter phase is absent and the central 8 km is occupied by the mixed phase. At higher frequency (nearer $\\Omega_K$) the star is relatively dilute in the center and centrifugally stretched.  Inflections at $\\epsilon=220$ and $950$ are the boundaries of the mixed phase. \\label{prof_k300b180} } \\end{minipage} \\hspace{\\fill} \\begin{minipage}[t]{75mm} \\makebox[74mm] {\\psfig{figure=ps.qm6,width=3in,height=3.6in} } \\caption {  Radial boundaries at various rotational frequencies separating  (1) pure quark matter,  (2) mixed phase, (3) pure hadronic phase, (4) ionic crust of neutron rich nuclei and surface of star. The pure quark phase appears only when the  frequency is below $\\Omega \\sim 1370$ rad/s. Note the decreasing radius as the frequency falls. The frequencies of two pulsars, the Crab and PSR 1937+21 are marked for reference. \\label{omega_r_k300B180} } \\end{minipage} \\end{figure} The redistribution of mass and shrinkage of the star change its moment of inertia and hence the characteristics of its spin behavior. The star must spin up to conserve angular momentum which is being carried off only  slowly by the weak electromagnetic dipole radiation. The star  behaves like an ice skater who goes into a spin with arms outstretched, is slowly spun down by friction, temporarily spins up by pulling the arms inward, after which friction takes over again.  It is that simple to describe, and that is the blatant signal I mentioned---the spontaneous spin-up of an isolated millisecond pulsar that is radiating angular momentum and ought otherwise to be slowing down. ",
        "conclusions": "An isolated millisecond pulsar will spin up over an epoch of $2 \\times10^7$  years out of a  spin-down life of $10^9$  years if it undergoes a phase transition obeying the two conditions (i) the  transition causes a substantial softening of the \\eosp, and (ii) the critical density is attained in stars very near the mass limit. The spin-up epoch, compared to the spin-down life of the pulsar, corresponds to an `event rate' of 1/50. The determination of whether a pulsar is spinning up or down is trivial. Of the presently known millisecond pulsars, about 25 are isolated. We are approaching the moment of truth for this observable signal of a phase transition.  We have emphasized that the transition need not be of first order so long as it is accompanied by a sufficient softening of the \\eosp. We do not have a measure of what we mean by this. Of course our model does possess the requisite softening, or our results would not have exhibited backbending.  If no pulsar is observed to produce the signal, little is learned. Just as in the search for deconfinement in high energy  nuclear collisions, failure to observe a signal does not inform us that the deconfined phase does not exist."
    },
    "9803/astro-ph9803192_arXiv.txt": {
        "abstract": "We construct evolutionary synthesis models for simple stellar populations using the evolutionary tracks from the Padova group (1993, 1994), theoretical colour calibrations from Lejeune et al. (\\cite{lejeune}, \\cite{lejeune1}) and fit functions for stellar atmospheric indices from Worthey et al. (\\cite{worthey}).  A Monte-Carlo technique allows us to obtain a smooth time evolution of both broad band colours in UBVRIK and a series of stellar absorption features for Single Burst Stellar Populations ({\\bf SSPs}). We present colours and indices for SSPs with ages from $1 \\cdot 10^9$ yrs to $1.6 \\cdot 10^{10}$ yrs and metallicities $[M/H]$ = -2.3, -1.7, -0.7, -0.4, 0.0 and 0.4.  Model colours and indices at an age of about a Hubble time are in good agreement with observed colours and indices of the Galactic and M 31 GCs. ",
        "introduction": "Colour distributions of Globular Cluster ({\\bf GC}) systems are observed for a large number of early-type galaxies (E, S0, dE, cD) using ground-based Washington or HST broad band photometry. In most cases double-peak or broad/mul\\-ti-peak colour distributions are seen (e.g. Zepf \\& Ashman \\cite{zepf}, Elson \\& Santiago \\cite{elson}, Kissler-Patig \\etal \\cite{kissler}).  If the different colour subpopulations of GCs are formed in different events then they may contain clues to the formation history of their parent galaxies.  For example a two-peak colour distribution may result, if in addition to a primary initial collapse population of GCs, a secondary population of GCs were formed either in a merger-induced starburst (Schweizer \\cite{schweizer}, Ashman \\& Zepf \\cite{ashman}, Fritze -- v. Alvensleben \\& Gerhard \\cite{fritze1}, Fritze -- v. Alvensleben \\& Burkert \\cite{fritze2}) or else in some distinct secondary phase of cluster formation within the original galaxy (Forbes \\etal \\cite{forbes}). Likewise the broad or multi-peaked colour distribution often observed in GC systems around cD galaxies may point to a series of GC formation events during the hierarchical assembly of the parent galaxy or to some protracted GC formation or accretion mechanism.  A well-known difficulty with the interpretation of colour distributions is the degeneracy of colours with respect to age and metallicity.  While for Washington photometry there are well established and reliable calibrations of colours in terms of metallicity, the situation with HST broad band observations of GC systems is less clear.  A better understanding of the formation of composite GC systems would be possible if separate age and metallicity distributions could be disentangled from an observed colour distribution.  A second issue concerns the interpretation of colours for young star cluster systems detected with HST in many interacting and starburst galaxies.  The question is, if these YSC systems -- at least some fraction of them -- are the progenitors of GC systems. In an attempt to answer this question star clusters are being imaged with HST in an age sequence of interacting galaxies -- from early stages of interaction through merger remnants up to E/S0s (eg. Schweizer \\etal \\cite{schweizer1}, Whitmore \\etal \\cite{whitmore}, Miller \\etal \\cite{miller}). With 10 m class telescopes, spectroscopy of the brighter members of young star cluster populations is becoming possible (Kissler-Patig \\etal \\cite{kissler1}, Brodie \\etal \\cite{brodie}, but see also 4 -- 5 m class spectra by Schweizer \\& Seitzer \\cite{schweizer} or Zepf \\etal \\cite{zepf1}).  Spectroscopy will only be possible for a subsample of YSCs. Thus the determination of ages and metallicities from broad band colors will still be necessary.  It is thus desirable to study the evolution of broad band colours and absorption indices for single burst stellar populations of various metallicities using the most recent and complete stellar evolutionary tracks as well as careful colour and index calibrations.This allows one to obtain theoretical calibrations of broad band colours and indices in terms of metallicity over the full range of ages under investigation, i.e. from $10^7$ yr to a Hubble time.  Since theoretical calculations for the evolution of stars are only available for a discrete grid of masses, some means for obtaining a smooth evolution of the composite population is needed. Applying the tracks as they are would create discontinuities because all stars of a given mass would reach the giant branch at the same time, dominating the integrated light until they die. This effect is large for populations with stars that have about the same age. The effect also increases with age of the whole population, since the differences in both the lifetimes and luminosities between the main sequence and the later stages increase with decreasing mass.  For this work, we use the \\emph{Monte Carlo} method to bypass this problem while still avoiding the interpolation of tracks with its accompaining danger of creating artificial states. This is described in detail in section \\ref{sec_nummethod}.  The star formation history of any stellar system can be described by a superposition of SSP models of different ages and metallicities. An example of this is given by Cellone \\& Forte et al (\\cite{cellone}) in their study of Low Surface Brightness galaxies or Contardo et al. (\\cite{contardo}) who investigate the formation and evolution of galaxies in a cosmological scenario. ",
        "conclusions": "In this work we present Monte-Carlo evolutionary synthesis models for SSPs which cover a wide range in metallicity, from $Z=0.0001$ to $Z=0.05$, using most recent and complete sets of input physics: stellar evolutionary tracks for stellar masses from $0.08 M_\\odot$ to $120 M_\\odot$, including post - helium flash evolution and mass loss, model atmosphere libraries also covering late stellar types and giving colours from $U$ to $K$ in agreement with observations and empirical calibrations for a series of absorption indices.  We obtain theoretical calibrations of colours and indices in terms of metallicity which for model ages of 10-15 Gys agree closely with observations of GCs. The theoretical calibrations extend beyond the range of observed GCs, i.e. to a metallicity up to $[M/H] \\leq 0.4$. Moreover, our models provide these theoretical calibrations for all ages from cluster formation to 15 Gyrs and thus can also be applied in the interpretation of young star cluster systems observed in many interacting and starburst galaxies.  The complete model files are available via WWW on {\\tt http://www.uni-sw.gwdg.de/\\~\\ okurth/ssp.html}. There are also models with different parameters for the IMF and for mass loss available."
    },
    "9803/astro-ph9803100_arXiv.txt": {
        "abstract": "I present several simple figures to illustrate cosmology and structure formation in a nutshell. Then I discuss the following argument: if we assume that $\\Omega_{\\Lambda} = 0$ then the CMB results favor high $\\Omega_{m}$ while the supernova results favor low $\\Omega_{m}$. This large inconsistency is strong evidence for the incorrectness of the $\\Omega_{\\Lambda}=0$  assumption. Finally I discuss recent CMB results on the slope and normalization of the primordial power spectrum. ",
        "introduction": "The Big Bang model became the standard cosmological model soon after the discovery of the cosmic microwave background (CMB). The Big Bang model has a hot, dense early epoch (see Figure 1) when nucleosynthesis occurred. It also has an opaque surface that can naturally produce the Planckian spectrum of the CMB. The Steady State universe was not hotter in the past, has no epoch of Steady State Nucleosynthesis and has no opaque surface to produce the CMB.  Gravitational collapse is the leading model of structure formation (Figure 2). Slight over-densities are gravitationally unstable and collapse under their own self-gravity. In an alternative family of models, structure forms from topological defects. In gravitational collapse models CMB anisotropies larger than $\\sim 1$ degree are acausal and rely on inflation to explain their existence. In defect models these large anisotropies are close by, causal and sub-horizon sized. \\thefirstfig  \\clearpage \\thesecondfig  \\setcounter{figure}{2} \\begin{figure}[hbt] \\centerline{\\psfig{file=wearehere.eps,height=70mm,width=110mm,angle=-90}} \\vspace{0cm} \\caption{Galaxies are CMB anisotropies are Quantum fluctuations. According to the inflationary scenario, quantum fluctuations of a scalar field are the origin of all structures. These quantum fluctuations are not caused by any preceeding event in the same sense as radioactive decay or quantum tunneling are not caused. They are non-deterministic prime movers. Inflation of the universe by a factor of more than $10^{26}$ transforms these quantum fluctuations into super-horizon classical density fluctuations. On their way to becoming galaxies we can monitor their progress by looking at CMB maps.} \\end{figure}  \\clearpage One of the most important questions in cosmology is: What is the origin of all the galaxies, clusters, great walls, filaments and voids we see around us? The inflationary scenario provides the most popular explanation for the origin of these structures: they used to be quantum fluctuations.  Figure 3 illustrates the metamorphosis of quantum fluctuations to CMB anisotropies to galaxies. Primordial quantum fluctuations of a scalar field get amplified and evolve to become classical seed perturbations and eventually large scale structure. This process can be monitored by CMB observations since matter fluctuations produce temperature fluctuations in the CMB: $\\frac{\\delta \\rho}{\\rho} \\propto \\frac{\\Delta T}{T}$.  How does a particular fluctuation know whether it will become a spiral or an elliptical galaxy? Does the density and irregularity of its environment determine its morphology by controlling its angular momentum and the amount of merging? With a full understanding of galaxy formation we may be able to look at CMB cold spots and their neighborhoods and predict where they will end up in the Hubble tuning fork diagram of galaxy types. The distribution of morphological types at high redshift discussed by Driver in these proceedings would then be a derivable function of the characteristics of the CMB anisotropies.  \\thefourthfig  \\subsection{There is no scale beyond which the universe is homogeneous}  It has been claimed that some recent, deep, galaxy redshift surveys have reached the scale at which the Universe becomes homogeneous. Strictly speaking however there is no scale beyond which the universe is homogeneous. The amplitude of the density contrast ($\\delta \\rho / \\rho \\propto k^{3}P(k)$ decreases for larger scales but is never zero. A more meaningful question is: Where is the turnover in the power spectrum? This turnover is due to a suppression of growth of a given k mode by $k^{4}$ relative to modes which enter the horizon during matter domination (assuming $\\oo = 1$). Thus, the horizon scale at matter-radiation equality is an important diagnostic of this fundamental scale. See Figure 4.     Lineweaver \\& Barbosa (1998) have used current CMB anisotropy measurements to determine the position of the adiabatic peak in the CMB spectrum under the assumption of open or critical density CDM dominated universes: $\\ell_{peak} = 260^{+30}_{-20}$.  Figure 5 illustrates how harmonic sound bumps appear in the CMB power spectrum driven by the wells and valleys of the CDM potentials. The epoch when matter and radiation densities are equal has a redshift of $z_{eq}$ while decoupling occurs at $z_{dec}$. The number of oscillations between $z_{eq}$ and $z_{dec}$ and thus the phase of the oscillations at $z_{dec}$ is determined by i) the physical size of the potential well, ii) the speed of sound and iii) the time interval between $z_{eq}$ and $z_{dec}$.  \\vspace{0.7cm}  \\thefifthfig \\clearpage ",
        "conclusions": ""
    },
    "9803/astro-ph9803336_arXiv.txt": {
        "abstract": "This paper presents the preliminary results of the ESO Imaging Survey (EIS), a public survey being carried out by ESO and member states, in preparation for the VLT first-light. The survey goals, organization, strategy and observations are discussed and an overview is given of the survey pipeline developed to handle EIS data and produce object catalogs. A report is presented on moderately deep I-band observations obtained in the first of four patches surveyed, covering a region of 3.2 square degrees centered at $\\alpha \\sim 22^h 40^m$ and $\\delta =-40^\\circ$. The products available to the community, including pixel maps (with astrometric and photometric calibrations) and the corresponding object catalogs, are also described. In order to evaluate the quality of the data, preliminary estimates are presented for the star and galaxy number counts, and for the angular two-point correlation function obtained from the available data.  The present work is meant as a preview of the final release of the EIS data that will become available later this year. ",
        "introduction": "With the advent of very large telescopes, such as the VLT, a largely unexplored domain of the universe becomes accessible to observations which may dramatically enhance on our understanding of different physical phenomena, in particular the origin and evolution of galaxies and large scale structures.  In the next few years a wide array of 8-m telescopes will become available world-wide.  Among these, the European VLT project is particularly striking because of its four 8-m telescopes and an impressive array of complementary instrumentation. Viewed as a unit, the VLT provides great flexibility by combining complementarity for certain programs with multiplexing capabilities for others. First-light for the VLT is scheduled for May 1998, with regular science operation starting in April 1999.  In order to take full-advantage of the VLT from the start of its operation, ESO and its Observing Programmes Committee (OPC) decided to coordinate an imaging survey to provide candidate targets well-suited to the first set of VLT instruments. The ESO Imaging Survey (EIS) has been conceived as a collaborative effort between ESO and astronomers in its member states. Following the recommendation of the OPC, the survey has been overseen by a Working Group (WG).  The EIS WG is composed of leading experts in different fields and has the responsibility of defining the survey science goals and strategy, and monitor its progress. In order to carry out the survey a dedicated team was assembled, starting March 1997. To stimulate cooperation between ESO and the astronomical community of the member states, EIS has sponsored the participation of experts as well as students and post-docs from the community in the development of software, observations and data reduction.   As described by Renzini \\& da Costa (1997) (see also ``http://www.eso.org/eis''), EIS consists of two parts: EIS-wide to search for rare objects (\\eg distant clusters and quasars) and EIS-deep to define samples of high-redshift galaxies.  These science goals were chosen to match as well as possible the capabilities of the first VLT instruments, FORS, ISAAC and UVES.  EIS is also an essential first step in the long-term effort, currently underway at ESO, to provide adequate imaging capabilities in support of VLT science (Renzini 1998). The investment made in EIS will be carried over to a Pilot Survey utilizing the ESO/MPIA 2.2m telescope at La Silla, with its new wide-field camera. This Pilot Survey, which will follow the model of EIS, has been recommended by the EIS WG and is being submitted to the OPC.  The goal of this paper is to describe the characteristics of I-band observations carried out in the fall of 1997 over a region of 3.2 square degrees (EIS patch A, da Costa \\etal 1998a) and of the corresponding data products, in the form of calibrated images and single frame catalogs.  These products have been made publicly available through the ESO Science Archive, as a first step towards the full distribution of the EIS data.  The purpose of the present release is also to provide potential users with a preview of the data, which may help them in the preparation of VLT proposals, and to encourage the community to provide constructive comments for the final release. It is important to emphasize that due to time limitations the results presented here should be viewed as preliminary and improvements are expected to be made before the final release of the EIS data later this year.  In section 2, a brief description is presented of the criteria adopted in the field selection, the strategy of observations and the characteristics of the data in patch A already completed. It also describes the filters used, the definition of the EIS magnitude system and its relation to other systems, and the data used for the photometric calibration of the survey. In section 3, a brief description of the data reduction pipeline is presented, followed in section 4 by a description of the data products made publicly available in this preliminary data release. In section 5, the algorithms used to detect and classify objects, and the information available in the catalogs being distributed are described. Preliminary results from a scientific evaluation of the data is presented in section 6. In section 7, future plans are presented, followed in section 8 by a brief summary. ",
        "conclusions": "The ESO Imaging Survey is being carried out to help the selection of targets for the first year of operation of VLT. This paper describes the motivation, field and filter selection, and data reduction pipeline.  Data for the first completed patch, in the form of astrometric and photometric calibrated pixels maps, single-frame catalogs, on-line coadded section images and further information on the project are available on the World Wide Web at ``http://\\-www.eso.org/eis''.  Preliminary evaluation of the data shows that the overall quality of is good and the completeness limit of the extracted catalogs is sufficiently deep to meet the science requirements of EIS. Furthermore, the results for the other patches should improve as the observing conditions were considerably better than those in the period patch A was observed.   The final and complete release of the data products of EIS is scheduled as follows: 1) EIS-wide, except the U-band : July 31, 1998, before the first call for proposals for the VLT; 2) EIS-deep and EIS-wide U-band on December 31, 1998."
    },
    "9803/astro-ph9803046_arXiv.txt": {
        "abstract": " ",
        "introduction": "Thermal radiation from surfaces of several radio pulsars has been  detected \\linebreak recently in soft X-rays by {\\it ROSAT\\/} and {\\it ASCA\\/} (see Becker and Tr\\\"umper,~1998). A number of the point-like X-ray sources has also been discovered and identified as radio-silent isolated cooling neutron stars (NSs) (see Caraveo et al.,~1996; Walter et al.,~1996; Vasisht et al.,~1997). In several cases, the observations seem to be better explained if the NSs possess an envelope of matter of low atomic weight, presumably hydrogen (Page,~1997), and if NS atmosphere models are applied to fit the data (Pavlov and Zavlin,~1998). There are indications that this may be the case for NSs with  different magnetic fields $B$: from weak, $B\\ll10^{10}$~G (Rajagopal and Romani,~1996; Zavlin and Pavlov,~1998; Zavlin et al.,~1998), to strong, $B\\sim10^{11-13}$~G (Page et al.,~1996; Zavlin et al.,~1998), and superstrong, $B>10^{14}$~G (Heyl and Hernquist,~1997). To justfy this, one needs in further observations and advanced atmosphere models of NSs with high $B$. Thermal motion of atoms accross the magnetic field induces an electric field in the frame of the atom. This affects atomic structure, atmospheric thermodynamics and opacities (Ventura et al.,~1992; Pavlov and M\\'esz\\'aros,~1993), being a critical point in the construction of models and data interpretation. Here we study this problem for hydrogen plasma at temperatures $T\\sim 10^{5.5}-10^{6}$~K, densities $\\rho\\sim10^{-3}-10^1{\\rm~g~cm}^{-3}$, and magnetic fields $B\\sim 10^{12}-10^{13}$~G, typical for atmospheres  of middle-aged cooling NSs. We construct an analytic model of the plasma free energy and derive a generalized Saha equation which is used to obtain the opacities. ",
        "conclusions": ""
    },
    "9803/astro-ph9803270_arXiv.txt": {
        "abstract": "Warped \\ion{H}{1} gas layers in the outer regions of spiral galaxies usually display a noticeably twisted structure.  This structure almost certainly arises primarily as a result of differential precession in the \\ion{H}{1} disk as it settles toward a preferred orientation in an underlying dark halo potential well that is not spherically symmetric.  In an attempt to better understand the structure and evolution of these twisted, warped disk structures, we have adopted the ``twist--equation'' formalism originally developed by Petterson (1977) to study accretion onto compact objects.  Utilizing more recent treatments of this formalism, we have generalized the twist--equation to allow for the treatment of non--Keplerian disks and from it have derived a steady--state structure of twisted disks that develops from free precession in a nonspherical, logarithmic halo potential. We have used this steady--state solution to produce \\ion{H}{1} maps of five galaxies (M83, NGC 300, NGC 2841, NGC 5033, NGC5055), which match the general features of the observed maps of these galaxies quite well.  In addition, the model provides an avenue through which the kinematical viscosity of the \\ion{H}{1} disk and the quadrupole distortion of the dark halo in each galaxy can be quantified.  This generalized equation can also be used to examine the time-evolutionary behavior of warped galaxy disks. ",
        "introduction": "In simplest terms, spiral galaxy disks can be described as geometrically thin, flat, and circular.  We understand that spiral disks are geometrically thin because the gas of which they are composed is cold (the sound speed of the gas is much smaller than its circular orbital velocity); and they are both circular and flat because, being dissipative, the gas is fairly efficient at both minimizing out-of-the-plane motions and radial excursions that would lead to departures from circular orbits.  Describing spiral disks as {\\it perfectly} circular and flat is clearly an oversimplification, however.  In addition to the nonaxisymmetric structures that are obvious in optical photographs of many spiral disks, 21-cm maps of the projected velocity fields of spiral disks often reveal isovelocity contours that are significantly twisted (\\cite{RLW}; \\cite{RN78}; \\cite{RCC79}; \\cite{N80a}, 1980b; \\cite{B81}; \\cite{S85}).   Kinematical tilted-ring models have been constructed in an effort to explain the presence of such twists in the velocity maps of \\ion{H}{1} disks.  The models indicate that the outer regions of many normal spiral disks are significantly warped out of the principal plane that is defined by the optically visible, central portion of each galaxy.  The line of nodes that defines the intersection of adjacent rings of gas in these kinematical models also usually must be twisted significantly as a function of radius in order to explain the observed contour maps (for recent reviews, see \\cite{Briggs90}; \\cite{Bosma91}; \\cite{CTSC93}, hereafter CTSC).  It is not particularly surprising that many galaxies are observed to possess extended, rotationally flattened disks because such structures appear to be fairly ubiquitous in gravitationally bound, astrophysical systems.  (There is strong evidence, for example, that rotationally flattened disks either exist now or have existed in the past around our sun, individual planets within our solar system, numerous protostars, the primary star of many mass--exchanging binary star systems, and active galactic nuclei.) What is peculiar about the \\ion{H}{1} disks of many galaxies is that the disks are significantly warped. It is not clear why natural dissipative processes similar to those which work effectively to minimize out-of-the-plane motions in stellar or protostellar ``accretion'' disks are unable to suppress warps in the gaseous disks of galaxies. As Binney (1992) has reviewed, there still is no generally accepted dynamical model of spiral galaxies that satisfactorily explains either the origin or the current structure of warped \\ion{H}{1} disks.  Almost thirty years ago, Hunter \\& Toomre (1969) examined whether normal, infinitesimal bending oscillations might be exhibited by thin, rotating disks of self-gravitating material, permitting them to sustain coherent warps for more than a Hubble time. They concluded that ``for any disk whose density tapers sufficiently gradually to zero near its edge,'' the frequency spectrum of such oscillations is at least partly continuous and, as a result, coherent warps cannot be sustained. Over the subsequent decade, a number of other ideas surfaced to explain the persistence of warps in galaxies, each one taking advantage of the demonstrated existence of dark matter halos around galaxies.  As Toomre (1983) has reviewed, however, models proposing to use the halo as an active agent to excite warps in otherwise flat disks -- for example, via the Mathieu instability (Binney 1981) or via a flapping instability (Bertin \\& Mark 1980) -- each present significant difficulties. Toomre (1983) and Dekel \\& Shlosman (1983) proposed, instead, that untwisted steady-state warps may be sustained as a result of steady forcing by a nonspherical, tilted halo. Building on the early work of Hunter \\& Toomre (1969) and the idea that forcing by a nonspherical dark halo can influence the dynamics of the visible disks of galaxies, Sparke (1986) and Sparke \\& Casertano (1988) have shown that a discrete warping mode can persist if the disk is sufficiently self-gravitating and if it is embedded in a nonspherical halo whose equatorial plane is tilted with respect to the centralmost regions of the disk.  Adopting this model of galaxy warps, the following evolutionary picture emerges.  During the galaxy formation process, gas which falls into a spheroidal dark matter halo generally will find that its angular momentum vector is tipped at some nonzero angle, $i$, away from the symmetry axis of the halo.  If the gas is cold, it will settle into a rotationally flattened disk that is tilted at the same angle $i$ with respect to the equatorial plane of the halo.  In the centralmost regions of the galaxy, where the self-gravity of the gas (or, ultimately, the combined gas/star system) dominates over the gravitational influence of the halo, the gas will be content to remain in orbits that preserve this original tilt.  In the outermost regions of the galaxy where the gravitational field of the halo dominates, however, the gas should settle into the halo's equatorial plane.  As Toomre (1983) and Dekel \\& Shlosman (1983) both sketched in their original concept papers, there also should be an intermediate region where the gas will be influenced significantly by the nonspherical gravitational fields of both the halo and the central gas (or gas/star) disk.  Through their modeling efforts, Sparke (1986) and Sparke \\& Casertano (1988) confirmed this earlier suspicion that in the intermediate region, the gas can reside in a steady-state, ``warped disk'' structure that provides a smooth radial transition between the separate ``flat disk'' orientations of the inner and outer regions of the galaxy.  Concentrating on the dynamics of the centralmost and intermediate regions -- that is, by building models in which there was effectively no gas in the outermost regions -- Sparke \\& Casertano (1988) showed that, in steady-state, the warped disk exhibits a straight line of nodes which precesses slowly and coherently in a direction retrograde to the orbital motion of the gas. In a time-dependent simulation, furthermore, Hofner \\& Sparke (1994) showed that settling to this steady-state warped disk structure occurs from the inside, out, and is driven not by dissipative processes that are similar to those which are thought to drive settling in most stellar or protostellar accretion disks but, rather, because ``bending waves carry energy associated with transient disturbances out toward the disk edge.'' They also showed that, during an evolution as the bending waves propagate outward through the intermediate region of the disk, a twisted structure can develop and persist until the gas has had sufficient time to settle into the steady-state (constant line of nodes) configuration. Sparke \\& Casertano (1988) and Hofner \\& Sparke (1994) have demonstrated that, with an appropriate choice of parameters, this model of disk warping matches well the observed properties of several galaxies with warped disks. (See also Kuijken 1991 and Dubinski \\& Kuijken 1995.)   As Hofner \\& Sparke (1994) have pointed out, in galaxies with extended \\ion{H}{1} disks ``the outermost gas cannot be expected to form part of a coherent warping mode.''  They did not include normal dissipative forces in their simulations and therefore were unable to comment on how such forces might influence the settling process.  In this paper, we examine the disk-settling process from the other extreme, ignoring the self-gravity of the gas but introducing an effective kinematical viscosity into the dynamical equations in order to simulate the effects of dissipative forces.  Hence, our effort is complementary to the work of Hofner \\& Sparke (1994) and is most relevant to galaxies with extended \\ion{H}{1} disks -- although there is one galaxy used for model comparisons (NGC 2841) that is shared by both works. We adopt the view that warps in extended \\ion{H}{1} disks which exhibit substantial twists are transient features. Independent of precisely what physical process was responsible for initially placing the gas in an orbit that is inclined to the halo's equatorial plane ({\\it e.g.}, gas infall at the time of formation or a recent tidal encounter with another galaxy), the twisted structure can be understood as the result of differential precession in the gaseous disk as it dissipatively settles toward that ``preferred plane.''  In the past, there has been considerable concern (first enunciated by Kahn \\& Woltjer [1959], but reiterated in the reviews by both Toomre [1983] and Binney [1991]) that differential precession will destroy any warped disk structure on a time scale that is short compared to a Hubble time and, therefore, that the mechanism we are examining cannot reasonably be used to explain the persistence of such structures. In the outermost regions of \\ion{H}{1} disks, however, precession times are relatively long and, as was first pointed out by Tubbs \\& Sanders (1979), a warped gas layer can persist for a Hubble time if the dark halo in which the disk is embedded deviates only slightly from spherical symmetry. By modeling carefully the process of disk settling that is driven by normal dissipative forces and comparing the models to the observed kinematical properties of galaxies with extended, warped \\ion{H}{1} layers, we hope to be able to more carefully examine the viability of such models.  Steiman-Cameron \\& Durisen (1988, hereafter SCD88) have developed this idea rather extensively. They have adopted a numerical, cloud-fluid model to simulate the time-dependent evolution of a galaxy disk that intially is tilted out of the equatorial plane of an underlying, spheroidal dark halo.  The disk is assumed to be composed of a set of annular mass elements, or ``clouds,'' which act like atoms in a viscous fluid.  The SCD88 model has offered some valuable physical insight into the time-dependent settling process that is driven by normal dissipative forces and their dynamically generated model of a twisted galaxy disk has been surprisingly successful at matching the peculiar optical image of one particular galaxy, NGC 4753 (\\cite{SCKD92}).  Our model is analogous to the one developed by SCD88 but it derives from an analytical prescription of the viscous settling process.  More specifically, we employ the ``twisted-disk'' equation formalism first developed by Bardeen \\& Petterson (1975) and Petterson (1977, 1978) to describe the time-dependent settling of a thin, viscous disk in a nonspherical dark halo potential.  This is a rather natural formalism to adopt because, as numerous kinematical ``tilted-ring'' models have demonstrated, warped \\ion{H}{1} galaxy disks display a structure that resembles, at least qualitatively, the twisted geometry that had once been thought to be important in accretion disks which surround certain compact stellar objects (Bardeen \\& Petterson 1975; see a recent rejuvenation of this idea put forward by Maloney \\& Begelman 1997).  In adapting the model to galaxy disks, we have replaced the approximate Keplerian gravitational potential used in earlier accretion disk work with a logarithmic potential appropriate to galaxy halos. (Pringle [1992] also recently described how the twisted--disk formalism may be adapted to galaxies.) In the limit of stress-free precession, our model reproduces the analytical prescription of disk settling first presented by SCD88, but our model is not constrained to this limit.  A more general solution to the governing equations predicts an exponential settling rate that depends on time to the first power, rather than on time to the third power as has been derived in the limit of stress-free precession.  Furthermore, an analytical, steady-state solution to the governing equations produces a twisted-disk structure that is very similar to previously constructed, kinematical models. We demonstrate that projected surface density maps and radial velocity maps derived from our analytical model match published \\ion{H}{1} maps of five well-studied warped disk galaxies (M83, NGC 300, NGC 2841, NGC 5033, and NGC 5055) very well. ",
        "conclusions": "Utilizing the formalism originally introduced by Petterson to describe warped and twisted accretion disk structures in Keplerian potentials, we have derived a single, complex ODE to describe time dependent settling of an \\ion{H}{1} disk in the logarithmic potential that appears to be typical of normal spiral galaxies.  Over the interval $1 \\lesssim x \\lesssim 20$, the analytical function $w_{ss}\\bigl(x\\bigr)$ -- derived from an analysis of the steady-state limit of the general twisted-disk equation -- appears to describe quite accurately the general warped and twisted structure that is exhibited by a number of galaxy disks.  It should be noted again that our analysis is based on the assumption that the warping angle $\\beta$ is $\\ll 1$. For galaxies with larger warps (including some that we have modeled) this is a simplifying assumption and should be disregarded in more proper treatments (see Pringle [1992]).  From our model fits we conclude that, quite generally, the effective kinematical viscosity in these neutral hydrogen disks is $\\nu\\sim0.6$ km s$^{-1}$ kpc. That is, the effective Reynolds number in these systems is \\begin{displaymath} R_e\\sim\\frac{r_{max} v_{\\psi}}{\\nu} \\sim \\frac{x_{max}^2 v_\\psi}{r_{max}} \\sim6000\\;. \\end{displaymath} According to equation (\\ref{radvsub}), this also implies that the ratio of the radial inflow velocity of the gas to its orbital velocity is \\begin{displaymath} \\frac{v_r}{v_{\\psi}}\\approx\\frac{1}{R_e}\\sim\\times{10}^{-4}\\;. \\end{displaymath}  This model also provides a mechanism by which the parameter $\\eta$ -- the quadrupole moment of the underlying dark halo potential well -- can be measured in spiral galaxies.  Our fits to five normal spirals with well-studied warped \\ion{H}{1} disks specifically indicate that $\\eta\\sim{10}^{-3}$.  (This value can be increased to $\\eta\\sim{10}^{-2}$, with a corresponding factor of ten increase in viscosity, only if the age of these disks is assumed to be 0.1$\\,{H_0}^{-1}$ -- which seems unlikely -- or $\\sigma$ in expression 38c is found to be $\\sim 0.1$.)  Hence, we conclude that the dark halos in which these warped disks sit are, to quite high accuracy, spherically symmetric.  This particular conclusion should not come as a surprise because some time ago Tubbs and Sanders (1979) pointed out that if warped disks are identified as transient structures, the warp can only be sustained for a Hubble time if the underlying halo potential is very nearly spherical. Although we have examined in detail here only the steady-state solution, the derived time-dependent twisted-disk equation provides a tool that can be utilized to model the time-evolution of warped \\ion{H}{1} disks without resorting to elaborate numerical techniques.  The twisted-disk formalism in general -- and the analytical function $w_{ss}\\bigl(x\\bigr)$ in particular -- provides an avenue through which models of warped \\ion{H}{1} disks can advance from purely kinematical fits (e.g. tilted-ring models) to dynamical models based on a reasonable physical model.  In the future, we also expect to use the time-dependent form of the twisted-disk equations as a point of comparison for fully three-dimensional gas dynamic simulations of \\ion{H}{1} disks settling into distorted halo potentials."
    },
    "9803/astro-ph9803028_arXiv.txt": {
        "abstract": "We present a total of 48~minutes of observations of the nearby, bright millisecond pulsar \\psr~taken at the Parkes radio observatory in Australia.  The data were obtained at a central radio frequency of 1380~MHz using a high-speed tape recorder that permitted coherent Nyquist sampling of 50 MHz of bandwidth in each of two polarizations. Using the high time resolution available from this voltage recording technique, we have studied a variety of single-pulse properties, most for the first time in a millisecond pulsar.  We show that individual pulses are broadband, have pulse widths ranging from $\\sim$10~$\\mu$s ($\\sim 0.6^{\\circ}$ in pulse longitude) to $\\sim$300~$\\mu$s ($\\sim 20^{\\circ}$) with a mean pulse width of $\\sim$~65$\\mu$s ($\\sim 4^{\\circ}$), exhibit a wide variety of morphologies, and can be highly linearly polarized.  Single pulse peaks can be as high as 205~Jy (over $\\sim$40 times the average pulse peak), and have a probability distribution similar to those of slow-rotating pulsars.  We observed no single pulse energy exceeding $\\sim$4.4 times the average pulse energy, ruling out ``giant pulses'' as have been seen for the Crab and \\ntts\\ pulsars.  \\psr\\ does not exhibit classical microstructure or show any signs of a preferred time scale that could be associated with primary emitters; single pulse modulation has been observed to be consistent with amplitude-modulated noise down to time scales of 80 ns. We observe a significant inverse correlation between pulse peak and width. Thus, the average pulse profile produced by selecting for large pulse peaks is narrower than the standard average profile.  We find no evidence for ``diffractive'' quantization effects in the individual pulse arrival times or amplitudes as have been reported for this pulsar at lower radio frequency using coarser time resolution (\\cite{amdv97}).  Overall, we find that the single pulse properties of \\psr\\ are similar to those of the common slow-rotating pulsars, even though this pulsar's magnetosphere and surface magnetic field are several orders of magnitude smaller than those of the general population.  The pulsar radio emission mechanism must therefore be insensitive to these fundamental neutron star properties. ",
        "introduction": "\\label{sec:intro}  Single pulse studies of millisecond pulsars are considerably more difficult to perform than are those of slow pulsars.  Faster data rates are needed to study millisecond pulsars with comparable pulse phase resolution, and finer radio frequency resolution is required to minimize the effect of interstellar dispersion.  Also, interstellar scattering time scales comparable to the pulse duration render studying individual pulse morphologies impossible, while steep millisecond pulsar spectra preclude observations of sufficient sensitivity at higher radio frequencies where scattering is less important.  Yet single pulse studies of millisecond pulsars are highly desirable for two reasons.  First, the origin of the radio emission that has made isolated neutron stars famous is, even 30 years after their discovery, still a mystery. The high brightness temperatures ($\\sim 10^{30}$K) associated with the radio emission point to coherent processes which are poorly understood even under less exotic conditions (\\cite{mel96}). Previous observations of slow pulsars have not constrained the emission mechanism sufficiently; the study of radio emission properties of millisecond pulsars may provide important new clues. Millisecond pulsars, because of their fast spin periods, have much smaller light-cylinder radii, and hence magnetospheres, than slow pulsars.  They also have lower surface magnetic field strengths than the general pulsar population (most likely resulting from their having been ``recycled'' by a binary companion through the accretion of mass and angular momentum). Were the radio emission mechanism at all dependent on such properties, millisecond pulsars should have different radio properties than the slower-spinning general population.  The second reason single pulse studies of millisecond pulsars are important is that millisecond pulsar timing is well-known to be an unparalleled source of precision astrometric and astrophysical information.  Among factors possibly limiting timing precision is the stability of the average profile, which depends on the properties of single pulses.  Only recently has a systematic study of single pulses from millisecond pulsars become possible, largely due to improving computational and data recording technologies.  To date, the only such investigation has been for the 1.5~ms pulsar \\ntts\\ (\\cite{wcs84,sb95,bac95,cstt96}). Interestingly, \\ntts\\ exhibits giant radio pulses like those seen elsewhere only in the Crab pulsar.  With the single pulse properties of only one millisecond pulsar having been studied in any detail, the question of whether  all millisecond pulsars show similar properties naturally arises.  Unfortunately, \\ntts\\ suffers interstellar scattering at time scales comparable to the duration of a single pulse at the radio frequencies at which it has been observed, rendering detailed study of individual pulse morphologies difficult.  Here we report on high-time-resolution single pulse studies of a second millisecond pulsar, \\psr.  This pulsar's large flux density and low dispersion measure (DM), and the corresponding scarcity of line-of-sight scattering material, render it an obvious target for single pulse work.  Some single pulse investigations of \\psr\\ have been reported (\\cite{jlh+93,amdv97}) but none have had sufficient time resolution to resolve most individual pulses. Using a fast recording device and powerful supercomputers, we have been able to resolve all pulses in our data, the narrowest being $\\sim$10~$\\mu$s. ",
        "conclusions": "\\label{sec:concl}  We have presented the first detailed single pulse study of a millisecond pulsar in which sufficient time resolution was available to resolve single pulses as they would be seen in the pulsar vicinity. The similarity of the single pulse properties to those of normal slow pulsars is remarkable given the dramatically reduced magnetospheric volume and magnetic field strength of \\psr.  Indeed, without being told the absolute sample rate, it is unlikely that one could distinguish between this being a millisecond or slow pulsar.  To summarize, our observations of \\psr\\ have: resolved individual pulses, and shown that they have a wide variety of morphologies, with multiple sub-pulse components not uncommon; shown that individual pulses are in general broadband; shown that individual pulses can have high linear polarization; provided no evidence for giant pulses as observed in the Crab pulsar and PSR B1937+21, nor for pulse nulling or drifting subpulse phenomena; found structure in the intensity fluctuation spectrum; revealed a correlation of pulse peak with pulse width so that the average profile formed from only the highest amplitude pulses is much narrower than the conventional average profile; not shown any evidence for micro-structure or preferred time scales $\\geq$80~ns; shown that the emission is consistent with an amplitude modulated noise model; provided no evidence to support the claim made by Ables \\etal\\ (1997) of the detection of coherent radiation patterns.  Because there exists no self-consistent radio emission mechanism context (see e.g. \\cite{mel96}) in which to discuss these results, it is difficult to say how such models are constrained.  Indeed previous slow pulsar single pulse studies suffered from the same difficulty. However, it is often the case that fundamental insights become apparent when known phenomenon are taken to extremes; this, and improved tape recording and computer technologies that permit single pulse studies of millisecond pulsars, provided our motivation in undertaking this analysis.  Similar studies of other millisecond pulsars may eventually lead to an understanding of the radio emission mechanism."
    },
    "9803/astro-ph9803264_arXiv.txt": {
        "abstract": "We present multi-frequency radio continuum VLBI observations of the gravitational lens system B0218+357 carried out using a global VLBI network and the VLBA. The source has been observed with resolutions from 0.2~mas to 5~mas and displays interesting structure. The spectral properties of various components show that the lensed object is a standard flat spectrum radio source which has many self-absorbed components. Based on the flux ratio of the lensed images as a function of frequency we propose a simple model for the background radio source. ",
        "introduction": "The radio source B0218+357 has been identified as a gravitationally lensed system (Patnaik et al. 1993). The source, which has a core-halo structure at low frequency (1.4~GHz) and low resolution (5~arcsec), consists of two compact flat-spectrum components, A and B, separated by 335 milliarcsec. The weaker of the two, B, is surrounded by a faint Einstein ring of similar diameter. The spectral and polarisation characteristics of the components show them to be the lensed images of a single flat spectrum `core'. The lensing galaxy, suggested to be a spiral galaxy, has been observed by the HST (Jackson et al. 1997). It has a redshift of 0.6847 (Browne et al. 1993). Atomic hydrogen and many molecular species have been detected in absorption against the background source (Carilli et al. 1993, Wiklind \\& Combes 1995, Menten \\& Reid 1996). The redshift of the background object is suggested to be 0.96 (Lawrence 1996).  In this contribution we explore the properties of the radio source from our multi-frequency VLBI observations. We give a summary of our observations and discuss the results in the context of a flat spectrum radio source. We note that the overall radio spectrum of B0218+357 is flat between 365~MHz and 43~GHz i.e. its flux density is constant within a factor of two between these frequencies. ",
        "conclusions": "We derive two basic results from these observations, namely the spectral decomposition of various components and the ratio of their flux densities as a function of frequency.  Since we resolve the `core-jet' structure in A and B at frequencies higher than 8.4~GHz, we have 4 measurements of flux density for these components. The spectra are plotted in Fig. 2. The spectrum of the core (A1 and B1) is self-absorbed with a peak around 15~GHz. The jet (A2 and B2) has steep spectra above 8.4~GHz. However, the total flux density of A and B at 1.7 and 5~GHz suggest that both A2 and B2 must also be self-absorbed between these two frequencies. The Einstein ring emission has a steeper spectrum between 5 and 22~GHz and it must also be self-absorbed at lower frequencies since the total flux density provides a limit (Patnaik et al. 1993).  In summary, the radio structure of B0218+357 consists of at least 3 distinct components, core, jet and the Einstein ring, which are all self-absorbed at different frequencies. The components are self-absorbed at progressively lower frequencies as one moves from the core along the jet. This behaviour of the radio source is consistent with its flat spectrum (Cotton et al. 1980).  The core, being self-absorbed around 15~GHz, contributes very little to the emission at 1.7 and 5~GHz, where the emission is dominated by the jet as evident from the spectra. Perhaps it is not surprising therefore that the radio structures of both A and B at 1.7~GHz appear diffuse. However, this poses a difficulty in that a self-absorbed component, which must be compact, has a large observed size. In this case one must consider the intrinsic size of the source rather than the observed size which has been magnified by the lens.  Another result from our observations is that the image flux ratios change with frequency. This is an apparent contradiction since gravitational lensing is achromatic. However, since the source is extended and spans areas of different image magnifications the observed flux ratio can indeed vary with frequency. The flux ratio of the core (i.e. A1/B1) is around 3.7 for frequencies higher than 8.4~GHz. This is similar to the ratio of A/B obtained from VLA observations at these frequencies (Patnaik et al. 1993). At lower frequencies the flux ratio A/B is around 2.6.  Of course, we point out that the source is variable at high frequencies and thus the measured flux ratio can be different due to differential time delay.  The above results lead to a rather simple model for B0218+357. The radio source consists of a number of different components which are self-absorbed at different frequencies thus conspiring to produce the observed flat spectrum. The core (imaged into A1 and B1) is located at the western-most edge, the jet (imaged into A2 and B2) lies between the core and the component giving rise to the Einstein ring. The core dominates the radio structure at frequencies higher than 8.4~GHz, and the jet and the Einstein ring dominate at lower frequencies. At lower frequencies one does not detect any compact feature in the source since the core contributes very little to the flux density."
    },
    "9803/astro-ph9803322_arXiv.txt": {
        "abstract": "We have determined detailed radio spectra for 26 compact sources in the starburst nucleus of M82, between $\\lambda\\lambda$ 74 and 1.3 cm.  Seventeen show low-frequency turnovers.  One other has a thermal emission spectrum, and we identify it as an \\HII region.  The low frequency turnovers are due to absorption by interstellar thermal gas in M82.  New information on the AGN candidate, 44.01+595, shows it to have a non-thermal falling powerlaw spectrum at the highest frequencies, and that it is strongly absorbed below 2 GHz.  We derive large magnetic fields in the supernova remnants, of order 1-2 $(1+k)^{2/7} \\phi^{-2/7}$ milliGauss, hence large pressures in the sources suggest that the brightest ones are either expanding or are strongly confined by a dense interstellar medium. From the largest source in our sample, we derive a supernova rate of 0.016 yr$^{-1}$. ",
        "introduction": "Supernovae (SNe) and supernova remnants (SNR's) are thought to be the main ``drivers'' of the starburst phenomenon (\\cite{kw75}). Detailed radio observations for supernovae and supernova remnants are few, due to the fact that only a small subset of the extragalactic SNe discovered each year produce detectable radio emission. Those events for which monitoring data is available (e.g. \\cite{weiler88}) have revealed considerable information regarding the evolution of the properties of the expanding shock wave and its emission processes.  In intense starbursts like M82, where a significant population of massive stars has had sufficient time to evolve to the SN stage, we have the unique opportunity to study a collection of SNR's of similar age and origin, and their interaction with the surrounding environment.  The observations are briefly outlined in Section 2.  We discuss the technique used to fit spectral models in Section 3.  Then in Section 4 we discuss the physical implications of our modelling. Throughout the paper, we assume that the distance to M82 is 3.63 Mpc (\\cite{freedman}, \\cite{tammann}), so that 1$\\farcs$0 corresponds to 17.6 pc. ",
        "conclusions": "  1. Of the 26 sources, one (42.21+590) is an \\HII region, 22 are most likely young radio supernovae or supernova remnants, and three further sources remain of uncertain nature.  2. High magnetic field strengths and pressures are derived for the sources we identify as SNR's.  The brighter ones do not appear to be in pressure equilibrium with the surrounding medium, and are probably expanding, or are strongly confined by the high pressure ISM. The source of the high nuclear pressure and halo winds are most likely the SNR's.  3. Sources that show a low-frequency turnover, do so as a result of absorption by ionized gas.  This ionized gas may be intimately associated with the source, or may be in clouds located along the line of sight to the sources. The gas is most likely in the form of clumpy, high density clouds.  4. The resolved ring-like source in our survey is the oldest. We derive a supernova rate in M82 of $\\sim$ 0.016 $\\left( \\frac{V}{5000 \\;\\;km/s} \\right)$yr$^{-1}$.  This value is a lower limit."
    },
    "9803/astro-ph9803114_arXiv.txt": {
        "abstract": "As the Cosmic Microwave Background (CMB) radiation is observed to higher and higher angular resolution the size of the resulting datasets becomes a serious constraint on their analysis. In particular current algorithms to determine the location of, and curvature at, the peak of the power spectrum likelihood function from a general $N_{p}$-pixel CMB sky map scale as $O(N_{p}^{3})$. Moreover the current best algorithm --- the quadratic estimator --- is a Newton-Raphson iterative scheme and so requires a `sufficiently good' starting point to guarantee convergence to the true maximum. Here we present an algorithm to calculate bounds on the likelihood function at any point in parameter space using Gaussian quadrature and show that, judiciously applied, it scales as only $O(N_{p}^{7/3})$. Although it provides no direct curvature information we show how this approach is well-suited both to estimating cosmological parameters directly and to providing a coarse map of the power spectrum likelihood function from which to select the starting point for more refined techniques. ",
        "introduction": "Planned observations of the Cosmic Microwave Background (CMB) will have sufficient angular resolution to probe the CMB power spectrum up to multipoles $l \\sim 1000$ or more (for a general review of forthcoming observations see \\cite{S}). If we are able to extract the multipole amplitudes $C_{l}$ from the data sufficiently accurately we will be able to obtain the values of the fundamental cosmological parameters to unprecedented accuracy. The CMB will then have lived up to its promise of being the most powerful discriminant between cosmological models \\cite{Kn,HSS,KKJS}.  Extracting the power spectrum is conceptually simple --- the raw data is cleaned and converted into a time-ordered dataset. This is then converted to a sky temperature map, which in turn is analysed to find the location of, and curvature at, the maximum of the likelihood function of the power spectrum. In practice as the size of the dataset increases the problem rapidly becomes intractable by conventional methods. This is particularly true of the final step --- the likelihood analysis of any reasonably general sky temperature map.  An observation of the CMB contains both signal and noise \\begin{equation} \\Delta_{i} = s_{i} + n_{i} \\end{equation} at each of $N_{p}$ pixels. For independent, zero-mean, signal and noise the covariance matrix of the data \\begin{equation} M \\equiv \\left< {\\bf \\Delta} \\, {\\bf \\Delta}^{T} \\right> = \\left< {\\bf s} \\, {\\bf s}^{T} \\right> + \\left< {\\bf n} \\, {\\bf n}^{T} \\right> \\equiv S + N \\end{equation} is symmetric, positive definite and dense. For any theoretical power spectrum $C_{l}$ we can construct the corresponding signal covariance matrix $S(C_{l})$; knowing the noise covariance matrix $N$ for the experiment we now know the observation covariance matrix for that power spectrum $M(C_{l})$. The probability of the observation given the assumed power spectrum is then \\begin{equation} \\label{eq.lf} P({\\bf \\Delta} \\, | \\, C_{l}) = \\frac{e^{-{\\small\\frac{1}{2}} \\, {\\bf \\Delta}^{T} M^{-1} {\\bf \\Delta}}} {(2 \\pi)^{N_{p}/2} \\left| M \\right|^{1/2}} \\end{equation} Assuming a uniform prior, so that \\begin{equation} P(C_{l} \\, | \\, {\\bf \\Delta}) \\propto P({\\bf \\Delta} \\, | \\, C_{l}) \\end{equation} we can restrict our attention to evaluating the right hand side of equation (\\ref{eq.lf}). Unfortunately unless the noise covariance matrix is unrealistically simple (eg. diagonal) both direct evaluation \\cite{G1,G2} and quadratic estimation \\cite{T,BJK} of the likelihood function scale as at best $O(N_{p}^{3})$ \\cite{B}, making them impractical for the forthcoming $10^4$ -- $10^6$ pixel datasets. ",
        "conclusions": "We have presented an algorithm to calculate probabilistic bounds on the power spectrum likelihood function from an $N_{p}$-pixel CMB map using Gaussian quadrature which scales as between $O(N_{p}^{7/3})$ and $O(N_{p}^{5/2})$ --- a very significant advance on existing algorithms for the exact calculation which scale as $O(N_{p}^{3})$. Since lowering the convergence constraint below the 1\\% level gains us only marginally tighter final bounds at the expense of increasing the scaling power, it is not recommended for the forthcoming $10^{4}$ -- $10^{6}$ pixel CMB maps. Our final algorithm of choice therefore gives better than 3\\% bounds on the logarithm of the likelihood function with $O(N_{p}^{7/3})$ operations with 99\\% confidence.  Since this algorithm gives no information about the local curvature of the likelihood function it is not as well suited as quadratic estimator techniques for searching a large multi-dimensional parameter space for its likelihood maximum. However for direct estimation of a small set of cosmological parameters this technique is certainly viable and fast. Moreover, even when the parameters are taken to be the multipole moments (individually or in bins), quadratic estimation, being a Newton-Raphson iteration, requires a starting point `sufficiently close' to the maximum to guarantee convergence; the algorithm presented here is then well suited to provide a coarse overall mapping of the likelihood function from which to select a starting point for more refined techniques."
    },
    "9803/astro-ph9803232_arXiv.txt": {
        "abstract": "We study the influence of a possible H dibaryon condensate on the equation of state and the overall properties of neutron stars whose population otherwise contains nucleons and hyperons. In particular, we are interested in the question of whether neutron stars and their masses can be used to say anything about the existence and properties of the H dibaryon. We find that the equation of state is softened by the appearance of a dibaryon condensate and can result in a mass plateau for neutron stars. If the limiting neutron star mass is about that of the Hulse-Taylor pulsar a  condensate of H dibaryons of vacuum mass $\\sim 2.2$ GeV and a moderately attractive potential in the medium could not be ruled out. On the other hand, if the medium potential were even moderately repulsive, the H,  would not likely exist  in neutron stars. If neutron stars of mass $\\sim 1.6 M_\\odot$ were known to exist, attractive medium effects  for the H could be ruled out. ",
        "introduction": "Since Jaffe proposed that there  may exist a stable dihyperon (a quark composite with baryon number two) \\cite{Jaffe77}, an ongoing  quest for this particle began \\cite{Carroll78}. Recent searches using kaon beams \\cite{Aoki90} or heavy ion beams \\cite{Belz96,Belz97,Stotz97} found no candidates or are still in progress \\cite{Craw98}. There exist some claims for evidence for the H dibaryon produced in proton-nucleus \\cite{Shahba93} and in heavy-ion collisions \\cite{Long95}. Nevertheless, these candidates might be misidentified $K^0_L$ as seen in \\cite{Belz97}. For a most recent overview on the search for the H dibaryon we refer to \\cite{HYP97}.  There are numerous mass estimates for the H dibaryon and they are reviewed in \\cite{Dover89}. The existence or nonexistence of the H dibaryon is strongly connected with the observation of double $\\Lambda$ hypernuclei which has been discussed in \\cite{Dalitz89}. Three double $\\Lambda$ hypernuclei have been reported in literature: $^{~~6}_{\\Lambda\\Lambda}$He \\cite{Danysz63}, $^{~10}_{\\Lambda\\Lambda}$Be \\cite{Prowse66}, and $^{~13}_{\\Lambda\\Lambda}$B \\cite{Aoki91,Dover91}. The two $\\Lambda$'s can decay by strong interactions to the H dibaryon. As this has not been seen in the above events, the H must either be heavier than $m_H> 2 m_\\Lambda+B_{\\Lambda\\Lambda}\\approx 2.22$ GeV \\cite{Kerb84} or the events are misidentified as an H hypernucleus with a shallow attractive nuclear potential \\cite{Dover89}. A more stringent condition is the observation of the weak mesonic decay of the double $\\Lambda$ hypernuclei giving $m_H> m_\\Lambda + m_p + m_{\\pi^-} + B_\\Lambda \\approx 2190$ MeV \\cite{Jaffe91} where $B_\\Lambda$ depends on the mass of the decay fragment and is $B_\\Lambda=-3.1$ MeV for $^5_\\Lambda$He. In all cases, a deeply bound H dibaryon seems to be ruled out by these events.  If the H dibaryon exists, it will have a certain impact also on the properties of dense matter. It is quite established nowadays, that neutron stars have a large hyperon fraction in the core and  might be described as giant hypernuclei, though bound by gravity \\cite{Glen85}. Here again, the presence of hyperons might restrict certain properties of the H dibaryon. Recently, studies for neutron stars have been done for nuclear matter without hyperons but including H dibaryon condensation \\cite{Fae97a} and limits have been set for the coupling constants of the H dibaryon \\cite{Fae97b}.  There might exist heavier partners of the H dibaryon, lumps of strange quark matter dubbed strangelets. There are several heavy-ion experiments dedicated to search for this novel form of matter \\cite{Arm97,Beavies95,Apple96}. In the MIT bag model, strangelets with $A\\leq 6$ are found to be unbound \\cite{Aerts78}. Nevertheless, light strangelet candidates in the range of $6<A<40$ might be stable against weak hadronic decay \\cite{Gilson93,Scha97} (an overview of the properties of strange matter can be found in \\cite{GS96}). The H dibaryon as well as these light strangelets can occur in dense matter as a precursor of the phase transition to a quark plasma.  In this paper, we study the influence of H dibaryons and other strangelet candidates on the composition and structure of neutron stars including the hyperon degree of freedom. We are particularly interested in the question of whether neutron stars and their masses can be used to say anything about the existence and properties of the H dibaryon. In section \\ref{sec:comp}, we discuss the condition for the occurrence of dibaryons and strangelets in neutron star matter. The relativistic mean field model with hyperons and the H dibaryon is presented in section \\ref{sec:model}. Implications for a H dibaryon condensate are discussed in section \\ref{sec:results} and summarized in the last section. ",
        "conclusions": "\\label{sec:summary}  We are particularly interested in the question of whether neutron stars and their masses can be used to say anything about the existence and properties of the H dibaryon. We have studied the influence of the possible occurrence of an H dibaryon condensate and strangelets in neutron stars including hyperons. Without in-medium modifications, it is quite likely that especially negatively charged strangelets, if they exists, will be present in the dense interior of neutron stars.  The appearance of  H dibaryons in the stellar core depends crucially on their mass and on the chosen potential of the H in nuclear matter. Hyperons tend to shift the onset of the H to higher density or to prevent H dibaryon condensation. If the condensation happens and if the potential of the H is attractive enough to provide a substantial number density in the neutron star, the maximum mass of the neutron star is reduced compared to the case without the H dibaryon. The decrease of the maximum mass is moderate and allows for the presence of H dibaryons in the interior of neutron stars in accord with present neutron star mass data. If the H dibaryon feels an attractive potential in matter, it can lead to a plateau in the mass of neutron stars, as there exist a region of very slowly rising pressure with energy density.  If the limiting neutron star mass is about that of the Hulse-Taylor pulsar a  condensate of H dibaryons of vacuum mass $\\sim 2.2$ GeV and a moderately attractive potential in the medium could not be ruled out. On the other hand, if the medium potential were even moderately repulsive, the H,  would not likely exist  in neutron stars. If neutron stars of mass $\\sim 1.6 M_\\odot$ were known to exist, attractive medium effects could be ruled out. For a mass limit of $1.44 M_\\odot$, attractive potentials for an H mass below the $\\Sigma$N threshold (1.3 GeV) are ruled out.  H dibaryon or strangelet condensation might happen as a precursor to the phase transition to a quark plasma. In this respect, we note that this phase transition is of first order \\cite{Glen92}. Hence, small bubbles of strange quark matter will appear in the mixed phase which are most likely negatively charged due to the isospin potential of the nuclear matter. This is in line with the results presented here. As the most stable strangelets have spin zero, the onset to a quark plasma will be initiated by a Bose condensation of strangelets (possibly including the H dibaryon). As the phase transition proceeds, the bubbles will overlap and will finally replace nuclear matter by essentially filling up the whole volume.   ~~\\\\[2ex] Acknowledgments:  J.S.B. acknowledges support by the Alexander-von-Humboldt Stiftung with a Feodor-Lynen fellowship. This work is supported by the Director, Office of Energy Research, Office of High Energy and Nuclear Physics, Nuclear Physics Division of the U.S. Department of Energy under Contract No.\\ DE-AC03-76SF00098."
    },
    "9803/astro-ph9803142_arXiv.txt": {
        "abstract": "Two fundamental empirical laws have been established in the analysis of galaxy space  distribution. First, recent analyses have revealed that the three dimensional distribution of galaxies and clusters is characterized by large scale structures and huge voids: such a distribution shows fractal correlations up to the limits of the available samples. This has confirmed the earlier de Vaucouleurs power-law density - distance relation, now corresponding to a fractal structure with dimension $D \\approx 2$, at least, in the range of scales $ \\sim 1 \\div 200 \\; Mpc$ ($H_0 = 55 km/sec/Mpc$). An eventual cut-off towards homogenization has not been yet identified. Second, since Hubble's discovery, the linear redshift-distance law has been well established within $200 Mpc$ and also much deeper. The co-existence of these laws within the same scales is a challenge for the standard cosmology where the linear Hubble law is a strict consequence of homogeneity of the expanding universe. This puzzle is now sufficiently strong to raise doubts for the standard cosmology. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803138_arXiv.txt": {
        "abstract": "We present new {$v\\sin i$}~ measurements for 235 low-mass stars in the Pleiades. The differential rotational broadening has been resolved for all the stars in our sample. These results, combined with previously published measurements, provide a complete and unbiased rotation data set for stars in the mass range from 0.6 to 1.2{$M_{\\odot}$}. Applying a numerical inversion technique on the {$v\\sin i$}~ distributions, we derive the distributions of {\\it equatorial\\/} velocities for low-mass Pleiades members. We find that half of the Pleiades dwarfs with a mass between 0.6 to 1\\,{$M_{\\odot}$}~ have rotation rates lower than 10{\\,km\\,s$^{-1}$}. \\\\ Comparison of the rotational distributions of low-mass members between IC 2602/2391 ($\\approx 35$\\,Myr) and the Pleiades ($\\approx 100$\\,Myr) suggests that G dwarfs behave like solid-bodies and follow Skumanich's law during this time span. However, comparison between Pleiades and  older clusters --M34 ($\\approx 200$\\,Myr) and Hyades ($\\approx 600$\\,Myr)--  indicates that the braking of slow rotators on the early main sequence is weaker than predicted by an asymptotical Skumanich's law. This strongly supports the view that angular momentum tapped in the radiative core of slow rotators on the zero age main sequence (ZAMS) resurfaces into the convective envelope between Pleiades and Hyades age. For the G-dwarfs, we derive a characteristic coupling time scale between the core and the envelope of about 100--200\\,Myr, which accounts for the observed evolution of surface rotation from the ZAMS to the Hyades.\\\\ The relationship between rotation and coronal activity in the Pleiades is in agreement with previous observations in other clusters and field stars.  We show that the Rossby diagram provides an excellent description of the X--ray activity for all stars in the mass domain studied. The Pleiades data for slow and moderate rotators fills the gap between the X-ray--rotation correlation found for slow rotators and the X-ray ``saturation plateau'' observed for young fast rotators.  The transition between increasing X-ray flux with rotation and X-ray saturation is observed at $\\log(P/\\tau)=0.8\\pm0.1$. These results strengthen the hypothesis that the ``saturation'' of the angular momentum loss process depends on the stellar mass. ",
        "introduction": "In the past 10 years, numerous rotational velocity measurements have been obtained for low-mass pre-main sequence stars and members of nearby young open clusters.  This set of observations has shown that the angular momentum of each star follows an evolutionary scenario from the star's birth to the age of the Sun.  Though several models have been developed to describe the rotational history of low-mass stars, the data available to confront models with observations were still insufficient to provide an unambiguous validation of these models. The measurements of stellar rotation for various ages and the understanding of its evolution shall provide a way to get a better knowledge of the physical processes experienced by the star through its history. In particular, this data might offer additional insight into the physics of young stars and their interactions with their proto-stellar and maybe their proto-planetary disks.  On the ZAMS, G type stars ($0.8-1.0${$M_{\\odot}$}) exhibit a large spread of rotational velocities. A difference of the order of 150{\\,km\\,s$^{-1}$}~ is observed for G dwarfs in Alpha\\,Per ($\\approx 50$\\,Myr) between the fastest and slowest rotators (Prosser 1992). At Hyades age (about 600\\,Myr), the spread has disappeared and the stars with masses in the 0.6 to 1 {$M_{\\odot}$}~ range show a monotically decreasing rotational velocity with mass (Radick et al. 1987; Stauffer et al. 1997a).  Over this time span, the fast rotators have been strongly braked but the slow rotators have only suffered a little braking. Moreover, the efficiency of the braking process also depends on the stellar mass. At the Pleiades age ($\\approx100$\\,Myr), almost all G type stars have converged down to slow rotation while K dwarfs stars still exhibit a large roational spread (Mayor \\& Mermilliod 1991; Soderblom et al. 1993).  At the Hyades age, a significant spread in rotational velocity spread is only observed for stars less massive than 0.6{$M_{\\odot}$} (Stauffer et al. 1997a).  Models have difficulties to simultaneously reproduce the large diversity of rotators at early ages and their strong convergence in a few 100\\,Myr years. The recent discovery of a bi-modal velocity distribution for pre-main sequence stars in Taurus and Orion (see Bouvier 1994; Choi \\& Herbst 1996) suggests that the magnetic coupling between the disk and the star (Camenzind 1990; K\\\"onigl 1991) could prevent the young star from spinning up during its PMS contraction, yielding a large spread in the initial angular momentum distribution.  On the main sequence, the star's rotation is mostly ruled by the angular momentum losses at the surface through a magnetized wind (Schatzmann 1962; Weber \\& Davis 1967) and by the angular momentum transfer in the inner parts of the star (Endal \\& Sofia 1978; Mac\\,Gregor \\& Brenner 1991).  Depending on the efficiency of the transfer and the disk lifetime, various scenarios can be predicted. The latest models of angular momentum evolution are from Krishnamurthi et al. (1997a), Bouvier et al. (1997b) and Allain (1998).  Numerous data on stellar rotation were available to these models, though not enough yet to derive the complete distributions of rotational velocity in young open clusters. The comparison between such distributions and the initial velocity distribution gathered from T\\,Tauri observations would provide an important additional constraint on the models.  Unbiased rotational velocity distributions are hard to build, numerous observations being needed to get a statistically significant distribution. In this case, measurements of the projected velocity ({$v\\sin i$}~) are a good alternative to the direct measurement of rotation periods which requires a huge number of photometric measurements, difficult to achieve for a large sample of stars. Indeed, since the pioneering work of van Leeuwen \\& Alphenaar (1982), there have been several attempts to measure photometric periods in the Pleiades (see O'Dell et al. 1995 for a census of the observations; Krishnamurthi et al. 1997b).  Up to now, 42 periods have been measured in this cluster, for stars with {(B$-$V)$_{0}$} in the range 0.5 to 1.4. The detection of rotational periods for the whole sample of Pleiades stars in this mass range (more than 200 catalogued members) is difficult and will still take a few years.  Measurements of {$v\\sin i$}~ are more straighforward but are obviously affected by projection effects. Yet, the knowledge of the {$v\\sin i$}~ for a large stellar sample does provide  rich {\\sl statistical} information. It is then possible to extract an accurate estimate of the distribution of {\\it equatorial} velocities in a cluster from a large number of observed {\\sl projected} rotational velocities if the stellar sample is statistically representative of the stellar population of the cluster and if it does not suffer from incompleteness (i.e., all the {$v\\sin i$}~ must be resolved). We consider that the inclination angles are randomly distributed in clusters.  In the next section we present new {$v\\sin i$}~ measurements for 235 Pleiades low-mass members.  We show that by taking into account the spectral type of stars, we can calibrate the intrinsic width of the spectral lines and measure the {$v\\sin i$}~ of very slow rotators down to 1.5{\\,km\\,s$^{-1}$}. The calibration process is described in Sect.~2. We detect rotational broadening for all observed stars. More than 98\\% Pleiades members of the original Hertzprung sample now have a resolved rotational velocity measurement up to {(B$-$V)$_{0}$}=1.35. The distributions of projected rotational velocities are then converted into distributions of equatorial velocities using a numerical algorithm described in Sect.~3.  The velocity distributions for various mass ranges are presented in Sect.~4. In Sect.~5 we compare our results with those obtained for other open clusters and discuss the implications for the models of angular momentum evolution. We also examine the relationship between X-ray flux and stellar rotation for Pleiades stars. ",
        "conclusions": "From a complete and unbiased set of {$v\\sin i$}~ measurements, we have computed the distribution of equatorial velocities in the Pleiades for various mass ranges between 0.5 and 1.5{$M_{\\odot}$}.  Comparison with the distribution of rotational velocities in the Hyades, M\\,34, IC\\,2391 and IC\\,2602 yields a coherent picture for the angular momentum evolution of the convective envelope.  The comparison with the younger clusters IC\\,2391 and IC\\,2602 suggests that most Pleiades G--dwarfs are in an unsaturated braking law regime.  The rotational evolution of moderate rotators in early stages is in agreement with a solid body rotation driven by Skumanich's relationship. The relationship between {$v\\sin i$}~ and X-ray emission indicates that the transition between saturated X-ray emission and its steady decrease with rotation occurs at about $P=2$\\,d or {$v\\sin i$}$=25${\\,km\\,s$^{-1}$} for solar-type stars. About 10\\% of Pleiades G dwarfs lie in the saturation domain.  The comparison with older cluster such as M\\,34 and the Hyades strongly suggests that angular momentum of slow and moderate rotators is transported from the fast rotating core to the convective enveloppe on a time scale of about 100--200\\,Myr on the early main sequence.  An alternative intepretation would call for intrinsic differences in the distribution of initial angular momenta from clusters to clusters, which is not currently supported by observations.  The advent, in the next decade, of multi-fiber spectrographs on large telescopes will offer the possibility of determining the {$v\\sin i$}~ distributions of more remote clusters of various ages and abundances as well as the rotational properties of very low-mass stars (0.1--0.5{$M_{\\odot}$}). This will undoubtedly improve our understanding of the rotational evolution of young stars and provide new clues to the physical mechanisms responsible for angular momentum transport in stellar interiors."
    },
    "9803/astro-ph9803248_arXiv.txt": {
        "abstract": "We have studied the variability of 6 low redshift, radio quiet `PG' quasars on three timescales (days, weeks, and months) using the ROSAT HRI.  The quasars were chosen to lie at the two extreme ends of the ROSAT PSPC spectral index distribution and hence of the H$\\beta$ FWHM distribution.  The observation strategy has been carefully designed to provide even sampling on these three basic timescales and to provide a uniform sampling among the quasars  We have found clear evidence that the X-ray steep, narrow H$\\beta$, quasars systematically show larger amplitude variations than the X-ray flat broad H$\\beta$ quasars on timescales from 2 days to 20 days. On longer timescales we do not find significant differences between steep and flat quasars, although the statistics are poorer.  We suggest that the above correlation between variability properties and spectral steepness can be explained in a scenario in which the X-ray steep, narrow line objects are in a higher $L/L_{\\rm Edd}$ state with respect to the X-ray flat, broad line objects.  We evaluated the power spectrum of PG1440+356 (the brigthest quasar in our sample) between $2\\times10^{-7}$ and $\\sim10^{-3}$ Hz, where it goes into the noise.  The power spectrum is roughly consistent with a 1/f law between $10^{-3}$ and $2\\times10^{-6}$ Hz.  Below this frequency it flattens significantly. ",
        "introduction": "\\bigskip   While X-ray variability in Seyfert galaxies has been the subject of intensive study (see review by Mushotzky, Done \\& Pounds 1993, and Green et al 1993, Nandra et al. 1997), the variability of quasars, being fainter, have received far less attention (the Zamorani et al. 1984 study remains the most extensive to date).  Higher luminosity AGN had been expected to be physically larger and so have slower, and most likely lower amplitude, variations. A typical quasar is 100--1000 times more luminous than the highly variable Seyferts, such as NGC~4051, and so would vary at a significantly lower rate, i.e. a minimum of several days. It was probably this expectation that deterred extensive observing campaigns.  This is unfortunate since we show here that X-ray variability is common, rapid and of quite large amplitude ($\\gs$ factor of 2). The variability most likely originates in the innermost regions of quasar, and so can help unravel the basic parameters of the quasar central engine (mass, geometry, radiation mechanisms, radiative transfer) none of which are yet well constrained.  These speculations are supported by compact galactic sources, for which the investigation of the ``variability properties vs. spectral shape'' and the ``variability properties vs. luminosity'' planes have brought a great improvement in our understanding (see e.g. van der Klis 1995). The same is likely to happen for quasars.  Compact galactic sources are usually bright and variable on timescales as short as 1 ms, which means that their variability timescales can be probed by just a few observations of individual objects.  For example, the Galactic black hole candidate (BHC) Cyg X-1 recently underwent two dramatic changes in its variability and spectral properties (see e.g. Cui et al. 1997), and it was possible to follow the whole cycle from the usual low and hard state, to a medium-high, softer state, and back to the low and hard state in the course of a few months.  This is unlikely to happen in a luminous AGN, even if there were similarities between AGN and black hole candidates, because the variability timescales (and the sizes and luminosities) are probably much longer (greater and higher) in AGN.  Instead, if the analogy between quasars and BHC holds, we will observe a population of quasars some in a `high and soft' state, and some in a `low and hard' state.  Hence the analogous observational way forward in quasars is to investigate the variability properties of samples of quasars, selected according to their spectral properties and luminosity.  The evidence for rapid, large amplitude, X-ray variability in a few AGN with unusually steep X-ray spectra and unusually narrow Balmer lines (mostly Narrow Line Seyfert 1 galaxies, NLSy1, given the strong correlation between these two quantities found by Laor et al., 1994, 1997, Boller, Brandt \\& Fink 1996) has recently grown:  \\begin {enumerate} \\item NGC4051 shows large variations (50 \\%) on timescales of $\\approx 100$ seconds and has very narrow optical and UV emission lines and steep 0.1-2 keV X-ray spectrum. It is also highly variable in the EUV (a factor of factor of 20 in 8 hours, Fruscione et al 1998).  \\item IRAS13224-3809 has a particularly steep 0.1-2 keV spectrum ($\\alpha_X>3$) narrow H$\\beta$ line, very strong FeII emission and shows X-ray variations of a factor 50 or more on timescales of a few days (Otani 1995, Brandt et al. 1995, Boller et al 1997) and a factor of 2 in less than 800 seconds (Boller et al. 1996);  \\item RE J 1237+264 showed one very large (factor of 50) variability event (Brandt, Pounds \\& Fink, 1995)  \\item PHL1092 varied by a factor of 4 in 2 days (Forster \\& Halpern 1996, Lawrence et al. 1997). PHL1092 is a high luminosity ($5\\times10^{46}$ erg s$^{-1}$) very steep 0.1-2 keV spectrum and narrow emission line quasar.  \\item The relatively high luminosity quasar NAB0205+024 ($L_{0.5-10 keV}=8\\times10^{44}$ erg s$^{-1}$) shows variations of a factor of 2 in less than 20 ks in an ASCA observation (Fiore et al. 1998a).  \\item Mark 478 (PG1440+356) shows factor of 7 variations in 1 day during a long EUVE monitoring (Marshall et al 1996)   \\item The extremely soft NLSy1 WPVS007 ($\\alpha_{0.1-2keV}=7.3$) showed a huge variation (factor of 400) variation between the RASS observation and a follow-up PSPC pointed observation taken 2 years later (Groupe et al. 1995).   \\end{enumerate}  At this point we believe that a systematic study of a sample of normal quasars spaning the observed range of properties is needed.  A systematic study of the variability properties of AGNs is not an easy task, since it requires:  \\begin {enumerate} \\item the selection of well defined and representative (i.e. unbiased) sample;  \\item the availability of numerous repeated observations;  \\item a carefully designed observational strategy. In particular, it is very important that the sampling time is regular and that it is similar for all objects in the sample, so that the results on each object can be easily compared with each other. Otherwise the differences in the observed variability may be induced just by differences in sampling patterns.  \\end{enumerate}  Because of these stringent requirements systematic studies are still lacking (for example the Zamorani et al. 1984 study on quasars, and the Green et al. 1993 study on Seyfert galaxies do not fulfill any of the three criteria above).  To explore in a systematic manner the possible relation between line width, soft X-ray spectrum and X-ray variability properties, we initiated a pilot program using the ROSAT HRI.  In the following we present the results from a campaign of observations of six PG quasars which addresses the above points. ",
        "conclusions": "In our sample of six PG quasars we found clear evidence that X-ray steep quasars show larger amplitude variations than X-ray flat quasars on timescales from 2 days to 20 days. On longer timescales we do not find significant differences between steep and flat quasars, although the statistics are poorer.  While the sample in this pilot study is small (and expanded monitoring of a larger sample of quasars is surely needed), the distinction is so clean cut that we feel justified in speculating on its origin.  Large narrow band flux variability in sources with a steep spectrum could be induced by small changes in the spectral index without large variations of the spectrum normalization. If this is the case each source should show large spectral variability, with $\\alpha_X$ anti-correlated (correlated) with the observed flux if the pivot is at energies lower (higher) than the observed band. Spectral variability of this kind has not been observed in the PSPC observations of these and other quasars (Laor et al. 1997, Fiore et al. 1994) and in NLSy1 (Boller et al 1997, Fiore et al. 1998a, Brandt et al. 1995).  Although our HRI observations cannot directly rule out this possibility, we therefore conclude that the observed temporal variability pattern is not likely the result of complex spectral variability.  Laor et al (1997) suggest that a possible explanation for the remarkably strong $\\alpha_X$--$H_{\\beta}$ FWHM correlation is a dependence of $\\alpha_X$ on $L/L_{\\rm Edd}$. The line width is inversely proportional to $\\sqrt {L/L_{\\rm Edd}}$ if the broad line region is virialized and if its size is determined by the central source luminosity (see Laor et al. 1997, \\S 4.7).  So narrow-line, steep (0.1-2 keV) spectrum AGNs emit close to the Eddington luminosity and have a relatively low mass black hole.  A similar dependence of spectral shape on $L/L_{\\rm Edd}$ is seen in Galactic black hole candidates (BHC) as they change from the `soft-high' state to the `hard-low' state.  A physical interpretation for this effect, as described by Pounds et al. (1995), is that the hard X-ray power-law is produced by Comptonization in a hot corona and that as the object becomes more luminous in the optical-UV, Compton cooling of the corona increases, the corona becomes colder, thus producing a steeper X-ray power-law. In BHC in `soft-high' states this power law component emerges above $\\sim 10$ keV, while the spectrum below this energy is dominated by softer emission often associated with optically thick emission from an accretion disk. In this model quasar disc emission is at a too low an energy to observe, since the disc temperature scales with the mass of the compact object as $M_{BH}^{-1/4}$, leaving the Comptonized power law to dominate the 2-10 keV spectrum. The 0.2-2 keV quasar emission could be due to Comptonization (see e.g. Czerny \\& Elvis 1987 and Fiore et al. 1995) by a second cooler gas component. The 0.2-2 keV and 2-10 keV spectral indices might well be correlated with each other, if the emitting regions are connected or if the emission mechanisms know about each other. We have undertaken a campaign of observations of the quasars in the Laor et al. (1997) sample with ASCA and BeppoSAX to clarify this point. Preliminary results (Fiore et al. 1998b) shows that this is indeed the case: steep $\\alpha_X$(PSPC) quasars tends to have a steeper hard (2-10 keV) X-ray power-law (although the spread of the 2-10 keV indices seems smaller than that of the PSPC indices (as also found by Brandt et al. 1997).   In a sample of quasars with similar luminosities, those emitting closer to the Eddington luminosity will also be those with the smaller black hole and hence smaller X-ray emission region.  Thus light travel time effects would smear intrinsic X-ray variability up to shorter time scales in high $L/L_{Edd}$ objects, compared with low $L/L_{Edd}$ objects.  Based on this interpretation, and on figures \\ref{stru_med} and \\ref{stru_mean} it appears that the emission region of the steep soft X-ray quasars is a factor of $\\approx 10$ smaller than that of flat soft X-ray quasars (about $\\approx 10^{16}$ cm and $10^{17}$ cm respectively). Alternatively, the higher variability of steep spectrum objects may be a true intrinsic property, which could be induced by some increased instability in high $L/L_{Edd}$ objects. The range of luminosity in our sample is too small to tell if the variability amplitude appears to be better correlated with black hole mass, as expected for light travel time smearing, or with $L/L_{Edd}$, which would indicate an intrinsic mechanism.   It is interesting to note that a completely different analysis, based on the interpretation of the optical to X-ray spectral energy distribution in terms of emission from accretion disks also suggest a small black hole mass and an high arretion rate in two NLSy1 (Siemiginowska et al. 1998).   It is also interesting to note that Cyg X-1 in the ``high and soft'' state (Cui et al. 1997) shows a total root mean square variability higher than that measured during periods of transitions and in the ``low and hard'' state.  This is due to strong 1/f noise, extending down to at least a few $10^{-3}$ Hz, when the source is the ``high and soft'' state (Cui et al. 1997).  When the source is in the ``low and hard'' state this 1/f noise is not present (see e.g. the review of van der Klis 1995).  Ebisawa (1991) found in a systematic study of Ginga observations of 6 BHC that the time scales of variability for the soft and hard components are often different.  The soft component is usually roughly stable on time scales of 1 day or less, while the hard component exhibits large variations down to msec time scales.  If the time scales of BHC and quasars scale with the mass of the compact object the above two time scales translate to $10^4$ years and 0.1 day respectively for our quasar sample. For a sample of quasars with similar redshifts and luminosities this predicts a rather small scatter in the soft component and a bigger scatter in the hard component.  ROSAT results go in this direction.  Laor et al. (1994, 1997) find that (for their sample of 23 low--z PG quasars) the scatter in the normalized 2 keV luminosity is significantly larger than that in the 0.3 keV luminosity.  We conclude that the analogy between AGN and Galactic BHC seems to hold qualitatively for their X-ray variability properties.  An alternative and intriguing possibility to explain the correlation between X-ray variability amplitude and spectral shape is that a component generated closer to the black hole dominates the emission of steep $\\alpha_X$ quasars, as in the spherically converging optically thick flow proposed by Chakrabarti and Titarchuk (1995) to characterize BHC in the high and soft state.  If this is the case then we would again expect that the spectrum of the steep $\\alpha_X$ quasars remain steep above 2 keV (and up to $m_ec^2$ according to Chakrabarti and Titarchuk). Observations with the ASCA and BeppoSAX satellites instruments, which are sensitive up to 10 keV, (Brandt et al 1997, Fiore et al 1998b, Comastri et al 1998) suggest that this is indeed the case.  Schwartz and Tucker (1988) suggested that AGN emission can be produced by an ensemble of acceleration sites (i.e. shock waves) with different electron spectral indices and therefore emitting power laws with different photon indices. In this picture the spectrum is a quadratic function in log E and the mean index at a given energy arises from the greatest number of individual acceleration sites. Variability would be greatest for both steepest energy index and flattest energy index sources, each of which are dominated by fewer individual regions.  This is contradicted by the present observations, unless the Soft X-ray flat quasars become still flatter (and more variable) above 2 keV. High energy X-ray spectroscopy and variability studies are again needed to obtain a definitive answer. A RossiXTE monitoring campaign of 4 PG quasars, with a sampling similar to that used in the HRI campaign, is in progress and could allow us to clarify this point."
    },
    "9803/astro-ph9803295_arXiv.txt": {
        "abstract": "We have searched for molecular gas towards the nucleus of four galaxies known to harbor a water vapor megamaser. CO(1\\raw0) emission of NGC 2639 and NGC 5506 was strong enough to allow us to map their inner regions. Weak emission from Mrk 1210 was detected and Mrk 1 was not detected at all. We report the tentative detection of the CO(2\\raw1) line in NGC 5506. After this work, 12 of the 18 known galaxies harboring a water vapor megamaser have been observed in CO.  The molecular gas content in the inner regions of water megamaser galaxies ranges from 5$\\times$$10^7$ to 6$\\times$$10^9~\\Msun$. The circumnuclear molecular gas surface density also extends over nearly two orders of magnitude. The maser luminosity is correlated neither with the total amount of molecular gas found in the inner few kpc of these galaxies nor with global properties of the molecular gas such as surface density or filling factor; it is also independent of the infrared and optical luminosities. The only significant correlation we have found involves the maser luminosity and the low frequency radio continuum flux density. We conclude that the maser activity is intrinsically related to the energy of the active galactic nucleus whereas the intensity and even the presence of a water megamaser is independent of the molecular gas global properties such as the molecular gas content and surface density in the inner galactic regions.  We have also found a pos\\-sible anti\\-cor\\-re\\-lation between the molecular gas surface density and the rate of the megamaser variations. A higher molecular gas abundance in the inner region could lead to higher maser variability because of larger nuclear flux variations due to the more variable gas infall, and/or because of more frequent interactions of the pumping agent with molecular gas condensations. ",
        "introduction": "The first water megamaser was discovered towards the nucleus of NGC 4945 (dos Santos \\& Lepine 1979). Strong emission at 22 GHz was detected, with a luminosity about 100 times higher than that of W49, the most powerful galactic water maser. Until 1985 four additional water megamasers were discovered, towards the nuclei of the Circinus Galaxy, NGC 1068, NGC 4258 and NGC 3079. During the following decade no more extragalactic water megamasers were found. In 1994 a survey towards active galactic nuclei was carried out by Braatz et al. (1994), leading to the discovery of five new water megamasers, doubling the number known to that date. Recently, Braatz et al. (1996) have reported the discovery of 6 additional water megamasers towards active galactic nuclei. The last ones up to now have been found in NGC 5793 (Hagiwara et al. 1997) and in NGC 3735 (Greenhill et al. 1997b).  \\begin{table*}[t] \\begin{center} \\caption{Adopted parameters for the observed water megamaser galaxies} \\begin{tabular}{lcccc} \\hline Parameter & NGC 2639 & NGC 5506 & Mrk 1 & Mrk 1210 \\\\ \\hline \\hline $\\alpha_{2000}$$^a$ & \\hms[8 43 38.0] &\\hms[14 13 14.9] &\\hms [1 16 7.25] & \\hms[8 4 6.0 ] \\\\ $\\delta_{2000}$$^a$ & \\gms[50 12 20 3] &\\gms[$-$3 12 26 7] &\\gms[33 5 22 2] & \\gms[5 6 50 4] \\\\ Morphological type$^a$  & SA(r)a & SA pec sp & S & S \\\\ Activity$^a$  & LINER & Sy 2 &Sy 2 &Sy 2  \\\\ Heliocentric velocity (\\kms)$^a$  & 3236 &1816 & 4824 & 4043 \\\\ \\VLSR (\\kms) & 3235 &1825  & 4825 & 4030 \\\\ Distance (Mpc)$^b$ & 44 & 24 &64  &54  \\\\ Position angle (deg)$^c$  & 140 & 91 & --- & ---\\\\ Inclination (deg) & 58 & 80 & 56 & 3 \\\\ $D_{25}$ (arcmin)$^c$ & 1.8 & 2.8 & 0.8 & 0.8 \\\\ Linear scale (pc arcsec$^{-1}$) &209 &118 &312 & 260\\\\ $\\log L_{\\rm IR}$ (\\Lsun)$^b$& 10.21  &10.35  & --- & 10.58\\\\ $\\log L_{\\rm FIR}$ (\\Lsun)$^b$& 9.95  &  9.85 &  ---& 9.85    \\\\ \\hline \\multicolumn{5}{l}{$^a$ Obtained from the NASA Extragalactic Database (NED)}\\\\ \\multicolumn{5}{l}{$^b$ See more information in the text (Sect. 2)}\\\\ \\multicolumn{5}{l}{$^c$ Obtained from the RC3 catalog (de Vaucouleurs et al. 1991)}\\\\ \\end{tabular}  \\end{center} \\end{table*}  Water megamasers, unlike the galactic and normal extragalactic masers, have been detected at the nuclei of distant galaxies. Interferometric observations (Claussen \\& Lo 1986; Greenhill et al. 1995a; Greenhill et al. 1996) have shown that megamaser emission comes from within a region around the galactic nucleus whose radius is in general smaller than 1 pc. Another important point is that all the galaxies harboring a water megamaser present some level of activity. This fact has led to suppose that the maser emission mechanism is identical to galactic masers. The huge difference in the energy involved is explained in terms of the pumping source: the energy source of the megamasers is the central object of the active nucleus, a much more powerful source than the central star powering the galactic masers.  Braatz et al. (1997) have recently examined the conditions for detectability of water megamasers in terms of a variety of properties of the active galaxies, but not including the molecular gas content. It is known (Heckman et al. 1989) that Seyfert 2 have abnormally large quantities of dust and gas if compared with Seyfert 1 galaxies. On the other hand LINER galaxies are thought to be simply extensions of Seyfert 2 galaxies to lower luminosities, photoionized by a weaker AGN spectrum (Osterbrock et al. 1993). Thus, the fact that no water megamasers have been found in Seyfert 1 nuclei, while all of them are in either Seyfert 2 or LINER galaxies seems to indicate that a high concentration of molecular gas in the inner regions of the galaxies is a key parameter for the megamaser emission to be produced.  To check this hypothesis and to find if the properties of the masers are related to those of the molecular gas, we have carried out a study of the molecular gas content and its properties in several of the water vapor megamasers for which no previous data existed or were incomplete. The megamasers we observed are four of those discovered by Braatz el al. in 1994: NGC 2639, NGC 5506, Mrk 1 and Mrk 1210. In our analysis we have used the available CO data for other water megamaser galaxies found in the literature (Heckman et al. 1989; Planesas et al. 1989; Sahai et al. 1990; Aalto et al. 1991; Wang et al. 1992 and Young et al. 1995).  \\begin{figure*}[t]  \\vspace{6cm} \\special{psfile=ms7133f1.ps hoffset=-40 voffset=-456} \\caption{CO(1\\raw0) emission profiles towards the central region ($23\\as$ diameter) of the three detected water megamaser galaxies. Velocity resolution is $21 \\kms$. The two NGC objects show double peak spectra, which may indicate the existence of a molecular gas ring surrounding the central AGN at a kpc scale. The spectrum of Mrk 1210 is single-peaked, a possible indication that the expected molecular gas ring is seen face-on. This spectrum was obtained in May 1997} \\end{figure*} ",
        "conclusions": "We have searched for molecular gas towards the nucleus of four galaxies known to harbor a water vapor megamaser, and detected the CO(1\\raw0) emission in three of them and the CO(2\\raw1) emission in one. With this work 12 of the 18 known water megamaser galaxies have been observed in CO, and only the most distant of the observed ones, Mrk 1, has not been detected yet.  \\begin{enumerate}  \\item The 12 water megamaser galaxies with molecular gas data available are not an homogeneous set regarding their molecular gas properties. The amount of H$_2$ in their circumnuclear regions ranges from 5$\\times$$ 10^7$ to 6$\\times$$ 10^9\\ \\Msun$. The extreme values of the H$_2$ surface density, $\\Sigma_{\\rm H_2}$, for the central kpc are 6 and 280 \\Msun\\ pc$^{-2}$. This parameter extends over a range of 2 orders of magnitude, a range similar to that of Seyfert galaxies, starburst galaxies or luminous infrared galaxies. The maser luminosity, $L_{\\rm maser}$, is not correlated to the total molecular gas mass. Therefore it seems that the total amount of molecular gas in the inner few kpc is not a fundamental parameter on which depends the existence and the average intensity of the water megamaser.  \\item $L_{\\rm maser}$ is not correlated with $\\Sigma_{\\rm H_2}$ or to the filling factor of giant molecular clouds. Apparently the maser luminosity does not depend on the content of molecular gas in the inner kpc. However, the accumulation of clouds of dense molecular gas is believed to be necessary for the generation of a water megamaser. Observations with higher angular resolution of the molecular gas in the inner regions would help to solve the issue.  \\item The only correlation we have found involving the maser emission and molecular gas parameters is between the rate of relative variation of the maser intensity and $\\Sigma_{\\rm H_2}$. This fact may indicate that a high abundance of molecular gas in the inner regions could lead to higher variability in the maser emission, on the one hand, due to the higher variability of the central pumping source produced by wider variations in the gas infall; on the other hand, due to the more frequent interactions of the pumping agent with molecular gas condensations.  \\item $L_{\\rm maser}$ is not correlated to any other luminosity (infrared, optical, X-ray, blue). However, we have found some correlation between $L_{\\rm maser}$ and the global radio continuum flux density at 1.4 and 8.4 GHz. This fact supports the idea that $L_{\\rm maser}$ is a property related to the galactic nucleus (characterized by the radio luminosities) rather than to the inner galactic regions (characterized by the infrared luminosity and the molecular gas content).    \\item Mrk 1210 stands as a peculiar object, having the highest star formation efficiency among water megamaser galaxies is spite of its relatively low molecular gas content.  \\end{enumerate}"
    },
    "9803/astro-ph9803156_arXiv.txt": {
        "abstract": "We study the FIR and UV-visible properties of star forming galaxies in the nearby Universe. This comparison is performed using the local luminosity functions at UV and FIR wavelengths and on individual starburst galaxies for which photometric data from UV to NIR and FIR are available.  The FIR and UV luminosity functions have quite different shapes~: the UV function exhibits a strong increase for low luminosity galaxies whereas the FIR tail towards ultra luminous galaxies ($\\rm L > 10^{11}  L\\odot$) is not detected in UV. The comparison of the FIR and UV local luminosity densities argues for a rather moderate extinction in nearby disk galaxies. The galaxies selected to be detected in FIR and UV are found to be located in the medium range of both luminosity functions.  An emphasis is made on starburst galaxies. For a sample of 22 of these objects, it is found that the UV (912-3650 $\\rm \\AA$), the visible (3600-12500 $\\rm \\AA$) and the NIR (12500-22000 $\\rm \\AA$) wavelength range contribute $\\sim 30\\%$,  $\\sim 50\\%$ and  $\\sim 20\\%$ respectively to the total emerging stellar emission (for a subsample of 12 galaxies for the NIR and  visible light). The mean ratio of the dust to bolometric luminosity  of these galaxies is 0.37$\\pm$0.22 similar to the ratio found for normal spiral galaxies. Only 4 out of the 22 galaxies exhibit a very large extinction with more than 60$\\%$ of their energy emitted in the FIR-submm range.  The mean extinction at 2000$\\rm \\AA$ is found to be $\\sim 1.2$ mag although with a large dispersion. The UV, visible and NIR emissions of our sample galaxies are  consistent with a burst lasting over $\\sim 1$ Gyr. The conversion factor of the stellar emission into dust emission is found to correlate  with the luminosity of the galaxies, brighter galaxies having a higher conversion factor. Since our sample appears to be representative of the mean properties of the galaxy  population in FIR and UV, a very large conversion of the stellar light into dust emission can no longer be assumed as a general property of starburst galaxies at least in the local Universe.  Instead a larger amount of energy emerging from the present starburst galaxies seems to come from the stars rather than from the dust.  We compare the UV properties of our local starburst galaxies to those  of recently detected high redshift galaxies. The larger extinction found in the distant galaxies is consistent with the trend we find for the nearby starburst galaxies namely the brighter the galaxies the lower the escape fraction of stellar light.  \\keywords { Galaxies: starburst --Galaxies: stellar content -- Infrared: galaxies-- Ultraviolet: galaxies -- dust, extinction} ",
        "introduction": "One of the most important challenge of modern astronomy is the detection of young primeval galaxies. Indeed, very significant progress has been made  with the detection of very high redshift galaxies either from ground based observations (Steidel et al. 1996b) or in the Hubble Deep Field (e.g. Steidel et al. 1996a, Lowenthal et al. 1997).  In order to understand the properties of high redshift galaxies and study the cosmological evolution of star-forming galaxies, it is crucial to properly characterize the properties of starburst galaxies in the local Universe. These nearby galaxies are forming stars with a very high rate and it it actually important to analyse their emission over the entire spectral range (from UV to FIR-submm) to study the efficiency of the dust extinction and know what is the spectral range (UV, visible, NIR or FIR) where most of their energy is emitted. If high redshift star forming galaxies are similar to their low-z counterparts studying the latter will bring some clues to detect the former. We can wonder whether the observation of the rest-frame UV continuum is the best way to detect high redshift galaxies or if the high obscuration from dust makes them emit more energy in the FIR.  IRAS discovered infrared bright galaxies with prodigious star formation rates, a very high extinction and therefore a low optical flux (Sanders \\& Mirabel 1996).  The most luminous FIR galaxies are produced by strong interactions or merging of molecular gas-rich galaxies which induce enormous starbursts. Such objects might be the progenitors of elliptical galaxies (Kormendy \\& Sanders 1992). In a \"bottom-up\" scenario of galaxy formation, numerous starbursts induced by merging are expected and the bulk of their emission would be in the FIR redshifted in the submm (e.g. van der Werf \\&  Israel 1996). Mazzei et al. (1994) predict that more than 90 \\% of the energy emitted by a starburst is in the FIR range during the first Gyr. This percentage rapidly drops to reach $\\sim 30\\%$ for $\\sim 5$ Gyr- old objects. Models aimed at explaining the galaxy counts in optical and FIR predict that during intense phases of star formation the quasi totality of the stellar light is absorbed and re-radiated in the FIR wavelength range (Franceschini et al. 1994, Pearson \\& Rowan-Robinson 1996).  Conversely, considerable effort has been carried out in the UV-optical study of star forming galaxies for some years. Calzetti, Kinney and co-workers have extensively used the IUE spectra of star-forming galaxies complemented with optical and IR data to characterize the star formation history and the extinction occurring in the central regions of these objects (Calzetti et al. 1994, Calzetti et al. 1995, Calzetti 1997a). Meurer et al. (1995) have used Faint Object Camera on board the {\\it Hubble Space Telescope} HST-FOC observations to study the morphology of some starburst galaxies. From these studies, a foreground distribution of dust and a rather grey extinction curve seems to be able to explain the spectral distribution of the central regions of starburst galaxies. The extinction found by Meurer et al. (1995) for nearby starburst galaxies is rather low: at 2000 $\\rm \\AA$ it lies between 0.08 and 1.9 mag (excluding NGC7552 at 3.13 mag). Nevertheless these studies deal with the central parts of starburst galaxies and may well not be valid for the global emission of these objects at least when longer wavelengths than UV are concerned (Buat et al. 1997).  Analyses of the UV-optical and FIR global emissions of nearby spiral and irregular galaxies selected to be observed both in UV and in FIR led to a rather low extinction (Xu \\& Buat 1995, Buat \\& Xu 1996, Wang \\& Heckman 1996). An important result of these studies is that the UV non-ionizing stellar emission is likely to be the major cause of dust heating. The contribution  of OB stars to the dust heating is estimated to amount to about 20$\\%$ of the total FIR emission in the Milky Way (Cox \\& Mezger 1989), almost the same contribution of the ionizing  radiation to the dust heating is found for spiral galaxies (Xu \\& Buat, 1995)  and for  starburst objects (Calzetti et al. 1995). The comparison  between the FIR emission (dust re-radiation) and the UV and optical emission  (escaped stellar light) constrains the extinction. As a consequence the FIR to  UV continuum ratio is a powerful indicator of the extinction occurring in  galaxies (Meurer et al. 1995, Buat \\& Xu 1996, Wang \\& Heckman 1996).  Recently, HST imaging of very high redshift galaxies complemented when possible by spectroscopic observations with the Keck Telescope have led to the discovery at high redshift (z$\\sim 2-3$) of compact star forming galaxies with a moderate size and a strong rest-frame UV emission (Steidel et al. 1996a) with sometimes more diffuse extended structures  Lowenthal et al. 1997). Depending on the intrinsic UV spectrum adopted, the  average extinction estimated at 1600 $\\rm \\AA$ from the rest-frame UV spectral  energy distribution of these galaxies is of the order of 1.7 to 3 mag (Meurer  et al. 1997, Calzetti 1997b). These significant average extinctions are  therefore larger than the values estimated for nearby starburst galaxies.  However, it must be noted that these high redshift galaxies are very luminous  ($\\rm M_B < -21$) when compared to the mean luminosity of nearby starburst  galaxies studied by Meurer et al. (1995) ($\\rm <M_B> = -18.6$) and the  extinction is known to correlate with the luminosity of galaxies  (Giovanelli et al. 1995, Wang \\& Heckman 1997). Moreover, the selection biases are very strong towards very luminous galaxies with a strong UV continuum and it cannot be excluded that high redshift galaxies almost entirely hidden by the dust are missing from these observations in the rest-frame UV (Mobasher et al. 1996, Burigana et al. 1997).  The selection biases in the recent detections of high z galaxies are difficult or even impossible to quantify in the absence of similar observations at other wavelengths corresponding to longer than UV rest frame emissions (NIR or FIR). A first step is to estimate the importance of such a bias in the local Universe. Such a study is also crucial to compare the properties of high redshift galaxies to those in the nearby Universe.  At this aim  we will adopt a global approach to study the local Universe which consists in comparing the luminosity functions and the luminosity densities in UV and FIR.  The UV wavelength range is particularly interesting since it is  a tracer of the recent star formation rate as already mentioned and  observations in the visible  range of high z ($>2$) galaxies correspond to their UV rest frame emission. More specifically we will compare the amount of energy locked up in FIR to the amount of energy directly emitted in UV in the local Universe. The comparison of these global values with  individual galaxies selected to be observed both in UV and FIR will allow to discuss how such samples of individual galaxies are representative of the mean properties of the local Universe.  We will also investigate the specific case of a sub sample of nearby starburst galaxies detected in UV, visible, NIR and FIR in order to compare their global dust and stellar emission and to estimate what fraction of the emission of stars is converted into dust emission as well as the relative contribution of the UV, visible and NIR spectral ranges to the observed stellar emission. Such estimates will lead to predict what spectral range is more favorable for the detection of high redshift starbursts under the hypothesis that they are similar to their nearby counterparts. The main limitation to this approach is that we deal with global fluxes integrated over the galaxies whereas the starburst often occurs in the central parts. Moreover the galaxies at high redshift so far detected seem to have compact morphologies. Nevertheless as it will be shown in section 4  a  large fraction of the UV emission of a starburst galaxy  is likely to come from the starburst itself making valid a study on the global fluxes as soon as this wavelength is concerned. It is also the case for the FIR emission since the UV (ionizing and non ionizing) emissions is the major contributor to the dust heating), especially in starbursting objects. Obviously, more care must be taken when dealing with  the visible and NIR emission: at these wavelengths  the contribution of the underlying old stellar population present in local starburst galaxies is very large even dominant. Endly, available FIR fluxes on large samples are integrated over the galaxies due to the poor resolution of the IRAS  satellite and dealing with global fluxes allows a reliable comparison of the emission of galaxies at different wavelengths.  Beyond the detection of high redshift star forming galaxies it is necessary to estimate a quantitative star  formation rate (SFR) for these galaxies. The deduction of such a quantity from  the observed rest-frame UV continuum relies almost entirely on the amount of  the extinction with only a moderate dependence on the star formation history  (e.g. Meurer et al. 1997, Calzetti 1997). More specifically  after $\\sim 5~10^7$ years of constant star  formation rate the UV flux (912-3650 $\\rm \\AA$) reaches 80 $\\%$ of its stationnary value calculated for $\\rm 10^{10}$ years of constant star formation (from Bruzual \\& Charlot  1993). From a  global energetic budget, we will try to constrain the amount of extinction and bring some clues to this difficult problem. ",
        "conclusions": "\\subsection{ The local Universe}  We have considered galaxies selected to be detected photometrically both in UV and FIR. An analysis of the local luminosity functions at both wavelengths shows that the galaxies selected in this way have a FIR and a UV emission representative of the mean population of galaxies in the local Universe. From a sample of 22 starburst galaxies, the mean escape fraction of the stellar light (ratio of the stellar luminosity to the bolometric one) is found to be $63\\%\\pm 22\\%$, i.e. similar to what is found for normal galaxies.  This escape fraction  exhibits a strong decrease with increasing galaxy luminosity ranging from  $\\sim 80\\%$ at faintest B magnitudes ($\\rm M_B~>-17$) to $\\sim 10\\%$ for $\\rm  M_B~\\sim~ -20$. This result is quite different to that of Pearson \\& Rowan-Robinson (1996) who found that the escape fraction of the stellar light in starburst galaxies cannot exceed 5-10$\\%$ from an analysis of deep counts in visible and NIR and assuming a very strong cosmological evolution for the starburst population. Such a low fraction of escaped light is consistent with the studies of FIR bright galaxies but seems not to be a generic property of starbursting objects, at least in the local Universe.  \\subsection {Comparison with high redshift galaxies}  We can compare our results found for nearby starbursts to high redshift galaxies recently detected. A limitation to this comparison might be that we deal with integrated fluxes on nearby galaxies  which contain an underlying stellar population pre-existing to the starburst.  To our knowledge,  there is no evidence for the   presence or not of such an older population in high redshift starburst  galaxies (z $\\sim$ 3). Nevertheless, as discussed in the paper, the UV range is  largely dominated by the emission of the newly formed stars and the comparison  of the high redshift  galaxies to present-day starburst ones  in this wavelength range  is justified.  From an analysis of the slope of the rest-frame UV continuum, an extinction has been estimated for  high z  galaxies (Meurer et al. 1997, Calzetti 1997b, Pettini et al. 1997). A mean extinction of 1.65 mag at 1650 $\\rm \\AA$ is found by Calzetti for a star formation rate constant over  $\\sim 1$ Gyr which implies an unreddened slope for the UV continuum $\\rm \\beta=-2$. This value is in agreement with Pettini et al.'s estimates. Meurer et al. have adopted a different star formation law with a starburst lasting 10 $\\rm Myr$, leading to a larger proportion of very young stars and a steeper UV slope ($\\rm \\beta=-2.5$). Such a very extreme scenario gives an extinction as large as 3 mag. at 1600 $\\rm \\AA$ which can be considered as a very upper limit. The extinction found for  nearby starburst galaxies is lower than the value found by Calzetti or Pettini et al. at high redshift. But the galaxies observed at high redshift are much more luminous that the nearby starburst galaxies studied by us or by Meurer et al. (1995) (Meurer et al. 1997, Lowenthal et al. 1997, Pettini et al. 1997). As an example, an extinction of 1.6 mag at 1600 $\\rm \\AA$ gives  $\\rm F_{dust}~/~F_{bol} = 0.6$ and $\\rm F_{dust}~/~F_{UV} = 4.5$ using the extinction curve of Calzetti et al. (1994), a foreground screen and a star formation rate constant over 1 Gyr. Extrapolating Fig.~5, these ratios correspond to the brightest galaxies with $\\rm M_B\\le -20$. If the trends found in Fig.5 are mainly due to the extinction,  the extinction estimated in high redshift star forming galaxies is in agreement with the values found for nearby starburst galaxies when the galaxy luminosity  is accounted for. We must point out, however, that  low luminosity galaxies ( $\\rm M_B \\sim -16$) would be detectable with a NGST-type telescope in a reasonable amount of time as far as $\\rm z\\sim 9$ (if they exist). In a hierarchical scenario for the formation of the galaxies, it is likely  that such galaxies would outnumber the large luminous galaxies  presently detected (Ellis 1997).  Obviously the knowledge of the local Universe properties in terms of the extinction occurring in galaxies as well as their spectral energy distribution, especially in the UV range, is essential to interpret the observations of the Universe at high redshift. It is also crucial to predict the best way to detect young galaxies during a  phase of intense star formation. We show that a comparison of the FIR and UV emissions in nearby galaxies can bring some clues since these two emissions are very sensitive to the current star formation rate and to the extinction. A  more straightforward method would be a systematic comparison of deep fields in UV and FIR. Such a preliminary comparison was already performed (Buat \\& Xu 1996) using FAUST (Deharveng et al. 1994) and IRAS observations on  two 7.6$\\rm ^o$-wide fields in the central region of the Virgo cluster. Although the UV observations  of FAUST were not very deep (limiting flux $\\rm 10^{14} - 10^{15}~ erg/cm^2/s/\\AA$ at 1600 $\\rm \\AA$), we have searched for  galaxies detected by IRAS without any UV detected counterpart. Given the low sensitivity of the FAUST experiment, the lower limits obtained for  the ratio of the FIR to UV luminosity are  within the range of  values found for the galaxies detected at both wavelengths (fig. 4  in  Buat \\& Xu 1996). This is in agreement with the results found in this paper but needs to be confirmed with   deeper observations. The UV observations carried out with the large field  ($\\rm \\sim 2^o$) FOCA telescope (e.g. Donas et al. 1990) reach the magnitude  18 at 2000$\\rm \\AA$. Unfortunately up to now the FIR observations (made by the  IRAS satellite) are not deep enough to be compared to these UV observations  and we have to wait for new FIR large field instruments like WIRE to perform  such a comparison."
    },
    "9803/astro-ph9803010_arXiv.txt": {
        "abstract": "Metal line ratios in a sample of 13 quasar spectra obtained with the HIRES spectrograph on the KeckI telescope have been analyzed to characterize the evolution of the metagalactic ionzing flux near a redshift of 3.  The evolution of \\ion{Si}{4}/\\ion{C}{4} has been determined using three different techniques: using total column densities of absorption line complexes, as in Songaila \\& Cowie\\markcite{sc96} (1996); using the column densities of individual Voigt profile components within complexes; and using direct optical depth ratios.  All three methods show that \\ion{Si}{4}/\\ion{C}{4} changes abruptly at $z \\sim 3$, requiring a jump in value of about a factor of 3.4, and indicating a significant change in the ionizing spectrum that occurs rapidly between $z = 2.9$\\ and $z = 3$, just above the redshift at which Reimers et al.\\markcite{reim} (1997) detected patchy \\ion{He}{2} ${\\rm Ly}\\alpha$\\ absorption.  At lower redshifts, the ionization balance is consistent with a pure power law ionizing spectrum  but at higher redshifts the spectrum must be very soft, with a large break at the He$^+$\\ edge.  An optical depth ratio technique is used to measure the abundances of ions whose transitions lie within the forest and \\ion{C}{3}, \\ion{Si}{3} and \\ion{O}{6} are detected in this way.   The presence of a significant amount of \\ion{O}{6} at $z > 3$\\ suggests either a considerable volume of \\ion{He}{3} bubbles embedded in the more general region where the ionizing flux is heavily broken, or the addition of collisional ionization to the simple photoionization models. ",
        "introduction": "\\label{intro}  Whereas we know from the absence of any significant Gunn-Peterson effect even in the highest redshift quasars that hydrogen reionization of the intergalactic gas must have taken place at $z > 5$, we have much less information about the period at which the bulk of singly ionized helium converted to doubly ionized helium.  Since late He$^+$\\ ionization may significantly change the temperature of the intergalactic gas it is critical to understand this heating if we are to correctly model the growth of structure in the IGM, and determine the mapping of the baryon density to observable quantities such as observations of the neutral hydrogen ${\\rm Ly}\\alpha$\\ forest and the He$^+$\\ ${\\rm Ly}\\alpha$\\ opacity. Phenomenological modelling of this event depends critically on the softness of the composite spectrum of the ionizing sources (e.g.\\ Miralda-Escud\\'e \\& Rees\\markcite{mr93} 1993; Madau \\& Meiksin\\markcite{mm} 1994), amplified by the subsequent radiative transfer, and such models cannot be considered reliable in predicting the high energy ($E > 54~{\\rm eV}$) metagalactic ionizing spectrum above the He$^+$\\ ionization edge, which determines the fraction of singly ionized helium.  Our most direct information on the He$^+$\\ opacity is through observations of quasars whose spectra extend to the He$^+$\\ ${\\rm Ly}\\alpha$\\ wavelength. Despite the extreme difficulty of these measurements, successful observations of the He$^+$\\ ${\\rm Ly}\\alpha$\\ absorption have been made toward $z > 2.8$\\ quasars with HST (Jakobsen et al.\\markcite{jak} 1994; Hogan et al.\\markcite{hog} 1997; Reimers et al.\\markcite{reim} 1997) and of the $z = 2.72$\\ quasar HS1700+6416 with HUT (Davidsen et al.\\markcite{dkz} 1996; Zheng et al.\\markcite{zdk} 1998).  The He$^+$\\ ${\\rm Ly}\\alpha$\\ opacity shows a marked decrease from a value of $\\tau = 3.2^{+\\infty}_{-1.1}$\\ in Q0302$-$003 at $z = 3.29$\\ to $\\tau = 1.0 \\pm 0.07$\\ in HS1700+6416.  More remarkably, the Reimers et al.\\markcite{reim} observation of the intermediate redshift quasar HE2347$-$4342 at $z = 2.89$\\ shows both `troughs' and `voids' in the \\ion{He}{2} ${\\rm Ly}\\alpha$\\ observations, suggesting that at $z \\sim 2.8$\\ we are seeing fully ionized \\ion{He}{3} bubbles interspersed among as yet unionized He$^+$\\ regions and that it is at this point that the porosity of \\ion{He}{3} regions is approaching unity.  The recent discovery that the bulk of ${\\rm Ly}\\alpha$\\ forest absorption with $N({\\rm H~I}) > 3\\times 10^{14}~\\rm cm^{-2}$\\ contains associated metals (Songaila \\& Cowie\\markcite{sc96} 1996, hereafter SC) gives us an alternative approach to the problem since ionization balance in these forest metals provides a diagnostic of the shape of the metagalactic flux in the neighborhood of the He$^+$\\ edge.  As was first noted in Songaila et al.\\markcite{shc} (1995), the value of \\ion{Si}{4}/\\ion{C}{4} in high ionization systems is critically dependent on the He$^+$\\ ionization edge break strength.  This has subsequently been investigated in more detail by SC, Savaglio et al.\\markcite{sav} (1997) and Giroux \\& Shull\\markcite{gs97} (1997), among others.  As will be discussed further here, the bulk of the forest metal line systems at the currently observed redshifts ($z \\sim 2 - 4$) are high ionization (\\ion{C}{2}/\\ion{C}{4} $\\ll 0.1$) and so, unless there is a strong break at the edge, they will have \\ion{Si}{4}/\\ion{C}{4} $\\ll 0.1$\\ even for the higher Si/C abundances characteristic of low metallicity systems.  Therefore, once the He$^+$\\ is fully ionized and the IGM becomes relatively transparent to the integrated quasar spectrum (e.g.\\ in the metagalactic spectrum of Haardt \\& Madau\\markcite{hm} 1996), the observed \\ion{Si}{4}/\\ion{C}{4} values should fall in this low range.  SC and Savaglio et al.\\markcite{sav} (1997) have shown that, whereas this is generally true at $z < 3$, much higher \\ion{Si}{4}/\\ion{C}{4} values are regularly seen at $z > 3$, suggesting a significant change above this redshift, which would be consistent with the interpretation of the Reimers et al.\\markcite{reim} (1997) observations as showing the redshift at which \\ion{He}{3} bubbles begin to overlap. Boksenberg\\markcite{bok} (1998) has recently questioned this result, based on an analysis of the the redshift evolution of the ion ratios in the separate Voigt profile components in complex systems rather than in the integrated column densities of the complexes.  However, his analysis is based on rather a small sample of systems without clear selection criteria.  In this paper, I shall use the largest sample to date to demonstrate unambiguously that, irrespective of the method of analysis, there is indeed a rapid jump in the value of \\ion{SI}{4}/\\ion{C}{4} at a redshift just below 3, and that the ionization stages in the metals are consistent with this being, for most systems, the point at which they change from being ionized by a metalgalactic spectrum that is, at $z > 3$, heavily broken above 54~eV to one that is only mildly broken at the lower redshifts.  The sample and the data reduction are described in \\S 2 and the reader who is primarily interested in the results could safely skip this section.  The evolution of the \\ion{C}{2}/\\ion{C}{4} and \\ion{Si}{4}/\\ion{C}{4} values with redshift is described in \\S 3 where I show that the presence of a jump in \\ion{Si}{4}/\\ion{C}{4} values at $z$\\ just under 3 is highly significant and does not depend on the method of analysis, whether by total column densities of complex, by the column densities of individual Voigt components, or by directly analyzing the distribution of optical depth ratios.  In \\S 4 I consider the overall ionization balance including intermediate ionization stages such as \\ion{C}{3} and \\ion{Si}{3} and high ions such as \\ion{N}{5} and \\ion{O}{6} using, for those lines that lie primarily within the forest, the optical depth distributions of the ensembles of such lines, which provides a new and robust technique for determining their properties.  The overall ionization balance is broadly consistent with unbroken power law photoionization at $z < 3$; however, the observation of significant amounts of \\ion{O}{6} at $z > 3$\\ requires either the presence of a considerable volume of \\ion{He}{3} bubbles permeating regions where the ionizing flux is heavily broken, or the addition of collisional ionization to the simple photoionization models.  Finally, the conclusions are briefly summarised in \\S 5. ",
        "conclusions": "\\label{conc}  I summarise the results of the paper by noting that, irrespective of the analysis methodology adopted, there is a significant change in the ionization balance of forest metal lines which occurs just below a redshift of 3.  At lower redshifts, the ionization balance in the forest lines is fully consistent with a pure power law ionization spectrum with an index of $-1.8$\\ but at higher redshifts the high values of \\ion{Si}{4}/\\ion{C}{4} seen in most of the forest clouds despite generally low \\ion{C}{2}/\\ion{C}{4} values implies that the ionizing flux must be very soft, with a large break at the He$^+$\\ edge.  The change occurs quite rapidly between $z = 2.9$\\ and $z = 3$, just above the redshift at which highly patchy \\ion{He}{2} ${\\rm Ly}\\alpha$\\ absorption is seen in the quasar HE~2347$-$4342 (Reimers et al.\\markcite{reim} 1997).  The simplest explanation seems to be that we are seeing the redshift at which \\ion{He}{2} ionizes completely to \\ion{He}{3} as the \\ion{He}{3} Str\\\"omgren spheres overlap."
    },
    "9803/astro-ph9803226_arXiv.txt": {
        "abstract": "The 6.4~s X-ray pulsar \\src\\ was observed by \\sax\\ in 1997 May. This source belongs to the class of ``anomalous'' pulsars which have pulse periods in range 5--11~s, show no evidence of optical or radio counterparts, and exhibit long-term increases in pulse period. The phase-averaged 0.5--10 keV spectrum can be described by an absorbed power-law and blackbody model. The best-fit photon index is 2.5$\\pm$0.2 and the blackbody temperature and radius are 0.64$\\pm$0.01~keV and $0.59 \\pm 0.02$~km (for a distance of 3~kpc), respectively. The detection of blackbody emission from this source strengthens the similarity with two of the more well studied ``anomalous'' pulsars, 1E\\,2259+586 and 4U\\,0142+614. There is no evidence for any phase dependent spectral changes.  The pulse period of $6.45026 \\pm 0.00001$~s implies that \\src\\ continues to spin-down, but at a slower rate than obtained from the previous measurements in 1994 and 1996. ",
        "introduction": "\\label{sec:introduction}  \\src\\ is an unusual pulsar, with a   period of 6.4~s and a soft spectrum, discovered during {\\it Einstein}\\/ observations of the Carina nebula (Seward et al.\\ 1983). The source has been spinning down for at least the last 17 years (Mereghetti 1995;  Corbet \\& Mihara 1997). The spectrum has been modeled by an absorbed power-law with a photon index, $\\alpha$, of $\\sim$2--3 (Seward et al.\\ 1983; Corbet \\& Mihara 1997). This spectral shape is softer than those of typical high-luminosity X-ray pulsars. Despite a small error box, no optical counterpart has been identified, with a limiting magnitude of $m_{V}\\sim20$ which excludes the presence of a massive companion (Mereghetti et al.\\ 1992). A recent observation with the {\\it RossiXTE}\\/ satellite provides a strong upper limit to the projected X-ray semi-major axis of 0.06~lt-s for orbital periods between 200~s and $\\sim$1~day (Mereghetti et al.\\ 1997). The lack of an optical counterpart and orbital Doppler shifts argue against a binary model for \\src, unless the companion has a low mass. Mereghetti et al.\\ (1997) find that a probable upper limit to the mass of a Roche lobe filling main sequence companion is $\\sim$0.3 \\Msun. Masses up to $\\sim$0.8 \\Msun\\ are allowed in the case of a helium-burning companion filling its Roche lobe.  The properties described above are similar to those of a small number of other X-ray pulsars with spin periods in the 5--11~s range (Mereghetti \\& Stella 1995), such as  1E\\,2259$+$586 and 4U\\,0142+614. These form a class of so-called ``anomalous'' pulsars, with clearly different properties from the majority of systems. Although accretion from a very low mass companion cannot be excluded, the lack of evidence for a binary nature from any of these systems has stimulated models where the X-ray emission originates from a compact object that is not in an interacting binary system. While an isolated, massive, white dwarf powered by the loss of rotational energy, as originally proposed for 1E\\,2259$+$586 (Paczynski 1990; Usov 1994), has been ruled out by the detection of a large increase in the spin-down rate of \\src\\  (Mereghetti 1995), other single object models, such as loss of magnetic energy of a strongly magnetized neutron star (Thompson \\& Duncan 1993), or an isolated neutron star accreting from a circumstellar disk (Corbet et al. 1995; van Paradijs et al. 1995), may be applicable.      We present a study of \\src\\, based on data obtained with the \\sax\\ satellite. We focus on the X-ray spectrum at energies $<$10 keV and on the pulse period history. As with some of the other ``anomalous'' pulsars, we find evidence for the presence of a blackbody spectral component. Since this component is not observed from the majority of other accreting X-ray pulsars, this strengthens the similarity between \\src\\ and the other better studied ``anomalous'' pulsars. ",
        "conclusions": "The 0.5--10~keV spectrum of \\src\\ can be described by the sum of an absorbed power-law and blackbody models. The existence of a blackbody component in \\src\\ was first suggested by the ASCA results of Corbet \\& Mihara (1997). However, these authors are unable to clearly discriminate between this two component model, which gives a \\rchisq\\ of 0.94 for 332 dof, and an absorbed power-law which also provides an acceptable description of the ASCA data with a \\rchisq\\ of 1.02 for 337 dof. The \\sax\\ results reported here clearly require that the \\src\\ spectrum differs significantly from an absorbed power-law. This deviation in consistent with a blackbody, but we cannot exclude the possibility that it has another form. A discussion about the physical interpretation of the alternate models can be found in White et al.\\ (1996). This discussion is applicable to \\src, since its spectral shape is roughly similar to that of 4U\\,0142+614. The reason for the significant \\sax\\ detection of spectral complexity is probably related to the combination of good energy resolution and extended low energy coverage of the LECS, together with the longer \\sax\\ exposure.  1E\\,2259+586  (Corbet et al.\\ 1995) and 4U\\,0142+614 (White et al.\\ 1996) have also been successfully fit with absorbed power-law and blackbody spectral models. Since the majority of X-ray pulsar spectra are fit by absorbed power-law models in the 0.5--10~keV energy range (e.g., White et al. 1983), our results strengthen the similarity between \\src\\ and the other ``anomalous'' pulsars. Table \\ref{tab:bbprops} is a compilation of the spectral properties of these ``anomalous'' pulsars. The blackbody component in these systems has been interpreted as evidence for quasi-spherical accretion onto an isolated neutron star formed after common envelope evolution and spiral-in of a massive X-ray binary (White et al.\\ 1996). In this case, the accretion flow results from the remaining part of the massive star's envelope and may consist of two components (Ghosh et al. 1997). A low-angular momentum component gives rise to the blackbody emission from a considerable fraction of the neutron star surface, while a high-angular momentum component forms an accretion disk and is responsible for the power-law emission and the long term spin-down evolution.  The area for the blackbody emitting surface obtained for \\src\\ ($\\sim$0.59 d$_{\\rm 3kpc}$ km) is smaller than those of the other ``anomalous'' systems. However, the distance to \\src\\ is poorly constrained and the assumed distance of 3~kpc could be considered as a lower limit since the  measured ${\\rm N_H}$ implies that it lies behind the Carina Nebula at 2.8~kpc (Seward et al.\\ (1986). White et al.\\ (1996) propose that the low pulsed fraction observed from 4U\\,0142+614 results from the large polar cap area in this system. This is consistent with the small polar cap area and high pulsed fraction reported here for \\src\\, but not with 1E\\,2259+586 and 1RXS{\\thinspace}J170849.0$-$400910, which both have large radii and a moderate pulsed fraction (see Table \\ref{tab:bbprops}). Interestingly, the best-fit blackbody temperature of 0.64~keV is somewhat higher than for 1E\\,2259+586 and 4U\\,0142+614 which may be related to the smaller area of the polar caps. Summarizing, we find that for \\src\\ (i) the blackbody temperature is higher, (ii) the blackbody radius is smaller, (iii) the power-law index is smaller (i.e. the spectrum is harder), and (iv) the pulsed fraction is higher than for the other ``anomalous'' pulsars (with the exception of RX{\\thinspace}J0720.4$-$3125 where the blackbody parameters are obtained from a one component fit).  The phase-resolved spectra are not well fit with a constant contribution from the blackbody component. This is unsurprising given the large pulse amplitude ($\\sim$70\\%), together with the large blackbody contribution (55\\%; 2--10~keV) to the phase averaged flux. The probable constancy of the pulsed fraction over the 1--10~keV energy range and the lack of any spectral dependence on pulse phase (see Sect.\\ \\ref{subsec:src_pulsetiming}) imply that the {\\it whole} spectrum is varying in a similar manner; i.e. the pulsed component cannot be attributed solely to either the blackbody or the power-law components. Furthermore, the long term variations in source flux (Fig.\\ \\ref{fig:fluxhistory}) and the approximate constancy of the pulsed fraction (Table~\\ref{tab:pdata}) during this time again implies that either the luminosities of the two components are correlated, or that the underlying spectral shape is a more complex ``single component'' that happens to mimic a power-law and a blackbody.  The spin-down rate of \\src\\ obtained from ROSAT and ASCA observations (Mereghetti 1995; Corbet \\& Mihara 1997), showed an 80\\% increase with respect to the value of $\\sim5 \\times 10^{-4}$ s yr$^{-1}$ measured before 1988 with \\ginga\\ and EXOSAT.  This increasing trend was further extended by the {\\it RossiXTE}\\/ observation of 1996 July ($P$ = $6.449769 \\pm 0.000004$ s, Mereghetti et al.\\ 1997), yielding a $\\dot{P}$ of 13$\\times10^{-4}$ s yr$^{-1}$ after the ASCA measurement. The pulse period obtained with \\sax\\ lies significantly below the linear extrapolation from the previous two measurements, indicating, for the first time in \\src, a decrease in the overall spin-down rate.  Table \\ref{tab:pdata} and Fig.\\ \\ref{fig:fluxhistory} indicate that there is no clear correlation between the observed long term spin-down rate ($\\dot{P}$) and the 2--10 keV source flux. At first sight this argues against an accretion hypothesis, but since $\\dot{P}$ is a long-term average and the flux an instantaneous measurement this comparison may not be valid. A much better comparison would be between the average flux and $\\dot{P}$ during the same time interval. Unfortunately, due to the faintness of the source, no such comparison is available. In view of the uncertainties in the average flux a good measure of the time variability of the source on timescales of months to years would be useful."
    },
    "9803/astro-ph9803082_arXiv.txt": {
        "abstract": "The EROS and MACHO collaborations have each published upper limits on the amount of planetary mass dark matter in the Galactic Halo obtained from gravitational microlensing searches. In this paper the two limits are combined to give a much stronger constraint on the abundance of low mass MACHOs. Specifically, objects with masses $10^{-7}~\\msun \\simlt m \\simlt 10^{-3}~\\msun$ make up less than $25$\\% of the halo dark matter for most models considered, and less than $10$\\% of a standard spherical halo is made of MACHOs in the $3.5\\ee{-7}~\\msun < m < 4.5\\ee{-5}~\\msun$ mass range. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803177_arXiv.txt": {
        "abstract": "We present a simple method for determining the (correlated) uncertainties of the light element abundances expected from big bang nucleosynthesis, which avoids the need for lengthy Monte Carlo simulations. Our approach helps to clarify the role of the different nuclear reactions contributing to a particular elemental abundance and makes it easy to implement energy-independent changes in the measured reaction rates. As an application, we demonstrate how this method simplifies the statistical estimation of the nucleon-to-photon ratio through comparison of the standard BBN predictions with the observationally inferred abundances. ",
        "introduction": "Big bang nucleosynthesis is entering the precision era \\cite{Sc97}. On the one hand, there has been major progress in the observational determination of the abundances of the light elements D \\cite{Bu97,We97}, $^3$He \\cite{Ba94,Gl96}, $^4$He \\cite{Ol97a,Iz97}, and $^7$Li\\cite{Ry96,Bo97}, although the increasing precision has highlighted discrepancies between different measurements (see Refs.\\cite{Mo97,Ho97,Le97} for recent assessments). Secondly, we have a sound analytical understanding of the physical processes involved \\cite{Be89,Es91} and the standard BBN computer code \\cite{Wa73,Ka92} which incorporates this physics is robust and can be easily altered to accomodate changes in the input parameters, e.g.\\ nuclear reaction rates \\cite{Sm93}. The comparison of increasingly accurate observationally inferred and theoretical abundances will further constrain the values of fundamental parameters, such as the nucleon density parameter (see, e.g., Ref.\\cite{Ol97b}) or extra degrees of freedom related to possible new physics beyond the Standard Model (see, e.g., Ref.\\cite{Sa96}). It goes without saying that error evaluation represents an essential part of such comparisons.  Because of the complex interplay between different nuclear reactions, it is not straightforward to assess the effect on a particular elemental yield of the uncertainties in the experimentally determined reaction rates. The authors of Ref.\\cite{Kr90} first employed Monte Carlo methods to sample the error distributions of the relevant reaction cross-sections which were then used as inputs to the standard BBN computer code. This enables well-defined confidence levels to be attached to the theoretically predicted abundances, e.g. the abundance range within which say 95\\% of the computed values fall correspond to 95\\% C.L. limits on the expected abundance. It was later realized that error correlations are also relevant, and can be estimated with the same technique \\cite{Ke94,KeKr}. The Monte Carlo (MC) approach has since become the standard tool for comparing theory and data \\cite{Sm93,Kr94,Co95,Ha95,Ol97c}. However, although it can include refinements such as asymmetric or temperature-dependent reaction rate uncertainties \\cite{Sm93}, it requires lengthy calculations which need to be repeated each time (any of) the input parameters are changed or updated. Since we may expect continued improvement in the determination of the relevant parameters, it is desirable to have a faster method for error evaluation and comparison with observations.  In this work we propose a simple method for estimation of the BBN abundance uncertainties and their correlations which requires little computational effort. The method, based on linear error propagation, is described in Sec.~II. A concrete application is given in Sec.~III, where theory and observations are compared using simple $\\chi^2$ statistics to obtain the best-fit value of the nucleon-to-photon ratio. In Sec.~IV we study with this method the relative importance of different nuclear reactions in determining the synthesized abundances. Conclusions and perspectives for further work are presented in Sec.~V. ",
        "conclusions": "We have shown that a simple method based on linear error propagation allows us to quantify the uncertainties associated with the elemental abundances expected from big bang nucleosynthesis, in excellent agreement with the results obtained from Monte Carlo simulations. This method makes transparent which nuclear reaction rate is mainly responsible for the uncertainty in the abundance of a given element. If determinations of the primordial abundances improve to the point where the observational errors become smaller than the theoretical uncertainties (say for $^7$Li), this will enable attention to be focussed on the particular reaction rate whose value needs to be experimentally better known.  We have also demonstrated that for standard BBN, our method enables the use of simple $\\chi^2$ statistics to obtain the best-fit value of $\\eta$ from the comparison of theory and observations. At present there are conflicting claims regarding the primordial abundances of, particularly, D and $^4$He, and different choices of input data sets imply values of $\\eta$ differing by a factor of $\\sim\\,3$. However this quantity can also be determined through measurements of the angular anisotropy of the cosmic microwave background (CMB) on small angular scales. Within a decade the forthcoming all-sky surveyors MAP and PLANCK are expected to pinpoint the nucleon density to within $\\sim5\\%$ \\cite{cmb}. Such measurements probe the acoustic oscillations of the coupled photon-matter plasma at the (re)combination epoch and will thus provide an independent check of BBN, assuming $\\eta$ did not change significantly between the two epochs.% \\footnote{New physics beyond the Standard Model can change $\\eta$, e.g. by increasing the photon number through massive particle decay \\cite{decay} or, more exotically, by {\\em decreasing} the photon number through photon mixing with a shadow sector \\cite{Ba91}. However such possibilities are strongly constrained by the absence of distortions in the Planck spectrum of the CMB \\cite{El92} and also, in the latter case, by the absence of Sakharov oscillations in the power spectrum of large-scale structure \\cite{Bi97}.} Nevertheless precise measurements of light element abundances, particularly $^4$He, are still crucial because they provide a unique probe of physical conditions, in particular the expansion rate at the BBN epoch. To illustrate, if $\\eta$ was determined by the CMB measurements to be $\\approx\\,2\\times10^{-10}$ (consistent with data set ``A''), but the abundance of $^4$He was established to be actually closer to its higher value of $\\approx\\,24\\%$ in data set ``B'', this would be a strong indication that the expansion rate during BBN was higher than in the standard case with $N_\\nu=3$ neutrinos. Although the number of $SU(2)$ doublet neutrinos is indeed 3, there are many light particles expected in extensions of the Standard Model, e.g. singlet neutrinos, which can speed up the expansion rate during nucleosynthesis \\cite{Sa96}. The generalization of our method to such non-standard cases is straightforward and we intend to present these results in a future publication \\cite{us}. It is clear that BBN analyses will continue to be important in this regard for both particle physics and cosmology."
    },
    "9803/astro-ph9803207_arXiv.txt": {
        "abstract": "Hubble Space Telescope observations of the gravitational lens PG~1115+080 in the infrared show the known $z_l=0.310$ lens galaxy and reveal the $z_s=1.722$ quasar host galaxy. The main lens galaxy G is a nearly circular (ellipticity $\\epsilon < 0.07$) elliptical galaxy with a de Vaucouleurs profile and an effective radius of $R_e = 0\\farcs59\\pm0\\farcs06$ ($1.7 \\pm 0.2 h^{-1}$ kpc for $\\Omega_0=1$ and $h = H_0/100$ \\kmm). G is part of a group of galaxies that is a required component of all successful lens models. The new quasar and lens positions (3 milliarcsecond errors) yield constraints for these models that are statistically degenerate, but several conclusions are firmly established. (1) The principal lens galaxy is an elliptical galaxy with normal structural properties, lying close to the fundamental plane for its redshift. (2) The potential of the main lens galaxy is nearly round, even when not constrained by the small ellipticity of the light of this galaxy. (3) All models involving two mass distributions place the group component near the luminosity-weighted centroid of the brightest nearby group members. (4) All models predict a time delay ratio $r_{ABC}\\simeq 1.3$. (5) Our lens models predict $H_0=44\\pm4$ \\kmm\\ if the lens galaxy contains dark matter and has a flat rotation curve, and $H_0=65\\pm5$ \\kmm\\ if it has a constant mass-to-light ratio. (6) Any dark halo of the main lens galaxy must be truncated near $1\\farcs5$ ($4 h^{-1}$ kpc) before the inferred \\Ho\\ rises above $\\sim 60$ \\kmm. (7) The quasar host galaxy is lensed into an Einstein ring connecting the four quasar images, whose shape is reproduced by the models. Improved NICMOS imaging of the ring could be used to break the degeneracy of the lens models. ",
        "introduction": "Gravitational lens time delays offer a means of determining the Hubble constant that is purely geometrical and hence completely avoids the complications of the local distance scale (Refsdal 1964). The time delay for Q~0957+561 is now well-measured (Schild \\& Thomson 1997; Kundi\\'c et al. 1997; Haarsma et al. 1997), but significant systematic uncertainties remain due to the degeneracy between the mass of the primary lens galaxy and its host cluster (e.g. Grogin \\& Narayan 1996; Bernstein et al. 1997; Romanowsky \\& Kochanek 1998). No single lens is likely to be completely free of systematic uncertainties, so a reliable estimate of $H_0$ should rely on an ensemble of lenses. There are now three more systems with time delay estimates: PG~1115+080 (Schechter et al. 1997), B~1608+656 (Fassnacht et al. 1996), and B~0218+357 (Corbett et al. 1996), which need detailed exploration of their lens models to examine the systematic uncertainties.  PG~1115+080 was the second gravitationally lensed quasar to be discovered (Weymann et al. 1980). The source is an optically selected, radio-quiet quasar at redshift $z_s=1.722$. Hege et al. (1981) first resolved the four quasar images (a close pair A1/A2, B and C), confirming the early model of Young et al. (1981) that the lens was a five-image system, one image being hidden in the core of the lens galaxy. Henry \\& Heasly (1986) detected the lens galaxy, followed by gradual improvements in the astrometry by Kristian et al. (1993; hereafter K93), and Courbin et al. (1997). The redshift of the lens galaxy was determined by Angonin-Willaime, Hammer \\& Rigaut (1993) and confirmed by Kundi\\'c et al. (1997) and Tonry (1998) to be $z_l=0.310$. Tonry also determined the central velocity dispersion of the lens galaxy: $\\sigma = 281\\pm25$ km s$^{-1}$. The spatial resolution of published data has always been insufficient to perform any surface photometry on the lens galaxy. Young et al. (1981) noted that the lens seemed to be part of a small group centered to the southwest of the lens, with a velocity dispersion of approximately $270\\pm70$ km s$^{-1}$ based on only four galaxy redshifts (Kundi\\'c et al. 1997). The group is an essential component of any model that successfully fits the lens constraints (Keeton, Kochanek \\& Seljak 1997; Schechter et al. 1997). Finally, Schechter et al. (1997) successfully determined two time delays between the images, which were reanalyzed by Barkana (1997) to give $\\Delta\\tau_{BC}=25.0^{+1.5}_{-1.7}$ days and the time delay ratio $r_{ABC}=\\Delta\\tau_{AC}/\\Delta\\tau_{BA}=1.13\\pm0.18$. These results were analyzed by Keeton \\& Kochanek (1997) and Courbin et al. (1997) to deduce $H_0=53_{-7}^{+15}$ km s$^{-1}$ Mpc$^{-1}$, with comparable contributions to the uncertainties from the time delay measurement and the models. The extreme variations are given in non-parametric form by Saha \\& Williams (1997), although some of these models may not be physical.  We present new near-infrared observations of the PG~1115+080 system obtained with the Hubble Space Telescope (HST) NICMOS camera. These are the first results of the CfA-Arizona Space Telescope Lens Survey (CASTLES).\\footnote{A summary of gravitational lens data and model results, including CASTLES data, is available at the URL http://cfa-www.harvard.edu/castles.} After summarizing the observations in \\S2, we present improved astrometry in \\S2.1, the first surface photometry of the lens galaxy in \\S2.2, a discussion of lens models and the Hubble constant in \\S2.3, photometry of the nearby group in \\S2.4, and comments on the quasar host galaxy in \\S2.5. In \\S3 we comment on the strengths and limitations of this system as a cosmological tool. ",
        "conclusions": "New infrared data on PG~1115+080 affirms multiple-component gravitational lens systems as powerful cosmological tools. The major puzzle remaining in the PG~1115+080 system is the anomalous A1/A2 flux ratio. Our observations rule out differential extinction as an explanation, and microlensing is ruled out by its lack of variability. Since a flux ratio near 0.9 is a generic feature of the large scale potential near a fold caustic, only a potential perturbation intermediate between that produced by isolated stars (microlensing) and by the overall galaxy can explain the flux ratio. The potential of PG~1115+080 must be perturbed either by a satellite galaxy or a globular cluster. Mao \\& Schneider (1998) showed that such perturbations alter the time delay -- and so the inferred value of the Hubble constant -- fortunately by no more than 2-3\\%.  Our improved astrometry greatly reduces some of the degeneracies in early models of the system. The group position is now well constrained and located near the luminosity centroid of the four bright group galaxies, and the lens galaxy is constrained to be nearly circular. Unfortunately, the degeneracies in the $H_0$ estimate have been exacerbated because with the revised astrometry the models no longer favor dark matter models over constant $M/L$ models for the main lens galaxy. Because all 4 images are located at nearly the same radial distance from the center of the lens galaxy, we do not expect the models to be sensitive to the radial mass profile of the lens galaxy (Kochanek 1991; Wambsganss \\& Pa\\'czynski 1994). We find \\Ho\\ ranging from $44\\pm4$ \\kmm\\ if the lens galaxy is modeled as a singular isothermal ellipsoid and the group as a singular isothermal sphere, to $65\\pm5$ ($72\\pm5$) \\kmm\\ for $\\Omega_0 = 1$ ($0.1$) if the lens galaxy has a constant $M/L$. Note that we find evidence for dark matter in our high value of $M/L$. A model with an adjustable truncation radius shows that the halo must be truncated on scales comparable to the ring diameter for $H_0$ to exceed $60$ \\kmm. Such a halo seems smaller than physically plausible given that the velocity dispersions of the group and the lens galaxy are comparable.  Further progress in reducing the uncertainties depends on improving the time delay measurements and on making more detailed studies of the Einstein ring formed by the quasar host galaxy. First, for any given mass profile, most of the current errors in \\Ho\\ are due to the uncertainties in the time delays. Second, all the best fit models predict time delay ratios near $r_{ABC}=1.3$, consistent with the current measurement of $1.13 \\pm 0.18$. If nothing else, a more accurate measurement of the delay ratio than is now available would be a powerful test of the models. For any given lens mass profile, the ratio constraint would further reduce the parameter space for the position and mass of the group, or could be used to constrain more complicated models (e.g., Saha \\& Williams 1997). Deep new observations to determine the surface brightness of the ring accurately, combined with direct measurement of the point spread function at the time of the observations would probably permit us directly to break the degeneracy of the models. Only NICMOS, however, has the ability to make these difficult observations within a decade.  No single technique or observation can tie down the Hubble constant -- the long history of unrecognized or underestimated systematic errors in this subject encourages humility. Nevertheless, PG~1115+080 demonstrates the potential of the gravitational lens approach. With recent reductions in age estimates of the oldest globular clusters (Chaboyer et al. 1998), and the likelihood that the mass density of the universe is lower than the Einstein-de Sitter case (e.g., Garnavich et al. 1998), the possibility of an age conflict in the big bang model has receded. The modeling of PG~1115+080 gives a plausible upper bound on the Hubble constant if we accept that the group is not a point mass and that the lens galaxy is unlikely to have a mass distribution that is more concentrated than its light distribution. This bound is $H_0 < 67$ ($72$) \\kmm\\ for $\\Omega_0 = 1$ ($0.1$). The most recent result of the HST Extragalactic Distance Scale Key Project is $H_0 = 72 \\pm5$ (random) $\\pm12$ (systematic) km s$^{-1}$ Mpc$^{-1}$ (Madore et al. 1998). Our upper limit is inconsistent with the upper end of the range from the Key Project, although it is consistent with the lower end of the range. There appears to be satisfactory concordance among the basic parameters of the big bang model, and between direct and indirect measures of the distance scale. Gravitational lenses can be expected to play an increasing role as versatile cosmological tools.  \\vskip 1truecm"
    },
    "9803/astro-ph9803031_arXiv.txt": {
        "abstract": "In this paper, we measure the ellipticities of 30 LSB dI galaxies and compare the ellipticity distribution with that of 80 dEs (Ryden \\& Terndrup 1994; Ryden et al.\\ 1998)\\markcite{ryden 1994; 1998} and 62 BCDs (Sung et al.\\ 1998).\\markcite{sung1998} We find that the ellipticity distribution of LSB dIs is very similar to that of BCDs, and marginally different from that of dEs.  We then determine the distribution of intrinsic shapes of dI galaxies and compare to those of other type dwarf galaxies under various assumptions.  First, we assume that LSB dIs are either all oblate or all prolate, and use non-parametric analysis to find the best-fitting distribution of intrinsic shapes. With this assumption, we find that the scarcity of nearly circular LSB dIs implies, at the 99\\% confidence level, that they cannot be a population of randomly oriented oblate or prolate objects.  Next, we assume that dIs are triaxial, and use parametric analysis to find permissible distributions of intrinsic shapes.  We find that if the intrinsic axis ratios, $\\beta$ and $\\gamma$, are distributed according to a Gaussian with means $\\beta_0$ and $\\gamma_0$ and a common standard deviation of $\\sigma$, the best-fitting set of parameters for LSB dIs is $(\\beta_0,\\gamma_0,\\sigma) = (0.66,0.50,0.15)$, and the best fit for BCDs is $(\\beta_0,\\gamma_0,\\sigma) = (0.66,0.55,0.16)$, while the best fit for dEs is $(\\beta_0,\\gamma_0,\\sigma) = (0.78,0.69,0.24)$.  The dIs and BCDs thus have a very similar shape distribution, given this triaxial hypothesis, while the dEs peak at a somewhat more spherical shape. Our results are consistent with an evolutionary scenario in which the three types of dwarf galaxy have a close relation with each other. ",
        "introduction": "Although the faint end of the galaxy luminosity function is not well determined, recent studies indicate that low-surface-brightness (LSB) dwarf galaxies are by far the most numerous type of galaxy, and contribute a significant fraction of the mass of the universe (Reaves 1983; Binggeli, Sandage, \\& Tammann 1985; Phillipps et al.\\ 1987)\\markcite{reaves1983, sandage1985, phillipps1987}.  Morphologically, dwarf galaxies, like their counterpart bright galaxies, are classified into several types.  The most common type of dwarf galaxy ($\\sim 80\\%$ of the total) is the dwarf elliptical (dE). These galaxies have regular elliptical isophotes and roughly exponential surface brightness profiles; they are often found in groups and clusters (Davies et al.\\ 1988).\\markcite{davies1988} The second type of dwarf galaxy is the blue compact dwarf (BCD) galaxy.  In contrast to gas-poor dEs, BCDs contain giant HII regions surrounding O and B stars within a massive HI reservoir; BCDs exhibit spectra slowly rising toward the blue, implying that they are undergoing intense star formation (du Puy 1970; Searle \\& Sargent 1972).\\markcite{dupuy1970, searle1972} Most BCDs have regular isophotes in the outer region, like dEs, but the inner isophotes are frequently distorted from ellipses, due to the presence of bright HII regions (Loose \\& Thuan 1986).\\markcite{loose1986}  The final type of dwarf galaxy is the LSB dwarf galaxy. LSB dwarfs include both irregular (dI) and more regular spiral (dS) galaxies. Like BCDs, they contain a large amount of HI, often with small OB associations, and have blue colors ($B-V\\sim 0.5\\ {\\rm mag}$), indicating a significant level of recent star formation (Staveley-Smith, Davies, \\& Kinman 1992)\\markcite{staveley1992}. However, they are distinguished from BCD galaxies by their amorphous shapes even in the outer region. Additionally, in contrast to dEs, these galaxies are more likely to be found outside of clusters (Bingelli, Tarenghi, \\& Sandage 1990).\\markcite{bingelli1990}  The evolutionary connections among the three different types of dwarf galaxies remain both elusive and confusing.  There are two major competing hypotheses for the evolutionary connection between BCDs and dEs.  The first hypothesis claims that BCDs are basically a different population from dEs, as evidenced by the spectroscopic and spectrophotometric differences.  According to this scenario, BCDs are truly young systems, in which the present star burst is the first in the galaxy's lifetime.  The second hypothesis suggests that BCDs, like dEs, are mainly composed of old stellar populations, and that their observed spectroscopic features and spectral energy distributions are the result of a recent burst of star formation (Staveley-Smith et al.\\ 1992).\\markcite{stavely1992} As an evidence for the second scenario, it is argued that the near-infrared emission in the vast majority of BCDs is attributable to old K and M giants, which are the major component of dEs (Thuan 1983; Hunter \\& Gallagher 1985).\\markcite{thuan1983, hunter1985}  Similarly, there exist two competing hypotheses to explain the evolutionary connection between dIs and dEs.  The first hypothesis states that dE galaxies are the faded remnants of previously actively star-forming dI galaxies whose gas has been lost. There exists circumstantial evidence that dEs have in fact evolved directly from dIs.  Faber \\& Lin (1983)\\markcite{faber1983} and Kormendy (1985)\\markcite{kormendy1985} have used the similarity in the surface brightness profiles of dIs and dEs, which are mostly exponential, to argue that gas-rich dIs are the progenitors of dEs.  The second hypothesis for the relation between dIs and dEs states that they represent parallel sequences of dwarf galaxies, fundamentally separated by the intrinsic difference in their structure.  The observational evidence for this hypothesis is based mostly on the differences in appearance between the two types of dwarf galaxies; for instance, dIs have a more diffuse light distribution than dEs, and lack the bright nucleation which is frequently found in dEs.  In addition to these differences, there is a dissimilarity in the flattening distribution of dEs and dIs; the apparent flattening of a galaxy is customarily given either by the apparent axis ratio $q$ or by the ellipticity $\\epsilon \\equiv 1 - q$. Bothun et al.\\ (1986) and Impey \\& Bothun (1997)\\markcite{bothun1986, impey1997} presented the results of Ichkawa, Wakamatsu, \\& Okamura (1986)\\markcite{ ichikawa1986} and Caldwell (1983)\\markcite{caldwell1983} as evidence for the different flattening distributions between dEs and dIs. However, the analysis of Ichikawa et al.\\ was based on the comparison between the flattening distributions of dEs and bright (non-dwarf) spiral galaxies; the situation is similar for Caldwell's analysis.  In addition, contrary to the claims of Bothun et al.\\ and Impey \\& Bothun, both Ichikawa et al.\\ and Caldwell showed that the flattening distribution of dEs is similar to that of bright irregular galaxies.   There have been previous attempts to compare the apparent axis ratio distributions between LSB dI galaxies and other types of dwarf galaxies.  For example, Staveley-Smith et al.\\ (1992)\\markcite{stavely1992} constructed the axis ratio distribution for 438 Uppsala Galaxy Catalogue (hereafter UGC, Nilson 1973)\\markcite{nilson1973} LSB galaxies, and compared it to that of BCDs whose ellipticities were measured from Palomar Observatory Sky Survey (POSS) plates by Gorden \\& Gottesman (1981).\\markcite{gorden1981} However, previous studies of axis ratio distributions suffer from large uncertainties for several reasons.  First, owing to the small dimensions and low surface brightness of dwarf galaxies, estimating their axis ratio is difficult and leads to large uncertainties. Second, the UGC sample used by Staveley-Smith et al.\\ (1992)\\markcite{stavely1992} is known to be inhomogeneous, containing galaxies ranging from true dwarf galaxies to more luminous very low surface brightness systems (Thuan \\& Seitzer 1979; McGaugh, Schombert, \\& Bothun 1995).\\markcite{thuan1979, mcgaugh1995} Finally, previous determinations of LSB dI axis ratios have been based on photographic plates; for comparison with recent CCD observations of other types of dwarf galaxy, it is essential to have measurements of the axis ratios of a homogeneous sample of LSB dIs based on modern CCD observations.  In this paper, we measure the ellipticities of 30 LSB dI galaxies and compare the ellipticity distribution with that of 80 dEs (Ryden \\& Terndrup 1994; Ryden et al.\\ 1998)\\markcite{ryden 1994; 1998} and 62 BCDs (Sung et al.\\ 1998).\\markcite{sung1998} We find that the ellipticity distribution of LSB dIs is very similar to that of BCDs, and marginally different from that of dEs. We then determine, under various assumptions, the distribution of intrinsic shapes of dI galaxies and compare it to that of other types of dwarfs. First, we assume that LSB dIs are either all oblate or all prolate, and use non-parametric analysis to find the best-fitting distribution of intrinsic shapes. With this assumption, we find that the scarcity of nearly circular LSB dIs implies, at the 99\\% confidence level, that they cannot be a population of randomly oriented oblate or prolate objects.  Next, we assume that dIs are triaxial, and use parametric analysis to find permissible distributions of intrinsic shapes.  We find that if the intrinsic axis ratios, $\\beta$ and $\\gamma$, are distributed according to a Gaussian with means $\\beta_0$ and $\\gamma_0$ and a common standard deviation of $\\sigma$, the best-fitting set of parameters for LSB dIs is $(\\beta_0,\\gamma_0,\\sigma) = (0.66,0.50,0.15)$, and the best fit for BCDs is $(\\beta_0,\\gamma_0,\\sigma) = (0.66,0.55,0.16)$, while the best fit for dEs is $(\\beta_0,\\gamma_0,\\sigma) = (0.78,0.69,0.24)$. The LSB dIs and BCDs thus have a very similar shape distribution, given this triaxial hypothesis, while the dEs peak at a somewhat more spherical shape. Our results are consistent with an evolutionary scenario in which the three types of dwarf galaxy have a close relation with each other. ",
        "conclusions": "We measure the ellipticities for a sample of 30 LSB dIs and compare the distribution of ellipticities with those for the samples of 62 BCDs and 80 dEs.  From this comparison, we find that the axis ratio distribution of LSB dIs is very similar to that of BCDs.  Compared to dEs, LSB dIs are slightly flatter, but the difference is marginal.  We also determine the intrinsic shape of LSB dIs from the distribution of apparent axis ratios.  From the non-parametric analysis, we find the hypothesis that our sample LSB dIs are randomly oriented oblate or prolate objects is rejected with strong confidence level.  On the other hand, the shape of LBS dI galaxies are well described by triaxial spheroids if their axis ratios, $\\beta$ and $\\gamma$, have a Gaussian distribution.  From the parametric analysis, we determine the best-fitting parameters are $(\\beta_0,\\gamma_0,\\sigma) = (0.66, 0.50, 0.15)$.  These results directly contradict the long-standing belief that LSB dIs have very flattened disky shapes, quite different from the spheroidal shapes of dEs and BCDs.  Therefore, our results are consistent with the scenario that the three major types of dwarf galaxies have very close evolutionary connections."
    },
    "9803/astro-ph9803025_arXiv.txt": {
        "abstract": " ",
        "introduction": "The 48~ms radio pulsar PSR~B1259$-$63 is in a highly eccentric 3.4~yr orbit with a Be star (Johnston et al. 1992).  Near periastron, for a sufficiently strong Be star wind, PSR~B1259$-$63 could possibly make a transition to an accretion regime, exhibiting X-ray pulsations of luminosity $>$10$^{35}$~erg~s$^{-1}$.  {\\it ASCA} X-ray observations at three epochs around the 1994 periastron passage (Kaspi et al. 1994; Hirayama 1996) found the source to be unpulsed, with modest luminosity $\\sim$10$^{34}$~erg~s$^{-1}$, for a distance of 2~kpc.  These properties imply that shock emission produces the observed X-rays, rather than accretion (Tavani \\& Arons 1997). However radio timing observations suggested sudden, brief accretion events near periastron that could not be ruled out by the X-ray observations (Manchester et al. 1995).  Indeed the pulsar's anomalously low period suggests it may have undergone occasional accretion episodes in the past.  {\\it ASCA} observations during the 1997 periastron passage were impossible due to solar constraints.  The ASM on {\\it RXTE} has been observing bright celestial X-ray sources regularly since early 1996. The instrument consists of three ``scanning shadow cameras,'' each containing a position-sensitive proportional counter that is mounted below a wide-field collimator, covered by a coded mask.  The instrument provides roughly 5 celestial scans per day, with diminished exposure in directions toward the Sun.  Further information on the ASM is given by Levine et al. (1996).  While PSR~B1259$-$63 is not routinely monitored for the ASM source-history database, we extracted an X-ray light curve as an archival analysis project using standard techniques. The derived light curve (Figure~1a) covers the time interval of MJD 50178--50762 (1996 Apr 5 to 1997 Nov 10).  The overall mean intensity of PSR B1259$-$63 during this time interval was $(0.5 \\pm 2.4)\\times 10^{-11}$~erg~s$^{-1}$~cm$^{-2}$ at 2--10 keV. This result is given with consideration of both the systematic bias and uncertainty in the ASM light curves for faint, yet fairly isolated X-ray sources. Figure~1b shows $3\\sigma$ upper limits obtained in the periastron vicinity in 2-day averages.  It shows that in the 2 weeks following periastron, the source was not observed to be brighter than $\\sim$10 times the 1994 periastron luminosity.  The $3\\sigma$ upper limits to the X-ray emission from PSR~B1259$-$63 from the ASM thus argue that the pulsar did not undergo even brief episodes of accretion during the 1997 periastron passage.  This is consistent with the shock emission model, as well as with new conclusions based on timing that rule out the sudden ``spin-ups'' that were claimed previously (Wex et al. 1998).  \\begin{figure}[t] \\centerline{ \\psfig{figure=psr1259.ps,height=5.5cm} \\psfig{figure=psr1259_upper.ps,height=5.5cm} } \\caption{(a) ASM light curve for PSR~B1259$-$63. (b) 3$\\sigma$ upper limits to the ASM flux near periastron from (a). The horizontal dashed line represents the brightest {\\it ASCA} 1994 periastron detection (Kaspi et al. 1994), and the vertical dotted line indicates periastron.} \\end{figure} ",
        "conclusions": "The 48~ms radio pulsar PSR~B1259$-$63 is in a highly eccentric 3.4~yr orbit with a Be star (Johnston et al. 1992).  Near periastron, for a sufficiently strong Be star wind, PSR~B1259$-$63 could possibly make a transition to an accretion regime, exhibiting X-ray pulsations of luminosity $>$10$^{35}$~erg~s$^{-1}$.  {\\it ASCA} X-ray observations at three epochs around the 1994 periastron passage (Kaspi et al. 1994; Hirayama 1996) found the source to be unpulsed, with modest luminosity $\\sim$10$^{34}$~erg~s$^{-1}$, for a distance of 2~kpc.  These properties imply that shock emission produces the observed X-rays, rather than accretion (Tavani \\& Arons 1997). However radio timing observations suggested sudden, brief accretion events near periastron that could not be ruled out by the X-ray observations (Manchester et al. 1995).  Indeed the pulsar's anomalously low period suggests it may have undergone occasional accretion episodes in the past.  {\\it ASCA} observations during the 1997 periastron passage were impossible due to solar constraints.  The ASM on {\\it RXTE} has been observing bright celestial X-ray sources regularly since early 1996. The instrument consists of three ``scanning shadow cameras,'' each containing a position-sensitive proportional counter that is mounted below a wide-field collimator, covered by a coded mask.  The instrument provides roughly 5 celestial scans per day, with diminished exposure in directions toward the Sun.  Further information on the ASM is given by Levine et al. (1996).  While PSR~B1259$-$63 is not routinely monitored for the ASM source-history database, we extracted an X-ray light curve as an archival analysis project using standard techniques. The derived light curve (Figure~1a) covers the time interval of MJD 50178--50762 (1996 Apr 5 to 1997 Nov 10).  The overall mean intensity of PSR B1259$-$63 during this time interval was $(0.5 \\pm 2.4)\\times 10^{-11}$~erg~s$^{-1}$~cm$^{-2}$ at 2--10 keV. This result is given with consideration of both the systematic bias and uncertainty in the ASM light curves for faint, yet fairly isolated X-ray sources. Figure~1b shows $3\\sigma$ upper limits obtained in the periastron vicinity in 2-day averages.  It shows that in the 2 weeks following periastron, the source was not observed to be brighter than $\\sim$10 times the 1994 periastron luminosity.  The $3\\sigma$ upper limits to the X-ray emission from PSR~B1259$-$63 from the ASM thus argue that the pulsar did not undergo even brief episodes of accretion during the 1997 periastron passage.  This is consistent with the shock emission model, as well as with new conclusions based on timing that rule out the sudden ``spin-ups'' that were claimed previously (Wex et al. 1998).  \\begin{figure}[t] \\centerline{ \\psfig{figure=psr1259.ps,height=5.5cm} \\psfig{figure=psr1259_upper.ps,height=5.5cm} } \\caption{(a) ASM light curve for PSR~B1259$-$63. (b) 3$\\sigma$ upper limits to the ASM flux near periastron from (a). The horizontal dashed line represents the brightest {\\it ASCA} 1994 periastron detection (Kaspi et al. 1994), and the vertical dotted line indicates periastron.} \\end{figure}"
    },
    "9803/astro-ph9803163_arXiv.txt": {
        "abstract": "We describe a concept for an imaging spectrograph for a large orbiting observatory such as NASA's proposed Next Generation Space Telescope (NGST) based on an imaging Fourier transform spectrograph (IFTS).  An IFTS has several important advantages which make it an ideal instrument to pursue the scientific objectives of NGST. We review the operation of an IFTS and make a quantitative evaluation of the signal-to-noise performance of such an instrument in the context of NGST. We consider the relationship between pixel size, spectral resolution, and diameter of the beamsplitter for imaging and non-imaging Fourier transform spectrographs and give the condition required to maintain spectral modulation efficiency over the entire field of view.  We give examples of scientific programs that could be performed with this facility. ",
        "introduction": "The Association of Universities for Research in Astronomy (AURA), with NASA support, recently appointed a committee to ``study possible missions and programs for UV-Optical-IR astronomy in space for the first decades of the twenty-first century.''  The report urged the development of a general-purpose, near-infrared observatory equipped with a passively cooled primary mirror ($T \\le 70$~K) with a minimum diameter of 4 meters (Dressler 1996). To enhance its performance, the report recommended that the observatory be placed as far from the Earth-Moon system as possible to reduce stray light and to maintain the telescope's relatively low temperature. With such a facility, it should be possible to learn in detail how galaxies formed, measure the large-scale curvature of space-time by measuring distant standard-candles, trace the chemical evolution of galaxies, and study nearby stars and star-forming regions for signs of planetary systems. A detailed discussion of the next generation space telescope (NGST) and its scientific potential given by Stockman (1997).  For NGST to attain these scientific objectives, it must have an instrument which is designed to execute panchromatic observations over the critical 1--15 $\\mu$m wavelength range of the faintest detectable objects.  With nJy sensitivity levels attainable at near-infrared (1--5 $\\mu$m) (NIR) and mid-infrared (5--15$\\mu$m) (MIR) wavelengths, NGST will be able to study the well-calibrated rest-frame optical diagnostics in distant ($z=3-10$) galaxies, thus probing for the first time, their stellar content, star-formation history and nuclear activity.  At the longer wavelengths, NGST can investigate these properties in $z=3-5$ galaxies using diagnostics that are unaffected by dust extinction and reddening, and also study the dust properties directly.  At the flux limits characteristic of NGST, the confusion limit is likely to be approached, with virtually every pixel having significant information (e.g., by extrapolation from counts in the Hubble Deep Field (Williams et al. 1996)).  As a result, one of the best ways to maximize the scientific output from NGST is to provide a wide-field imaging spectrograph that is efficient in this limit.  An imaging Fourier transform spectrometer (IFTS) provides these capabilities in a low-cost, high throughput, compact design.  It provides the only efficient means of conducting {\\it unbiased} spectroscopic surveys of the high-$z$ Universe, i.e., without object preselection (e.g., using broad band colors) and without the restrictions imposed by spectrometer slit geometry and placement.  An IFTS also allows spectroscopy over a wide bandpass, affords flexibility in choice of resolution, is easy to calibrate, and is ideal for wide-field spectroscopic surveys.  Bennett et al. (1993) and Bennett, Carter, \\& Fields (1995) describe the operating principles of imaging Fourier transform spectrographs and compare their performance with alternative imaging spectrometers.  A comprehensive review of the application of interferometers and the techniques of Fourier spectroscopy to astrophysical problems is given by Ridgway \\& Brault (1984), and a recent summary of the field, including a description of an astronomical IFTS is given by Maillard (1995).  Spaceborne Fourier transform spectrometers have been responsible for spectacular results in the fields of planetary exploration and cosmology.  Infrared FT spectrometers developed at the Goddard Space Flight Center (GSFC) flew on board the Mariner 9 mission to Mars, and were carried to the outer planets by the Voyager spacecraft (Hanel et al. 1992).  The instruments provided superb data revealing, for the first time, the composition of the atmospheres of the giant gaseous planets (e.g. Jupiter, Hanel et al. 1979).  The Composite Infrared Spectrometer (CIRS), currently traveling to Saturn onboard the Cassini spacecraft, is another instrument developed at GSFC.  CIRS is the first step towards an imaging FTS as it has a linear array of detectors, rather than a single element detector, in order to map the temperature and composition of the atmospheres of Saturn and Titan as a function of altitude during limb soundings.  (Kunde et al. 1996).  The definitive measurement of the spectrum of the cosmic microwave background radiation (CMBR) was one of the most dramatic experimental measurements of this decade (Mather et al. 1990; Gush, Halpern \\& Wishnow 1990).  The FIRAS instrument onboard the NASA satellite COBE which first performed this measurement, and the COBRA rocket experiment conducted by the University of British Columbia which confirmed it a few months later, were both liquid-helium cooled, differential Fourier transform spectrometers.  These instruments used a dual-input, dual-output configuration where one input viewed the sky and the other viewed a blackbody calibrator (Mather, Fixen \\& Shafer 1993; Gush \\& Halpern 1992).  Absolute photometric measurements were obtained by reference to the blackbody calibrator, and the CMBR was observed to have an undistorted Planck spectrum corresponding to a temperature of $2.728\\pm0.004$ K (Fixen et al. 1996).  The IFTS proposed here can be thought of as an extension of these experiments where focal plane detector arrays yield simultaneous imaging and spectral information.  In the next decade, missions such as WIRE, AXAF, and SIRTF will expand astrophysical horizons, possibly unveiling entirely new populations of objects.  An IFTS offers the flexibility (e.g., spectral resolution) that may prove essential in investigating the nature of these sources. Due to its flexibility and its ability to provide simultaneous imaging and spectroscopy of every object in the field of view (FOV), an IFTS is a {\\it necessary} instrument for the NGST mission.  \\section {IFTS CONCEPT} \\label{concept}  An IFTS (Fig. \\ref{fts}) is axis-symmetric, and the optical path difference (OPD) is the same for all the points of the image with the same angle of incidence from the axis of the interferometer. Hence, the FOV is circular. On the object side, an entrance collimator illuminates the interferometer with parallel light. The interfering beams are collected by the output camera, creating a stigmatic relation between the object and image planes. By placing a detector array in the output focal plane the entrance field is imaged on the array and each pixel works as a single detector matched to a point on the sky.  \\begin{figure}[thb] \\centerline{\\psfig{figure=fig1.eps,width=3.3in,angle=-90}} \\caption{\\small A sketch of the optics of a simple single-beam imaging Fourier transform spectrograph consisting of a collimating lens, a beam splitter, two mirrors (one movable), and a camera lens. The optical path difference is $x$.} \\label{fts} \\end{figure}   Retrieving spectral information involves recording the interferogram generated by the source imaged onto the focal plane array (FPA). The OPD is scanned in discrete steps since FPAs are integrating detectors. Scanning in this way generates a data cube of two-dimensional interferograms.  The signal from the same pixel in each frame forms an independent interferogram.  These interferograms are Fourier transformed individually yielding a spectral data cube composed of the same spatial elements as the image.  \\subsection{A Perfect Match to NGST Science} \\label{perfect}  The features of an IFTS which make it the instrument of choice for NGST are efficiency, flexibility, and compactness. The most compelling reason for choosing an IFTS is that in the dual port design (see Fig. \\ref{optical-layout}) virtually every photon collected by the telescope is directed towards the focal plane for detection. Other solutions are inefficient, inflexible, and wasteful of mass, power, and volume. Cameras equipped with filters admit only a restricted bandpass at low spectral resolution. To compete with the spectral multiplex advantage of an IFTS, a camera system needs multiple dichroics and FPAs. The additional mass and thermal load is a severe penalty. Classical dispersive spectrographs have slit losses, grating inefficiencies due to light lost in unwanted orders, and limited free spectral range (the same is true for a Fabry-Perot).  {\\em An IFTS acquires full bandpass imaging simultaneously with higher spectral resolution data.} Therefore, a high SNR broad-band image always accompanies full spectral sampling of the FOV with no penalty in integration time.  An IFTS is a true imaging spectrograph and measures a spectrum for every pixel in the FOV.  It is not necessary to choose which regions in the image are most deserving of spectroscopic analysis.  Overheads are eliminated because no additional observing time is needed for imaging prior to object selection, and there is no delay in positioning slit masks, fibers, or image slicing micro-mirrors. Thus, an IFTS will produce a rich scientific legacy with tremendous potential for serendipity.  \\begin{figure}[thb] \\centerline{ \\psfig{figure=fig2.ps,width=3.3in,angle=-90}} \\caption{\\small Schematic optical layout of a 60$^\\circ$ dual-input, dual-output Michelson interferometer. } \\label{optical-layout} \\end{figure}   Table \\ref{capabilities} details the capabilities of a putative IFTS suitable for NGST.  We use the instrument described by this table to illustrate the potential of an IFTS.  Two points in Table \\ref{capabilities} must be stressed: 1) An IFTS is spectrally multiplexed, therefore all spectral channels are obtained simultaneously within the stated integration time. 2) The free spectral range of an IFTS is limited only by the band-pass filter and the detector response. Consequently, the usual definition of resolution, $R = \\lambda/\\delta \\lambda$, is of limited use. It is conventional to scan the OPD of an IFTS in equal steps so that the resolution is constant in wavenumber, $k$. Thus, we use $M$ to denote the number of spectral channels. For example, in the NIR with a 1-5 $\\mu$m band-pass, $M=5$ means that $\\delta k = (k_{max} - k_{min})/M = 1600$~cm$^{-1}$, and a scan yields 5 bands centered 1.1, 1.3, 1.7, 2.3, \\& 3.6 $\\mu$m.   \\begin{deluxetable}{lll} \\tablewidth{0pt} \\tablecaption{Capabilities of a Putative NGST IFTS} \\tablehead{ \\colhead{} & \\colhead{NIR Channel} &\\colhead{MIR Channel}} \\startdata Design\t\t\t& Dual-port & Dual-port \\\\ Bandpass\t\t& 1-5 $\\mu$m & 5-15 $\\mu$m \\\\ Resolution  \t\t& 1 cm$^{-1}$\t& 1 cm$^{-1}$ \\\\ FOV \t\t\t& $200''$\t& $100''$ \\\\ Pixel size \t\t& $0.''05$\t& $0.''1$ \\\\ Array format \t\t& 4k$\\times$4k & 1k$\\times$1k \\\\ Detector\t\t& InSb \t\t& HgCdTe \\\\ Throughput\t\t& $> 0.5$\t& $> 0.5$ \\\\ Sensitivity\\tablenotemark{a} \\\\ ~$M=1$\t\t& 200 pJy & 13 nJy \\\\ ~$M=5$\t\t\t& 1 nJy   & 65 nJy  \\\\ ~$M=100$\t\t& 35 nJy  & 1.3 $\\mu$Jy \\\\ \\enddata \\tablenotetext{a}{SNR = 10 for a $10^5$~s integration over the entire  spectral band for a point source. $M$ is the number of simultaneous spectral channels in the band-pass ---  see \\S \\ref{perfect}. Note that all spectral channels are obtained simultaneously. The spectrum is assumed to be flat in $F_\\nu$ and the SNR is quoted at 2 $\\mu$m for the NIR channel, and at 10 $\\mu$m for the MIR channel.} \\label{capabilities} \\end{deluxetable}   The throughput of an IFTS with ideal optics is only limited by the efficiency of the beam splitter.  In a dual-input, dual-output port design no light is wasted and the throughput approaches 100\\%.  An IFTS has no loss of light or spatial information because there is no slit, hence an IFTS is perfectly adapted to doing multi-object spectroscopy in crowded or confusion limited fields.  A IFTS uses every photon whereas traditional cameras and spectrographs throw away photons (either spectrally with a filter or spatially with a slit), so at a very fundamental level an IFTS is superior.  On blaze, a good grating is 80\\% efficient, but averaged over the free spectral range this drops to about 65\\%.  An IFTS is not optimized for single-object spectroscopy because the broad-band photon shot noise is associated with every frame in the interferogram. Hence, for a single object a slit spectrograph is $\\eta_g \\eta_s M$ times faster than an IFTS of the same resolution in background limited operation, where $\\eta_g$ is the grating efficiency averaged over the blaze function, and $\\eta_s$ is the slit loss, where typically the product $\\eta_g \\eta_s \\approx 0.3$.  This disadvantage is more than compensated for by the spatial-multiplexing capability of an IFTS. A typical deep background limited exposure of an IFTS will reach $K=29.5$, SNR=10 and will contain at least 3500 and possibly, depending on cosmology, up to 11,000, objects per field (see Fig. \\ref{number-counts}).  A grating spectrograph with a fiber feed or multi-slit capability can perhaps record spectra for only a few percent of these objects at a time, requiring hundreds of pointings to make an unbiased survey of a single field, as opposed to the single IFTS imaging-cum-spectroscopic observation.      \\begin{figure}[thb] \\centerline{\\psfig{figure=fig3.ps,width=3.4in,angle=90}} \\caption{\\small $K$-band number counts (Djorgovski et al. 1995; Gardner et al. 1993, 1996; Glazebrook et al. 1994; Huang et al. 1997; Mobasher et al. 1986; Moustakis et al. 1997; McLeod et al. 1995; \\& Metcalfe et al. 1996) together with models of the luminosity function modeled using the formalism of Gardner (1998), which has been used to extrapolate the number counts into the NGST domain. The solid lines include the effects of passive evolution, while the dashed lines include only K-corrections. The upper line in each case is for $q_0 = 0.1$, and the lower lines are for $q_0 = 0.5$.  Current number counts imply at least 3500 objects per $3.'3$ NGST field, while the extrapolations shown here suggest as many as 11,000 to $K=29.5$.  } \\label{number-counts} \\end{figure}   An IFTS is tolerant of detector noise because it always operates under photon limited conditions due to the broad spectral bandpass transmitted to the FPA. This is illustrated in Table \\ref{readoutexamples} which shows a break-down of the noise sources in the NIR and MIR channels corresponding to the performance listed in Table \\ref{capabilities}.  Table \\ref{readoutexamples} also shows that the read-out rates required to avoid saturation are modest ($1-10$~mHz), since typical well depths for NIR InSb or HgCdTe arrays are a few $10^5$ e$^-$ and $10^7$ e$^-$ for MIR Si:As arrays.   \\begin{deluxetable}{llll} \\tablewidth{0pt} \\tablecaption{Signal and Noise Budget\\tablenotemark{a}} \\tablehead{ \\colhead{} & \\colhead{} & \\colhead{NIR Channel} &\\colhead{MIR Channel}} \\startdata &       & $F$ = 1 nJy & $F$ = 65 nJy \\\\ &       & $t$ = 1000 s & $t$ = 100 s \\\\ &       & $\\Delta\\lambda = 1-5\\mu$m & $\\Delta\\lambda = 5-15\\mu$m \\\\ \\hline Signal \\\\ & Source     &  610 & 2709   \\\\ & Background\\tablenotemark{b} & 3724 & 463669 \\\\ &            &      &        \\\\ Total Signal  &            & 4334 & 466378 \\\\ \\\\ Noise\\\\ & Signal Shot Noise  & 24.7 & 52.0 \\\\ & Background Shot Noise & 60.8 & 680.9 \\\\ & Dark Shot Noise   & 5.5  & 10.0 \\\\ & Read Noise        & 5    & 5 \\\\ Total Noise &           & 66.0 & 683.0 \\\\ \\enddata \\label{readoutexamples} \\tablenotetext{a}{In electrons} \\tablenotetext{b} {Background includes zodiacal foreground and thermal emission from the telescope as calculated as described in \\S \\ref{snrcal}.} \\end{deluxetable}   Similarly, orders of magnitude higher thermal emission from the instrument, or thermal radiation leaks from outside the instrument bay, can be tolerated compared to the case for dispersive spectrometers or fixed filter cameras.  As a pragmatic demonstration of this principle, the IFTS instruments LIFTIRS and HIRIS are routinely operated with ambient temperature optics in the 8-14 $\\mu$m band (Bennett et al. 1995), whereas dispersive spectrometers, like SEBASS (Bennett, {\\it private communication}), operating in the same spectral region must have the slit and all following optics cooled far below ambient temperatures.  The reason is that in a dispersive spectrometer the thermal emission of all the elements and optics downstream of the slit reach the detector at full spectral range determined by the bandpass limiting element at or near the coldstop, whereas only the narrow spectral range corresponding to the width of a spectral channel for the signals of interest reach the detector pixels. For the IFTS, both the signals of interest and the thermal emission are seen over the full spectral range determined by the bandpass limiting filter, and thus it is only necessary that the thermal emission of the optical elements along the optic axis integrated over the bandpass of interest, either 1-5 $\\mu$m or 5-15 $\\mu$m, be somewhat less than that of the integral of the zodiacal foreground, telescope emission, and source signal level integrated over the same broad spectral region.  An IFTS is potentially immune to cosmic ray hits because the ``energy'' of a single upset pixel in one OPD frame appears as a sinusoidal signal divided among all bins in the spectral transform of the interferogram for that pixel.  We can ignore cosmic ray hits only if the counts generated are at or below our noise level. A minimum ionizing cosmic ray proton ($E \\simeq 1$ GeV) has ionization losses of $dE/dx \\simeq 400$ eV $\\mu$m$^{-1}$ in Si.  Assuming that 3.6 eV is required to produce an electron-hole pair, a cosmic ray will yield at least a few thousand events, since typical pixels have sensitive layers that are tens of $\\mu$m thick. We would obtain similar number for a hybrid device, i.e., InSb or HgCdTe on a Si multiplexer. If a cosmic ray hit produces a significant signal in a certain number of pixels, those pixels must be ``repaired'' by interpolating the interferogram between the previous few ``good'' frames, and the following few ``good'' frames which are not contaminated by cosmic ray hits. The same sort of processing would be needed for any other system as well, be it an imager or a spectrometer. Comparison with the noise sources listed in Table \\ref{readoutexamples} indicates that cosmic ray hits will have to be repaired in the NIR channel, while the MIR channel will be more tolerant.  A dual port design (Fig. \\ref{optical-layout}) delivers the complementary symmetric and antisymmetric interferograms.  In this dual-input dual-output design, the field of the complementary input (labeled ``Calibration Input'' in Fig. \\ref{optical-layout}) is also imaged and superimposed on each image of the ``Primary Input''.  This property is often used to cancel the sky emission.  In operation, when observing the sky in the primary input, the secondary input would be fed with a cold blackbody load, having negligible radiance.  The final interferogram is constructed from the difference between the two outputs (which is therefore also immune to common mode electrical noise) while the normalized ratio reveals systematic variation due to detector drifts.  The wavelength scale and the instrumental line shape (a $sinc$ function if there is no apodizing) are precisely determined and are independent of wavelength. Absolute wavelength calibration is done by counting fringes of an optical single-mode laser.  Compared to a dispersive system the broad-band operation of an IFTS means that there are $M$ times more photons for flat fielding and determining signal-dependent gain (linearity). Hence, high signal-to-noise calibration images can be acquired faster or with lower power internal sources.  \\subsection{Pixel Size, Spectral Resolution, and Field of View} \\label{pixelsize}  Spatial multiplexing renders the performance of an IFTS equal to that of an ideal multi-slit spectrograph (Bennett 1995). Hence, even if we ignore slit losses and blaze inefficiency the other advantages of an IFTS are overwhelming.  The spectral resolution can be varied arbitrarily from the coarsest case of a small number of bands up to a spectral resolution limit determined only by the maximum OPD characteristic of the instrument. The proposed instrument has a maximum OPD of 1 cm and hence can operate over a range of resolutions from full band up to $M$=8000 in the NIR  The spectral resolution limit, $R = k/\\delta k$, of a Michelson interferometer is  \\begin{equation} R = 8(d/ \\phi D)^2, \\end{equation}  \\noindent where $\\phi$ is the angular diameter of the FOV, $d$ is the diameter of the beam splitter, and $D$ the telescope primary mirror diameter (e.g., Jacquinot 1954; Maillard 1995). Classically, $\\phi$ refers to the entire field, but in the case of an IFTS, $\\phi$ is the FOV of an on-axis pixel.  Although convenient if a single fringe fills the FPA, just as with imaging Fabry-Perots, there is no reason why each pixel should record the same apparent wavenumber.  Fringes crowd together with increasing field angle.  Therefore, the need to maintain modulation efficiency over the entire field of view requires that the spatial separation of the fringes at the edge of the FPA, for a given retardance, is significantly greater than the pixel spacing.  If $x$ is the OPD for a normally incident beam with wavenumber $k$, and $\\theta$ is the field angle of off-axis rays at the beam splitter, then the path difference at $\\theta$ is $x_\\theta = x cos\\theta$ and the apparent wavenumber of this beam is  \\begin{equation} k_\\theta = k/cos\\theta. \\end{equation}  \\noindent The angles $\\theta$ and $\\phi$ are related by the angular magnification, $D/d$. If $\\delta \\theta$ is the angular width, also at the beam splitter, corresponding to a single pixel, the spectral resolution limit for off-axis points can be found by differentiating Eq (2),  \\begin{equation} 1/R_\\theta = \\delta k_\\theta/k_\\theta = tan\\theta \\delta \\theta, \\end{equation}  \\noindent Fig. \\ref{pixel_fov} shows the pixel size for a given field of view for a range of resolutions.  For example, for an 8~m diameter primary aperture and a beam splitter of diameter 10~cm, a FOV of $200''$, and a pixel size of $0.''05$, the corresponding angles at a beam splitter of diameter 10~cm are $2.^\\circ 2$ and $4''$ respectively, leading to a resolution limit of $R = 1.3 \\times 10^6$.  Since this resolution is two orders of magnitude greater than we are proposing, it is clear that spectral resolution is not the principal factor determining pixel size.  An alternative way to view this constraint is that $d$, i.e., the size of the optics, is determined not by spectral resolution, but by the requirement that there be no vignetting over the field of view. Thus, the optics for an IFTS are similar to that of a simple re-imaging camera, and are smaller and slower than those of an equivalent dispersive spectrograph.  \\begin{figure}[thb] \\centerline{\\psfig{figure=fig4.ps,width=3.4in,angle=90}} \\caption{\\small Pixel size as a function of the field of view required to spatially fully sample fringes at the edge of the FPA, and hence maintain modulation efficiency.  Curves are plotted for resolutions $R = k/\\delta k = 10^4$, $10^5$ and $10^6$ assuming an 8~m diameter primary aperture and a beam splitter of diameter 10~cm. } \\label{pixel_fov} \\end{figure}   We therefore have broad freedom to choose the pixel size by trading off field of view and spatial sampling.  Given that NIR arrays of $4096 \\times 4096 $ pixels are likely to be available in the near future, a pixel size of $0.''05$ yields a $3.'3$ field of view and $\\lambda / 2D$ sampling at 4 $\\mu$m.  This choice of pixel size does not preclude diffraction limited imaging at shorter wavelengths. If pixels have sharp boundaries, then it is possible to extract information at spatial frequencies above the cut-off in the pixel-sampling modulation transfer function if the spacecraft can offset and track at the sub-pixel level (cf. Fruchter and Hook 1998). Similar reasoning suggests $0.''1$ pixels would be a satisfactory compromise for the MIR channel. ",
        "conclusions": "An IFTS instrument can perform a wide variety of NGST science.  The advantages of the IFTS concept are:  \\begin{enumerate}  \\item Deep imaging acquired simultaneously with higher spectral resolution data over a broad wavelength range.  \\item ``Hands-off'', unbiased, multi-object, slitless spectroscopy (ideal for moving objects).  Efficient in confusion limit.  \\item Flexible resolution ($M=1-10,000$).  \\item High throughput (near 100\\%) dual-port design.  \\item Tolerant of cosmic rays, read-noise, dark current, and light leaks.  \\item Simple and reliable calibration.  High SNR determination of flat-fields and detector non-linearity.  \\item Compact, lightweight design. Slow reimaging optics.    \\end{enumerate}"
    },
    "9803/astro-ph9803119_arXiv.txt": {
        "abstract": "We present a study of 581 Hz oscillations observed during a thermonuclear X-ray burst from the low mass X-ray binary (LMXB) 4U 1636-54 with the Rossi X-ray Timing Explorer (RXTE). This is the first X-ray burst to exhibit both millisecond oscillations during the rising phase as well as photospheric radius expansion. We measure an oscillation amplitude within 0.1 s of the onset of this burst of $75 \\pm 17 \\%$, that is, almost the entire thermal burst flux is modulated near onset. The spectral evolution during the rising phase of this burst suggests that the X-ray emitting area on the neutron star was increasing, similar to the behavior of bursts from 4U 1728-34 with 363 Hz oscillations reported recently. We argue that the combination of large pulsed amplitudes near burst onset and the spectral evidence for localized emission during the rise strongly supports rotational modulation as the mechanism for the oscillations. We discuss how theoretical interpretation of spin modulation amplitudes, pulse profiles and pulse phase spectroscopy can provide constraints on the masses and radii of neutron stars. We also discuss the implications of these findings for the beat frequency models of kHz X-ray variability in LMXB. ",
        "introduction": "Large amplitude millisecond oscillations have now been observed during thermonuclear X-ray bursts from six low mass X-ray binary (LMXB) systems with the Rossi X-ray Timing Explorer (RXTE) (see \\markcite{SZS}Strohmayer, Zhang \\& Swank 1997; \\markcite{SMB}Smith, Morgan \\& Bradt 1997, \\markcite{Z96}Zhang {\\it et al.} 1996; \\markcite{Swank97}Swank {\\it et al.} 1997; and \\markcite{Stroh97}Strohmayer {\\it et al.} 1997). The thermonuclear instability which triggers an X-ray burst burns in a few seconds the nuclear fuel which has been accumulated on the neutron star surface over several hours. This $>$ 10$^3$ difference between the accumulation and burning timescales means that it is extremely unlikely that the conditions required to trigger the instability will be achieved simultaneously over the entire stellar surface. This realization, first emphasized by \\markcite{Joss78}Joss (1978), led to the study of lateral propagation of the burning instability over the neutron star surface (see \\markcite{FW82}Fryxell \\& Woosley 1982, \\markcite{NIF84}Nozakura, Ikeuchi \\& Fujimoto 1984, and \\markcite{B95}Bildsten 1995). The subsecond risetimes of thermonuclear X-ray bursts suggests that convection plays an important role in the physics of the burning front propagation, especially in the low accretion rate regime which leads to large ignition columns (see \\markcite{B98}Bildsten (1998) for a review of thermonuclear burning on neutron stars). \\markcite{B95}Bildsten (1995) has shown that pure helium burning on neutron star surfaces is in general inhomogeneous, displaying a range of behavior which depends on the local accretion rate. Low accretion rates lead to convectively combustible accretion columns and standard type I bursts, while high accretion rates lead to slower, nonconvective propagation which may be manifested in hour long flares. These studies emphasize that the physics of thermonuclear burning is necessarily a multi-dimensional problem and that {\\it localized} burning is to be expected, especially at the onset of bursts.  There is now good evidence that the oscillations seen during the rising phase of bursts from  4U 1728-34 are produced by spin modulation of such a localized thermonuclear hotspot on the surface of the neutron star, and that the observed oscillation frequency is a direct measure of the neutron star spin frequency (see \\markcite{SZS}Strohmayer, Zhang \\& Swank 1997). These observations provide the most compelling evidence to date that neutron stars in LMXB are rotating with near millisecond periods.  In this Letter we present new burst data from the LMXB 4U 1636-54 which provides further evidence in support of the spin modulation hypothesis for the millisecond burst oscillations. We present data from a thermonuclear burst from 4U 1636-54 which reveals a strong transient oscillation during the burst rise at 1.723 ms with an initial amplitude of $75 \\pm 17 \\%$. We also discuss the implications of these findings for the spin modulation interpretation and how they can be used to place constraints on the mass and radius of neutron stars. Finally, we discuss some implications of our observations for the current theories of kilohertz quasiperiodic oscillations (QPO) in LMXB. ",
        "conclusions": "If the millisecond oscillations seen during bursts are in fact due to spin modulation, then detailed study and modelling of the oscillation amplitudes, pulse profiles and spectral variability with pulse phase during X-ray bursts can provide a wealth of information on the mass and radius of the neutron star. For example, the maximum modulation amplitude that can be obtained from a hotspot of a given angular size on a rotating neutron star is set by the strength of general relativistic light bending. For the case that rotation is not rapid enough to substantially distort the exterior spacetime from the Schwarzchild spacetime, and this is the case even for spin periods of a few milliseconds (Lamb \\& Miller 1995), then the maximum amplitude depends only on the compactness of the neutron star, that is, the ratio of stellar mass to radius $GM/c^2R$. Stars which are more compact produce lower amplitudes due to flattening of the pulse by light bending (see \\markcite{PFC}Pechenick, Ftaclas \\& Cohen 1983; \\markcite{S92}Strohmayer 1992; and \\markcite{ML97}Miller \\& Lamb 1997). Since the intrinsic rotational modulation amplitude can only be decreased by other effects such as photon scattering (see \\markcite{MLP}Miller, Lamb \\& Psaltis 1997; \\markcite{BL}Brainerd \\& Lamb 1987; and \\markcite{KP}Kylafis \\& Phinney 1989) or the viewing geometry of the spot, the maximum observed oscillation amplitude represents a lower limit to the intrinsic amplitude. Thus an observed amplitude can be used to place an upper limit on the compactness of the neutron star, that is, if the star were more compact than some limit it would not be able to produce a modulation amplitude as large as that observed.  In principle, stellar rotation will also play a role in the observed properties of spin modulation pulsations. For example, assuming the oscillation frequency of 581 Hz represents the spin frequency of the neutron star in 4U 1636-54, then for a 10 km radius neutron star the spin velocity is $v_{spin}/c = 2\\pi \\nu_{spin} R  \\approx 0.12$ at the rotational equator. The motion of the hotspot produces a Doppler shift of magnitude $\\Delta E / E \\approx v_{spin}/c = 0.12$, thus the observed spectrum is a function of pulse phase (see \\markcite{CS}Chen \\& Shaham 1989). Measurement of a pulse phase dependent Doppler shift in the X-ray spectrum would provide additional evidence supporting the spin modulation model and would also provide a means of constraining the neutron star radius. The rotationally induced velocity also produces an aberration which results in asymmetric pulses, thus the pulse shapes also contain information on the spin velocity and therefore the stellar radius (\\markcite{CS}Chen \\& Shaham 1989). The component of the spin velocity along the line of site is proportional to $\\cos\\theta$, where $\\theta$ is the lattitude of the hotspot measured with respect to the rotational equator. The modulation amplitude also depends on the lattitude of the hotspot, as spots near the rotational poles produce smaller amplitudes than those at the equator. Thus we expect a correlation between the observed oscillation amplitude and the size of any pulse phase dependent Doppler shift. Dectection of such a correlation in a sample of bursts would definitively confirm the rotational modulation model in our opinion. We will present calculations of mass - radius constraints for 4U 1636-54 and 4U 1728-34 including the effects of light bending, rotation and angle dependent emission, based on the observed properties of burst oscillations as well as spectroscopy of Eddington limited bursts in a subsequent paper.  Several models for the kilohertz QPO seen in 13 LMXB systems (see \\markcite{vdk}van der Klis 1997 for a recent review) invoke some sort of beat-frequency interpretation for the twin kHz peaks seen in many of the sources (see \\markcite{MLP}Miller, Lamb \\& Psaltis 1997; \\markcite{S96}Strohmayer {\\it et al.} 1996). So far only in 4U 1728-34 does the frequency difference between the twin kHz peaks match the frequency observed during X-ray bursts (see \\markcite{S96}Strohmayer {\\it et al.} 1996). In two other sources (KS 1731-26 and 4U 1636-54) the separation of the twin kHz peaks is closer to 1/2 the frequency of oscillations observed during bursts (see \\markcite{WV}Wijnands \\& van der Klis 1997; and \\markcite{Z96}Zhang {\\it et al.} 1996). For example, \\markcite{W97}Wijnands {\\it et al.} (1997) report a frequency difference for the twin kHz QPO in 4U1636-54 of $276 \\pm 10$ Hz. The effort to reconcile these observations with a beat-frequency interpretation has led to speculation that the oscillation frequency observed during bursts may sometimes be twice the spin frequency of the star, although in 4U1636-54 the difference frequency appears to be a bit less than 1/2 the burst oscillation frequency, and in some Z sources the frequency difference is not constant (\\markcite{van97}van der Klis {\\it et al.} 1997). If this scenario is correct it implies the existence of two antipodal spots on the neutron star surface during X-ray bursts. In addition to the daunting requirement of initiating the thermonuclear flash nearly simultaneously on opposite sides of the neutron star, the observation of large oscillation amplitudes shortly after burst onset in 4U1636-54 places severe constraints on the two hotspot scenario. To see this one can ask the following question. What maximum amplitude can be produced by a star with antipodal hotspots? For two antipodal spots light bending strongly constrains the amplitudes that can be achieved (see \\markcite{PFC}Pechenick, Ftaclas \\& Cohen 1983). Calculations using the Schwarzchild spacetime and isotropic emission from the stellar surface indicate that even a neutron star with an implausibly small compactness of $M/R = 0.1$, recall that rotational modulation amplitude increases with decreasing compactness, can only achieve a maximum amplitude of about 30 \\%, whereas we measured an amplitude of 75\\% from the burst described here. We note that with $M/R = 0.1$ a 1.4 $M_{sun}$ neutron star would have a radius of 21 km. This is far stiffer than any neutron star equation of state that we are familiar with. Rapid rotation could in principle modify this result, but at the modest inferred spin period of 290.5 Hz for the two spots to produce the observed 581 Hz frequency one would still require implausibly large neutron star radii to make a significant rotational correction to the amplitude. Another process which can increase the amplitude is beaming of radiation at the stellar surface. For example, to the extent that electron scattering is the dominant opacity process in neutron star atmospheres then one should expect a specific intensity distribution which approximates that from a grey atmosphere. Such a distribution is proportional to $\\cos\\delta + 2/3$, where $\\delta$ is the angle from the normal to the stellar surface (see \\markcite{Mihal}Mihalas 1978), so this will modestly increase the rotational modulation amplitude. In addition, it is likely that the emergent spectrum will also be a weak function of $\\delta$ (\\markcite{ML97}Miller \\& Lamb 1997). With more detailed modelling of the emission from hotspots of finite size it will be possible to place strong constraints on the antipodal hotspot interpretation for the burst oscillations. Finally, if further analysis such as pulse phase spectroscopy continues to support the interpretation of the burst oscillation frequency as the neutron star spin frequency in 4U1636-54, then kHz QPO frequency separations near 1/2 the spin frequency would suggest that the beat frequency interpretation may be untenable."
    },
    "9803/astro-ph9803096_arXiv.txt": {
        "abstract": "{ To examine the evidence for hierarchical evolution on mass scales of $\\sim 10^{13}$-$10^{14} \\Mdot$, we apply a statistic that measures correlations between galaxy velocity and projected position (Dressler \\& Shectman 1988) to data for six poor groups of galaxies, HCG 42, HCG 62, NGC 533, NGC 2563, NGC 5129, and NGC 741.  Each group has more than 30 identified members (Zabludoff \\& Mulchaey 1998ab).  The statistic is sensitive to clumps of galaxies on the sky whose mean velocity and velocity dispersion deviate from the kinematics of the group as a whole.  The kinematics of galaxies within $\\sim 0.1$\\lith\\inv\\ Mpc of the group center do not deviate from the global values, supporting our earlier claim that the group cores are close to virialization or virialized.  We detect significant substructure (at $\\geq 99.9$\\% confidence) in the two groups with the most confirmed members, HCG 62 and NGC 741, that is attributable mostly to a subgroup lying $\\sim 0.3$-0.4\\lith\\inv\\ Mpc outside of the core. We conclude that at least some poor groups, like rich clusters, are evolving via the accretion of smaller structures from the field. With larger poor group surveys, the incidence of such accretion and the distribution of subgroup masses are potential constraints of cosmological models on mass scales of $\\simless 10^{13}$-$10^{14} \\Mdot$ and on physical scales of $\\simless 0.5$\\lith\\inv\\ Mpc.  \\vskip 0.5cm \\noindent{\\it Subject headings}:  galaxies: clustering --- cosmology: large-scale structure of Universe }  \\vfill\\eject ",
        "introduction": "The evolution of structure on different mass scales is one of the outstanding issues in cosmology.  For example, although galaxy clusters of $\\sim 10^{15} \\Mdot$ (including Virgo (Binggeli 1993; Bohringer \\etal 1994), Coma (Mellier \\etal 1988; Briel \\etal 1992; White \\etal 1993), and Abell 754 (Zabludoff \\& Zaritsky 1995)) are clearly evolving from the accretion of smaller groups, it is uncertain whether poor groups of $\\sim 10^{13}$-$10^{14} \\Mdot$ also evolve hierarchically.  There is some indirect evidence that the evolution of poor groups is similar to rich clusters; the galaxies and hot gas in poor groups follow the same relationships found among the X-ray temperature, X-ray luminosity, and galaxy velocity dispersion for rich clusters (Mulchaey \\& Zabludoff 1998, hereafter MZ98).  Historically, however, the number of known poor group members has been too small to examine individual groups for direct evidence of hierarchical evolution.  Multi-object spectroscopy now makes it possible to obtain ``cluster-size'' samples of galaxies in poor groups and to identify substructure, if it exists, in the same manner as for rich clusters. Substructure in clusters was not detected until the number of spectroscopically-confirmed cluster members exceeded $\\sim 30$-50 galaxies.  Recent poor group surveys have reached this membership level (Zabludoff \\& Mulchaey 1998ab; hereafter ZM98a and b).  The discovery of substructure in poor groups would provide new cosmological constraints by establishing that hierarchical evolution is occurring on mass scales of $\\sim 10^{13}$-$10^{14} \\Mdot$ and on physical scales of $\\sim 0.5$\\lith\\inv\\ Mpc.  The detection of substructure would also support the picture in which at least some poor groups evolve as low-mass analogs to rich clusters.  In this Letter, we describe the results from applying a substructure statistic (Dressler \\& Shectman 1988; hereafter DS88) to the six best sampled poor groups in ZM98ab, the most detailed spectroscopic survey of poor groups to date. ",
        "conclusions": "To search for substructure in the six poor groups, we use the method applied by DS88 to rich clusters.  This test identifies a fixed number of nearest neighbors on the sky around each galaxy, calculates the local mean velocity and velocity dispersion of the subsample, and compares these values with the mean velocity and velocity dispersion of the entire group.  The kinematic deviations of the subsamples from the global values are summed.  This sum is larger for a group with a kinematically distinct subgroup than for a similar group without substructure.  For each galaxy $i$, the deviation of its nearest projected neighbors from the kinematics of the group as a whole is defined as ${\\delta_i} \\equiv (n^{1/2}/ \\sigma_r) \\lbrack (\\upsilon_{loc} - \\overline \\upsilon)^2 + (\\sigma_{r,loc} - \\sigma_r)^2 \\rbrack ^{1/2}$, where $\\overline \\upsilon$ is the mean velocity for the group, $\\sigma_r$ is the group velocity dispersion, and $n$ is the number of nearest neighbors (including the galaxy) used to determine the local mean velocity $\\overline \\upsilon_{loc}$ and local velocity dispersion $\\sigma_{r,loc}$. The total deviation for the group is defined as the sum of the local deviations, $\\Delta \\equiv \\sum |\\delta_i|$ for all $i \\leq N_{grp}$, the number of group members. As pointed out in DS88, the $\\Delta$ statistic is similar to the $\\chi^2$ statistic, except that the $\\delta_i$'s are not squared before summation in order to reduce the contributions of the largest, rarest deviations. If the galaxy velocity distribution of the group is close to Gaussian, and the local variations are only random fluctuations, $\\Delta \\simeq N_{grp}$.  To calculate $\\delta$, we choose $n = 11$ (as in DS88). This choice allows robust determinations of $\\overline \\upsilon_{r,loc}$ and $\\sigma_{r,loc}$.  Silverman (1986) argues that using $n \\sim {N_{grp}}^{1/2}$ nearest neighbors (= 6-8 for these groups) maximizes the sensitivity of such a test to small scale structures while reducing its sensitivity to fluctuations within the Poisson noise (also see Bird 1994b). To check the robustness of the $n=11$ assumption, we compare the results below with those for $n=6$.  The conclusions drawn from the $n=6$ and $n=11$ cases are the same.  Calibration of the $\\Delta$ statistic for each group is required because 1) the $\\delta_i$'s are not statstically independent and 2) the velocity distribution may not be intrisically Gaussian even if there are no subgroups ({\\it e.g.}, the group members may follow predominantly circular or radial orbits). We determine the significance of the observed $\\Delta$ by comparing it with the results of 1000 Monte Carlo trials in which galaxy velocities are drawn randomly from the observed distribution and assigned to galaxy positions.  This scrambling technique effectively destroys any substructure (DS88). If the probability is low that a group without substructure has a $\\Delta$ value at least as large as that observed, then we consider the substructure detection significant.  In two groups, HCG 62 and NGC 741, the observed value of $\\Delta$ is significant at the $\\geq 99.9\\%$ confidence level. Such high $\\Delta$ values might arise from substructure, but also could result from smooth variations in the group's velocity field ({\\it e.g.}, rotation or velocity shear (Malumuth \\etal 1992)) and/or from a dependence of $\\sigma_r$ on radius (Bird 1994b). ZM98a show that $\\sigma_r$ is constant out to radii of $\\sim 0.5$\\lith\\inv\\ Mpc in the combined velocity dispersion profile for the sample groups. To determine whether substructure is in fact responsible for the high $\\Delta$ values in HCG 62 and NGC 741, we examine the local deviation $\\delta$ for each group member. A concentration of large $\\delta$ values on the sky indicates a kinematically distinct subgroup.  Figure 1 shows the projected spatial distribution of group members for each group (top panel). The second panel shows this distribution with the radii of the circles weighted by $e^\\delta$ (as in DS88). Because each point is not statistically independent, a few very deviant galaxies can boost the $\\delta$ values for a large number of nearby points.  Therefore, a visual comparison with the Monte Carlo simulations is required to assess the significance of any structures. The third panel shows the results of the Monte Carlo trial (out of 1000) with the largest $\\Delta$ value, or greatest total deviation. The results of the trial with the median $\\Delta$ value are in the bottom panel.  The significance of the seven large clustered circles to the northeast of HCG 62 and five large clustered circles to the south of NGC 741 is high.  In each case, the large $\\delta$ values show that a subgroup not in equilibrium with the global group potential is the principal source of the significant $\\Delta$ value. Each of the two subgroups lies a projected distance of $\\sim 0.3$-0.4\\lith\\inv\\ Mpc outside of the group core.  On the other hand, the clustering of small circles within $\\sim 0.1$\\lith\\inv\\ Mpc of all of the group centers indicates that the core mean velocity and velocity dispersion are similar to the global values for the group.  This result suggests that the group cores are close to virialization or virialized and is consistent with the conclusions from our earlier studies of group dynamics (ZM98a, MZ98)."
    },
    "9803/astro-ph9803089_arXiv.txt": {
        "abstract": "Aperiodic variability and Quasi Periodic Oscillations (QPOs) are observed from accretion disks orbiting white dwarfs, neutron stars, and black holes, suggesting that the flow is universally broken up into discrete blobs. We consider the interaction of these blobs with the magnetic field of a compact, accreting star, where diamagnetic blobs suffer a drag. We show that when the magnetic moment is not aligned with the spin axis, the resulting force is pulsed, and this can lead to resonance with the oscillation of the blobs around the equatorial plane; a resonance condition where energy is effectively pumped into non--equatorial motions is then derived. We show that the same resonance condition applies for the quadrupolar component of the magnetic field. We discuss the conditions of applicability of this result, showing that they are quite wide. We also show that realistic complications, such as chaotic magnetic fields, buoyancy, radiation pressure, evaporation, Kelvin--Helmholtz instability, and shear stresses due to differential rotation do not affect our results. In accreting neutron stars with millisecond periods, we show that this instability leads to Lense--Thirring precession of the blobs, and that damping by viscosity can be neglected. ",
        "introduction": "The interaction between the magnetic field anchored to rotating bodies and matter orbiting around them plays an important role in a variety of astrophysical situations, ranging from the Jupiter-Io system (Stern and Ness 1982) to the Jovian ring (Burns \\etal\\/ 1985), to protoplanetary disks around young magnetic stars (Bodenheimer 1995) and accreting degenerate stellar remnants in binary systems, such as white dwarfs in intermediate polars or neutron stars in X-ray binaries (Frank, King and Raine 1995). Straightforward applications of the laws of electrodynamics are possible in cases, such as the Jovian ring, in which the orbiting matter is made of solid, dust particles. Due to a variety of poorly known MHD and plasma effects, modelling this interaction in the case of gaseous disks that orbit different classes of magnetic stars is far more  uncertain and difficult. In particular, theoretical descriptions of viscous accretion disks around magnetic rotating stars have been largely based on a number of simplifying assumptions. Two of these are especially relevant to the present work: (a) that the magnetic field (assumed dipolar) and rotation axes of the star are coaligned;  (b) that the disk is continuous, smooth, azimuthally symmetric and lies in the equatorial plane of the rotating star. For example Ghosh and Lamb (1979) adopted these assumptions and introduced an {\\it ad hoc} effective diffusivity of the disk plasma to derive a stationary disk solution, in which the star magnetic field lines thread the disk and slip across its material. Models of this kind proved very useful for evaluating the basic properties of the disk-magnetic field interaction, such as the size of the magnetosphere and the balance of the material and non-material torques.  It has always been recognised that the coalignement of the magnetic field and rotation axes of the star is an unrealistic hypothesis as a finite magnetic colatitude is required for the generation of the periodic signals at the star spin frequency  that are often observed in these systems. On the other hand, the pronounced aperiodic (or, in some cases, quasi-periodic) variability that is frequently detected on a variety of timescales ranging from days to milliseconds in disk accreting compact stars of all classes (intermediate polars: King and Lasota 1990; accreting neutron stars and black holes: van der Klis 1995, van der Klis 1997) testify that the disk cannot be regarded as smooth, continuous and azimuthally symmetric. In a number of cases the quasi-periodic timing signature of blobs orbiting over a limited range of radii close to the compact star is clearly seen. A highly inhomogeneous and clumpy disk, possibly comprising two distinct and coexisting phases (hot and cold), is also envisaged as the end product of the instabilities predicted by current models of viscous disks (e.g the so-called secular and viscous instabilities of radiation-pressure dominated $\\alpha$-disks, Lightman and Eardley 1974, Shakura and Sunyaev 1976, or the magnetic amplification and buoyancy of flux tubes in dynamo-driven disks, Vishniac and Diamond 1992), as discussed by Krolik (1998).  In all accreting objects endowed with a strong magnetic field, the presence of inhomogeneities and/or blobs that can be regarded as discrete entities suggests that a novel mode of interaction between the compact star and the disk is possible. This occurs because individual blobs are most likely strongly diamagnetic (see for instance King 1993, Wynn and King 1995 and references therein); when moving through a magnetic field, strong surface currents develop on the blob the main effect of which is the generation of a drag opposite to the  component of the blob velocity perpendicular to the field (Drell, Foley and Ruderman 1965). The acceleration acting on each blob is thus \\begin{equation} \\vec{a} = -\\frac{\\vec{v}_\\perp^{(rel)}}{t_d}\\;; \\end{equation} the drag time--scale $t_d$ is given by \\begin{equation} t_d = \\frac{c_A m}{B^2 l^2}\\;, \\end{equation} where $c_A$ is the Alfv\\`en speed in the magnetic field $B$, and $m$ and $l$ are the mass and characteristic radius of the blob. Here $\\vec{v}^{(rel)}$ is the relative velocity between the blob and a magnetic field line. It is easy to see (Fig.1) that this acceleration has a component along the disk axis which tends to lift the blob off the equatorial plane where it is, at least initially, lying. It is the purpose of this paper to study this dynamical interaction, and to show how this alters the conventional view of accretion disks onto magnetized compact stars. In order to bring out the physical meaning of the instability most clearly, we shall at first idealize these plasma blobs as point masses. In the next section, it will be shown that this interaction leads to the lifting of blobs off the equatorial plane, at a resonance radius; the conditions under which this result applies are discussed in Section 3. In Section 4, we relax the hypothesis that blobs are point masses, and establish that a variety of effects, all related to the blobs having a finite size, do not modify our results. As the only concrete application of this instability, we discuss in Section 5 the generation of modulation of the X--ray flux in Low Mass X--ray Binaries (LMXBs) exhibiting millisecond QPOs at frequencies comparable to the single--particle Lense--Thirring (1918) precession frequency. The last section summarizes the results. ",
        "conclusions": "The processes described above are generic: they apply to every accretion disk broken up into discrete units, surrounding a magnetized non--aligned rotator provided the ordering of time--scales (Eq. \\ref{ordering}) holds and the blobs have density contrasts $(\\rho_b-\\rho_d)/ \\rho_d \\ga 1$. The instability is independent of the detailed diamagnetic properties of the blobs (here enshrined in the parameter $t_d$ which is just modulated at the relative frequency $\\omega_K-\\omega_s$), of their masses and densities, their interactions with radiation emitted by the accreting source, and the exact form of the viscosity. Provided the hierarchy of timescales $t_K \\ll t_d \\ll t_v$ (Eq. \\ref{ordering}) is established and blobs have non--negligible density contrasts, it should apply to both accreting white dwarfs and neutron stars.  The existence of resonances is essentially due to the fact that a spatially periodic magnetic field, such as that due to a single multipole, will produce a temporally--variable electromagnetic force on a single blob, as it passes through the various frequency components of the field. Minor corrections to the exact locations of the resonances might derive from effects we opted not to consider, such as an axially offset field, or a rotation rate different for different multipoles, as it happens to the non--dipolar part of the magnetic field of the Earth. Another phenomenon we neglected to investigate, and which might generate further resonances, is the non--linear coupling of radial, latitudinal and longitudinal motions. It should however be borne in mind that the exact locations of resonances might be irrelevant since blobs are expected to have a finite size.  The evolution of the accretion disk after the resonance of Eq. \\ref{resonance} is very difficult to predict; it seems clear enough that blobs will tend to move closer to the magnetic equatorial plane, and oscillate around it, but the further evolution is uncertain because the details depend on a very poorly known quantity, the $z$--axis viscosity. This is the reason why we did not push the analysis performed in this paper into the non--linear regime. It seems quite clear that the disk must puff--up, but whether, and when, viscosity will manage to bring it back to the equatorial plane remains an open question. Very interesting consequences of the above--discussed instability occur in a specific class of accretors, \\ie\\/ neutron stars exhibiting millisecond QPOs. Blobs that have acquired $z$~axis motions will be acted upon by the torque which causes Lense--Thirring (1918, LT) precession. We have shown here that, in both Atoll and Z--type sources, viscosity is surely unable to bring blobs back down to the equatorial plane. For motions in the plane, the net effect of the instability is to induce epicyclic motions, or, given that in $1/r$ potentials all orbits close, to perturb the blobs' motions into ellipses. Then periastron precession induced by general relativistic effects ought to ensue.  A related mechanism for the generation of off--plane motions has been proposed for the Jovian dust ring (Burns \\etal\\/ 1985). It differs from the present one in that the individual units are not blobs of plasma, but single dust grains with nonzero electric charge, which are then acted upon by the normal Lorenz force $q \\vec{v}^{(rel)}\\wedge\\vec{B}/c$. It is easy to see that this differs from our case because the forces differ: in our case, the acceleration (Eq. 1) is in the plane of the two vectors $\\vec{v}^{(rel)}$ and $\\vec{B}$, while in the Jovian case the force is perpendicular to this plane. This difference has as a consequence that, in the Jovian case, each magnetic field multipole has its own, single resonance, while in our case the two multipoles we explored have the same array of resonances, and the array is an infinite one.  In short, in this paper we have shown that a simple interaction lifts blobs off the equatorial plane of a viscous accretion disk, provided the accreting object is not an aligned rotator. This occurs in a thin radial annulus, identified by Eq. \\ref{resonance}. This conclusion is rather general, being independent of the exact form of the drag time--scale, of the blobs' diamagnetic properties, and also of the assumed form for the disk viscosity. We have also shown why viscosity is unable to damp these motions (Section \\ref{lt}), and that this leads to modulation of the X--ray flux at the single particle precession frequency, or possibly twice as much. Thus, the proposed identification of some QPOs in the spectra of both Atoll and Z sources as due to LT--precession (Stella and Vietri 1998) is made more likely by the existence of a relevant mechanism exciting blobs' motions off the equatorial plane.  Helpful comments from F.K. Lamb and J. Imamura are gratefully acknowledged. This work was begun while one of us (M.V.) was visiting the Institute for Advanced Study; its Director, John Bahcall, is thanked for the warm hospitality. The work of L.S. was partially supported through ASI grants."
    },
    "9803/astro-ph9803040_arXiv.txt": {
        "abstract": "With the new generation of instruments for Cosmic Microwave Background (CMB) observations aiming at an accuracy level of a few percent in the measurement of the angular power spectrum of the anisotropies, the study of the contributions due to secondary effects has gained impetus. Furthermore, a reinvestigation of the main secondary effects is crucial in order to predict and quantify their effects on the CMB and the errors that they induce in the measurements. \\par In this paper, we investigate the contribution, to the CMB, of secondary anisotropies induced by the transverse motions of clusters of galaxies. This effect is similar to the Kaiser--Stebbins effect. In order to address this problem, we model the gravitational potential well of an individual structure using the Navarro, Frenk \\& White profile. We generalise the effect of one structure to a population of objects predicted using the Press-Schechter formalism. We simulate maps of these secondary fluctuations, compute the angular power spectrum and derive the average contributions for three cosmological models. We then investigate a simple method to separate this new contribution from the primary anisotropies and from the main secondary effect, the Sunyaev-Zel'dovich kinetic effect from the lensing clusters. \\par ",
        "introduction": "During the next decade, several experiments are planned to observe the Cosmic Microwave Background (CMB) and measure its temperature fluctuations (Planck surveyor, Map, Boomerang, ...). Their challenge is to measure the small scales anisotropies of the CMB (a few arcminutes up to ten degrees scale) with sensitivities better by a factor 10 than the COBE satellite (Smoot et al. 1992). These high sensitivity and resolution measurements will tightly constrain the value of the main cosmological parameters (Kamionkowski et al. 1994). However, the constraints can only be set if we are able to effectively measure the {\\it primary} temperature fluctuations. These fluctuations, present at recombination, give an insight  into the early universe since they are directly related to the initial density perturbations which are the progenitors to the cosmic structures (galaxies and galaxies clusters) in the present universe; but which are first and foremost the relics of the very early initial conditions of the universe.\\\\ Between recombination and the present time, the CMB photons could have undergone various interactions with the matter and structures present along their lines of sight. Some of these interactions can induce additional temperature fluctuations called, {\\it secondary} anisotropies because they are generated after the recombination. Along a line of sight, one measures temperature fluctuations which are the superposition of the {\\it primary } and {\\it secondary} anisotropies. As a result, and in the context of the future CMB experiments, accurate analysis of the data will be needed in order to account for the foreground contributions due to the secondary fluctuations. Photon--matter interactions between recombination and the present time are due to the presence of ionised matter or to variations of the gravitational potential wells along the lines of sight. \\par The CMB photons interact with the ionised matter mainly through Compton interactions. In fact, after recombination the universe could have been re-ionised globally or locally. Global early re-ionisation has been widely studied (see Dodelson \\& Jubas 1995 for a recent review and references therein). Its main effect is to either smooth or wipe out some of the primary anisotropies; but the interactions of the photons with the matter in a fully ionised universe can also give rise to secondary anisotropies through the Vishniac effect (Vishniac 1987). This second order effect has maximum amplitudes for a very early re-ionisation. The case of a late inhomogeneous re-ionisation and its imprints on the CMB fluctuations has been investigated (Aghanim et al. 1996) and found to be rather important. In this case, the secondary anisotropies are due to the bulk motion of ionised clouds with respect to the CMB frame. When the re-ionisation is localised in hot ionised intra-cluster media the photons interact with the free electrons. The inverse Compton scattering between photons and electrons leads to the so-called Sunyaev-Zel'dovich (hereafter SZ) effect (Sunyaev \\& Zel'dovich1972, 1980). The Compton distortion due to the motion of the electrons in the gas is called the thermal SZ effect. The kinetic SZ effect is a Doppler distortion due to the peculiar bulk motion of the cluster with respect to the Hubble flow. The SZ thermal effect has the unique property of depressing the CMB brightness in the Rayleigh-Jeans region and increasing its brightness above a frequency of about 219 GHz. This frequency dependence makes it rather easy to observe and separate from the kinetic SZ effect. In fact, the latter has a black body spectrum which makes the spectral confusion between kinetic SZ and primary fluctuations a serious problem. The SZ effect has been widely studied for individual clusters and for populations of clusters. For full reviews on the subject we refer the reader to two major articles: Rephaeli 1995 and Birkinshaw 1997. These investigations have clearly shown that the SZ effect in clusters of galaxies provides a powerful tool for cosmology through measurements of the Hubble constant, the radial peculiar velocity of clusters and consequently the large scale velocity fields. \\par Besides the interactions with the ionised matter, some secondary effects arise when the CMB photons traverse a varying gravitational potential well. In fact, if the gravitational potential well crossed by the photons evolves between the time they enter the well and the time they leave it, the delay between entrance and exit is equivalent to a shift in frequency, which induces a temperature anisotropy on the CMB. This effect was first studied by Rees \\& Sciama (1968)  for a potential well growing under its own gravity. Numerous authors have investigated the potential variations due to collapsing objects and their effect on the CMB (Kaiser 1982, Nottale 1984, Martinez-Gonz\\'alez, Sanz \\& Silk 1990, Seljak 1996). Similarly, a gravitational potential well moving across the line of sight is equivalent to a varying potential and will thus imprint secondary fluctuations on the CMB. This effect was first studied for one cluster of galaxies by Birkinshaw \\& Gull (1983) (Sect. 2). Kaiser \\& Stebbins (1984) and Bouchet, Bennett \\& Stebbins (1988) investigated a similar effect for moving cosmic strings. Recent work (Tuluie \\& Laguna 1995, Tuluie, Laguna \\& Anninos 1996) based on N-body simulations has pointed out this effect in a study of the effect of varying potential on rather large angular scales ($\\simeq 1^{\\circ}$). A discussion of some of these results and a comparison with ours will follow in the next sections. \\par In this paper, following the formalism of Birkinshaw \\& Gull (1983) and Birkinshaw (1989), we investigate the contribution of secondary anisotropies due to a population of collapsed objects moving across the line of sight, these objects range from small groups to rich clusters in scale ($10^{13}$ to $10^{15}$ $M_{\\sun}$). In section 2., we first study in detail the case of a unique collapsed structure. We use a structure model to compute in particular the deflection angle and derive the spatial signature of the moving lens effect. We then account (Section 3.) for the contribution, to the primordial cosmological signal, of the whole population of collapsed objects using predicted counts and we simulate maps of these secondary anisotropies. In section 4., we analyse the simulated maps and present our results. We give our conclusions in section 5. ",
        "conclusions": "In our work, we investigate the secondary fluctuations induced by moving lenses with masses ranging from those of groups of galaxies to those of clusters of galaxies in a simple way, based on predicted structure counts and simulated maps. This method allows us to explore a rather wide range of scales ($>10$ arcseconds) in various cosmological models. The analysis, in terms of angular power spectra, show the scales for which the primary fluctuations are dominant (Fig. \\ref{fig:pspec}). In the standard and lambda CDM models, the primary anisotropies are dominant respectively for scales $l<4000$ and $l<4500$ whereas in the Open CDM model they are dominant for $l<6000$. In practice, it is thus impossible to detect the secondary anisotropies due to moving lenses in the open model. The standard CDM model shows the smallest cut-off scale with an intermediate SZ kinetic pollution, compared to the other two models. It is therefore the ``best case'' framework for making an analysis and predicting the detection of fluctuations and the contributions that they induce. One must keep in mind that the results quoted in this particular case represent the ``best'' results we get from the analysis. \\par The results of our analysis are obtained under the assumption of a universe that never re-ionises, which is of course not the case. The re-ionisation, if it is homogeneous, is supposed to somewhat ease the task of extraction of the pattern. In fact, its main effect is to damp the angular power spectrum of the primary anisotropies on small scales, shifting the cut-off towards larger scales. In this case, the effect of moving lenses dominates over the CMB fluctuations, and the SZ kinetic is not as high as it is on very small scales. However, if the re-ionisation is late and inhomogeneous, it generates additional SZ kinetic-type secondary fluctuations (Aghanim et al. 1996) without damping the power spectrum by more than a few percent. Here, the re-ionisation might worsen the analysis at small scales. In any case, there could be some other additional secondary fluctuations principally due to the Vishniac effect, that arise in a re-ionised universe. Our work thus gives a ``best case'' configuration of the problem, with all other effects tending to worsen the situation. \\par We found that the secondary fluctuations induced by the moving gravitational lenses can be as high as $1.5\\,10^{-5}$; with {\\it rms} contributions of about 5 to $3.\\,10^{-7}$ in the three cosmological models. Even if the moving lens fluctuations have a particular dipolar pattern and even if they are ``perfectly'' located through their SZ thermal effect, the detection of the moving lens effect and its separation from the SZ kinetic and primary fluctuations are very difficult because of the very high level of confusion, on the scales of interest, with the point--like SZ kinetic anisotropies and because of spectral confusion. \\par We nevertheless analysed the simulated maps using an adap\\-ted wavelet technique in order to extract the moving lens fluctuations. We conclude that {\\it the contribution of the secondary anisotropies due to the moving lenses is thus negligible whatever the cosmological model. Therefore it will not affect the future CMB measurements except as a background contribution. We have highlighted the fact that the moving lens fluctuations have a very significant spatial signature but we did not succeed in separating this contribution from the other signals.}"
    },
    "9803/astro-ph9803276_arXiv.txt": {
        "abstract": "Images recorded through broad ($J, H, K$), and narrow (CO, and $2.2\\mu$m continuum) band filters are used to investigate the photometric properties of bright ($K \\leq 13.5$) stars in a $6 \\times 6$ arcmin field centered on the SgrA complex. The giant branch ridgelines in the $(K, J-K)$ and $(K, H-K)$ color-magnitude diagrams are well matched by the Baade's Window (BW) M giant sequence if the mean extinction is A$_K \\sim 2.8$ mag. Extinction measurements for individual stars are estimated using the M$_K$ versus infrared color relations defined by M giants in BW, and the majority of stars have A$_K$ between 2.0 and 3.5 mag. The extinction is locally high in the SgrA complex, where A$_K \\sim 3.1$ mag. Reddening-corrected CO indices, CO$_o$, are derived for over 1300 stars with $J, H$, and $K$ brightnesses, and over 5300 stars with $H$ and $K$ brightnesses. The distribution of CO$_o$ values for stars with $K_o$ between 11.25 and 7.25 can be reproduced using the M$_K -$ CO$_o$ relation defined by M giants in BW. The data thus suggest that the most metal-rich giants in the central regions of the bulge and in BW have similar photometric properties and $2.3\\mu$m CO strengths. Hence, it appears that the central region of the bulge does not contain a population of stars that are significantly more metal-rich than what is seen in BW.  \\vspace{0.3cm} \\noindent{Key words: Galaxy: center -- stars: abundances -- stars: late-type} ",
        "introduction": "The stars in the central regions of the Galaxy have been the target of numerous photometric and spectroscopic investigations. The brightest, best-studied, objects belong to a young population (e.g. reviews by Blum, Sellgren, \\& DePoy 1996a, and Morris \\& Serabyn 1996) that dominates the near-infrared light within 0.1 -- 0.2 parsec ($\\sim 2.5 - 5.0$ arcsec) of SgrA* (Saha, Bicknell, \\& McGregor 1996), and is concentrated within the central $\\sim 1$ parsec (Allen 1994). However, there is also a large population of older stars near the Galactic Center (GC) belonging to the inner bulge, and it is only recently that efforts have been made to probe the nature of these objects.  Minniti {\\it et al.} (1995) reviewed metallicity estimates for a number of bulge fields, and a least squares fit to these data indicates that $\\Delta$[Fe/H]/$\\Delta$log(r) $= -1.5 \\pm 0.4$ (Davidge 1997). The presence of a metallicity gradient suggests that the material from which bulge stars formed experienced dissipation, so it might be anticipated that the central regions of the bulge will contain the most metal-rich stars in the Galaxy. While the fields considered by Minniti {\\it et al.} are at relatively large distances from the GC, and hence do not monitor trends at small radii, observations of other galaxies indicate that population gradients can extend into the central regions of bulges. This is clearly evident in spectroscopic studies of M31 (e.g. Davidge 1997), a galaxy that shares some morphological similarities with the Milky-Way (e.g. Blanco \\& Terndrup 1990).  It is not clear from the existing observational data if the bright stellar contents in the inner bulge and Baade's Window (BW), a field that is dominated by stars roughly 0.5 kpc from the GC, are similar. The $K$ luminosity function (LF) of moderately faint stars within an arcmin of SgrA* has a power-law exponent similar to that seen in BW (Blum {\\it et al.} 1996a; Davidge {\\it et al.} 1997). However, the significance of this result is low, as the LFs of bright giant branch stars are insensitive to metallicity (e.g. Bergbusch \\& VandenBerg 1992). The brightest inner bulge stars, which are evolving on the asymptotic giant branch (AGB), have $2\\mu$m spectroscopic properties reminiscent of bright giants in BW, although detailed measurements reveal that for a given equivalent width of near-infrared Na and Ca absorption, the inner bulge stars have deeper CO bands than giants in BW (Blum, Sellgren, \\& DePoy 1996b). While suggestive of differences in chemical composition, it should be recalled that these objects are the brightest, most highly evolved members of the bulge, and Na can be affected by mixing (e.g. Kraft 1994). Hence, the spectroscopic properties of these bright red giants may not be representative of fainter objects.  As the strongest features in the near-infrared spectra of cool evolved stars, the $2.3\\mu$m first-overtone CO bands provide an important means of probing the stellar content of the inner bulge. Unfortunately, the crowded nature of the inner bulge, coupled with the low multiplex advantage offered by the current generation of cryogenically-cooled spectrographs, most of which use a single long slit, makes a $2\\mu$m spectroscopic survey of a large sample of moderately faint objects a difficult task at present. Narrow-band imaging, using filters such as those described by Frogel {\\it et al.} (1978), provides a highly efficient alternate means of measuring the strength of CO absorption in a large number of objects. In the current paper $J, H, K$, CO and $2.2\\mu$m continuum measurements of moderately faint ($K \\leq 13.5$) stars are used to measure the strength of CO absorption in stars within 3 arcmin of the GC. To the best of our knowledge, this is the largest survey of stellar content in the central regions of the bulge conducted to date. The observations and reduction techniques are described in \\S 2, while the photometric measurements are discussed in \\S 3. In \\S 4 the line-of-sight extinction to these sources is estimated by assuming that they follow the same M$_K -$ color relations as M giants in BW. Reddening-corrected CO indices are derived, and the distribution of CO indices is compared with that predicted if stars in the inner bulge and BW follow similar M$_K -$ CO relations. A summary and brief discussion of the results follows in \\S 5. ",
        "conclusions": "Moderately deep near-infrared images have been used to probe the bright ($K \\leq 13.5$) stellar content in the central $6 \\times 6$ arcmin field of the Galaxy. The ridgeline of the bulge giant branch on the $(K, J-K)$ and $(K, H-K)$ CMDs is well matched by the BW M giant sequence, reddened according to the Rieke \\& Lebofsky (1985) extinction curve. This similarity in photometric properties suggests that the extinctions to individual stars in the inner bulge can be estimated by adopting the M$_K -$ color relations defined by M giants in BW. The mean extinction outside of the SgrA complex, where A$_K = 3.1$ mag, is A$_K = 2.8$ mag. The extinction estimates for individual stars have been used to generate reddening-corrected CO indices, and the histogram distribution of CO$_o$ values can be reproduced using the M$_K -$ CO$_o$ relation defined by M giants in BW. Therefore, M giants near the GC and in BW have similar $2.3\\mu$m CO strengths.  A potential source of systematic error in the procedure used to compute A$_K$ is that a single set of M$_K -$ color relations have been used, and no attempt has been made to allow for a dispersion in the metallicities of giants in the central regions of the bulge. There is a selection effect for magnitude-limited samples, in the sense that the brightest stars in a field containing an old composite population are likely the most metal-rich, so the current data likely do not sample the full range of metallicities near the GC. In any event, the CO$_o$ distribution does not change substantially when M$_K -$ color relations defined by red giants in 47 Tuc ([Fe/H] $\\sim -0.7$) are used to estimate extinction, indicating that the main conclusions of this paper are insensitive to the adopted intrinsic colors of GC giants.  Another source of systematic error is that the A$_K$ values derived in \\S 4 require measurements in $J, H$, and $K$. This broad wavelength coverage introduces a bias against heavily reddened objects, which are faint, and hence may not be detected, in $J$. In fact, there are regions where the extinction is so high that stars are not detected in $K$. The tendency to miss the most heavily reddened stars skews the A$_K$ distribution to lower values. One way to reduce, but not entirely remove, this bias is to estimate extinction using only $H-K$ colors, for which the number of objects is over $3 \\times$ greater than those with $J$ measurements. Following the procedure described earlier, $(H-K)_0$ was assigned to each star using the $M_K - (H-K)$ relation for BW M giants, and the results are shown in the lower panels of Figures 7 -- 12, while the radial distribution of CO$_o$ values is plotted in Figure 14. It is evident from the lower panel of Figure 8 that a number of sources with relatively high A$_K$ are added to the sample when measurements in $J$ are not required. Nevertheless, the impact on the CO$_o$ distribution is negligible, indicating that stars with higher than average obscuration near the GC have $2.3\\mu$m CO strengths that are similar to less heavily reddened stars.  It should be emphasized that the $2.3\\mu$m CO bands provide only one diagnostic of chemical composition. Indeed, studies of the radial behaviour of the $2.3\\mu$m CO bands in the bulges of other galaxies reveal weak or non-existant gradients (e.g. Frogel {\\it et al.} 1978), even though many of these systems show line strength gradients at optical wavelengths. When interpreting this ostensibly contradictory result it should be recalled that the CO measurements are dominated by the brightest red stars which, in old populations, will be those that are the most metal-rich. If the most metal-rich population in the bulges of other galaxies follows a single M$_K -$ CO relation with no radial dependence, as appears to be the case in the inner regions of the Galactic bulge, then this would help to explain why the wide-aperture CO measurements of other systems do not show gradients.  There are indications that the abundances of some species in the spectra of metal-rich giants may vary with distance from the GC. In particular, the infrared spectra obtained by Blum {\\it et al.} (1996b) indicate that $2.3\\mu$m CO absorption in GC giants is {\\it stronger} than in BW stars having the same $2\\mu$m Na and Ca absorption line strengths. Therefore, given that the CO line strengths in giants near the GC are similar to those in BW, then the Blum et al. measurements are suggestive of radial changes in [Na/Fe] and [Ca/Fe] among bulge stars.  While this paper has concentrated on bulge stars, a modest disk population is also evident in the CMDs, and these data can be used to estimate the contribution disk objects make to the near-infrared light output near the GC. Objects with $(J-K) \\leq 2.5$ and brightnesses between $K \\sim 7.5$ (the approximate saturation limit of the CTIO data) and $K \\sim 12$ (the faintest point at which the CTIO $K, J-K$ CMD is complete over a broad range of colors) account for 2.6\\% of the total light from resolved sources within 3 arcmin of the GC. The relative contribution made by disk stars would be much lower if dust did not obscure the bulge. Assuming that (1) the disk stars are not heavily reddened, and (2) stars near the GC are obscured by A$_K \\sim 2.8$ mag then, after correcting for this extinction, the contribution from disk stars drops to only 0.05\\% when $K_0 \\leq 8.5$.  \\vspace{0.3cm} Sincere thanks are extended to the referee, Jay Frogel, for providing comments that greatly improved the paper.   \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{cccr} \\tableline\\tableline IRS \\# & K$_{TD}$ & K$_{BSD}$ & $\\Delta$ \\\\ \\tableline 7 & 6.44 & 6.55 & $-0.11$ \\\\ 9 & 8.80 & 8.57 & 0.23 \\\\ 16NE & 8.56 & 9.01 & $-0.45$ \\\\ 28 & 9.41 & 9,36 & 0.05 \\\\ \\tableline \\end{tabular} \\end{center} \\caption{Comparison with $K$ measurements obtained by Blum {\\it et al.} (1996a)} \\end{table*}  \\clearpage"
    },
    "9803/astro-ph9803106_arXiv.txt": {
        "abstract": "Lithium is an excellent tracer of mixing in stars as it is destroyed (by nuclear reactions) at a temperature around $\\sim 2.5\\times 10^6$ K.  The lithium destruction zone is typically located in the radiative region of a star.  If the radiative regions are stable, the observed surface value of lithium should remain constant with time.  However, comparison of the meteoritic and photospheric Li abundances in the Sun indicate that the surface abundance of Li in the Sun has been depleted by more than two orders of magnitude.  This is not predicted by solar models and is a long standing problem.  Observations of Li in open clusters indicate that Li depletion is occurring on the main sequence. Furthermore, there is now compelling observational evidence that a spread of lithium abundances is present in nearly identical stars. This suggests that some transport process is occurring in stellar radiative regions.  Helioseismic inversions support this conclusion, for they suggest that standard solar models need to be modified below the base of the convection zone. There are a number of possible theoretical explanations for this transport process. The relation between Li abundances, rotation rates and the presence of a tidally locked companion along with the observed internal rotation in the Sun indicate that the mixing is most likely induced by rotation.  The current status of non-standard (particularly rotational) stellar models which attempt to account for the lithium observations are reviewed. ",
        "introduction": "Li\\footnote{In this review I will use Li to represent $^7$Li, the isotope which is produced by big bang nucleosynthesis.  $^7$Li accounts for $\\sim 93\\%$ of the total Li abundance in meteorites. Observers typically measure the total Li abundance, while theoretical models determine depletion factors for $^7$Li and $^6$Li The $^6$Li isotope is destroyed at much lower temperatures and $^7$Li.  When making the comparison between the observations and theory, it is usually assumed that the $^6$Li contribution to the stellar Li content is neglible. However, see the discussion on halo stars (\\S \\ref{secthalo}) where observations of $^6$Li may be used to elucidate the mixing mechanism operating in these stars.}  is a sensitive tracer of mixing in stellar radiative regions as it is easily destroyed at temperatures above $\\sim 2.5 \\times 10^6\\,$K.  For solar type stars, the Li destruction region is located below the surface convection zone in standard models.  As a consequence, standard stellar models predict that Li should not be depleted at the surface of solar type stars.  This is a rather robust prediction of stellar evolution theory, which has been known for 40 years \\cite{schwar}. However, comparisons between the solar photospheric Li abundance and the Li abundance in meteorites show that the Sun has depleted a substantial amount of Li at its surface \\cite{green51}. The solar Li depletion problem has posed a challenge to stellar evolution theory for 40 years, and the solution to this puzzle is still open to debate.  The Sun is unique in that helioseismic observations allow us to probe the interior structure and rotation of the Sun.  These observations can put constraints on possible solutions to the solar Li depletion problem, but by themselves solar observations cannot uniquely determine the cause of solar Li depletion.  Observations of stellar Li abundances allow one to study the Li depletion problem as a function of age, metallicity and stellar mass.  As such, they provide a powerful test for mechanisms which attempt to explain the solar Li depletion. The discovery of a large dip in Li abundances around 6600 K in the Hyades \\cite{lidip} was not predicted by theorists, and remains a major challenge to theoretical stellar evolution models.  There is increasing evidence that a dispersion in Li abundances exists among stars with similar ages, metallicities and masses (\\citeauthor{soder93} \\citeyear{soder93}; \\citeauthor{boes98} \\citeyear{boes98}).  Such a dispersion suggests that another stellar property is important in determining the amount of Li which is depleted in stars.  There is mounting observational evidence that rotation plays a key role in determining the amount of Li depletion in a star (\\citeauthor{hyadbin} \\citeyear{hyadbin}; \\citeauthor{jones97} \\citeyear{jones97}). In this review, I will discuss the relationship between mixing, rotation and Li abundances in stars. ",
        "conclusions": "} There is a wealth of data on Li abundances and rotation velocities in stars with a variety of ages, masses and metallicities.  This data clearly indicates that Li depletion occurs on the main sequence for all stars with a surface convection zone.  This is in direct contradiction with standard stellar evolution theory. A number of possible mechanisms which lead to extra Li depletion have been put forth. The dispersion in Li abundances at a given age, metallicity and temperature, the correlation between Li abundances and rotation velocities in Pleiades, the fact that tidally locked binary stars in the Hyades have an excess Li abundance as compared to single stars, and the detection of moderate Be deficiencies among Li dip stars with detectable Li abundances, all imply that that rotation induced mixing is leading to Li depletion on the main sequence.  Helioseismic observations of the Sun support this hypothesis, for they show that slow form of slow mixing is operating below the base of the solar convection zone \\cite{basu}. Current stellar models which incorporate rotation induced mixing explain many, by not all of the observations.  Models which are able to account for all of the data are likely to include diffusion, rotation induced mixing, angular momentum transport by gravity waves and/or magnetic fields and modest stellar winds."
    },
    "9803/astro-ph9803330_arXiv.txt": {
        "abstract": "We have used a coupled time-dependent chemical and dynamical model to investigate the lifetime of the chemical legacy left in the wake of C-type shocks. We concentrate this study on the chemistry of \\HtwoO\\ and \\Otwo , two molecules which are predicted to have abundances that are significantly affected in shock-heated gas. Two models are presented: (1) a three-stage model of pre-shock, shocked, and post-shock gas; and (2) a Monte-Carlo cloud simulation where we explore the effects of stochastic shock activity on molecular gas over a cloud lifetime.   For both models we separately examine the pure gas-phase chemistry as well as the chemistry including the interactions of molecules with grain surfaces.  In agreement with previous studies, we find that shock velocities in excess of 10 km s$^{-1}$ are required to convert all of the oxygen not locked in CO into \\HtwoO\\ before the gas has an opportunity to cool.   For pure gas-phase models the lifetime of the high water abundances, or ``\\HtwoO\\ legacy'', in the post-shock gas is $\\sim 4 - 7 \\times 10^{5}$ years, independent of the gas density.  A density dependence for the lifetime of \\HtwoO\\ is found in gas-grain models as the water molecules deplete onto grains at the depletion timescale.  Through the Monte Carlo cloud simulation we demonstrate that the time-average abundance of \\HtwoO\\ -- the weighted average of the amount of time gas spends in pre-shock, shock, and post-shock stages -- is a sensitive function of the frequency of shocks. Thus we predict that the abundance of \\HtwoO , and to a lesser extent \\Otwo , can be used to trace the history of shock activity in molecular gas.  We use previous large-scale surveys of molecular outflows to constrain the frequency of 10 km s$^{-1}$ shocks in regions with varying star-formation properties and discuss the observations required to test these results. We discuss the post-shock lifetimes for other possible outflow tracers (e.g. SiO, \\CHthreeOH) and show that the differences between the lifetimes for various tracers can produce potentially observable chemical variations between younger and older outflows. For gas-grain models we find that the abundance of water-ice on grain surfaces can be quite large and is comparable to that observed in molecular clouds.  This offers a possible alternative method  to create water mantles without resorting to grain surface chemistry: gas heating and chemical modification due to a C-type shock and subsequent depletion of the gas-phase species onto grain mantles. ",
        "introduction": "The importance of molecular gas as the material from which stellar and planetary systems are made has led numerous investigators to examine the chemical processes that combine atoms into molecules.  Over the past two decades these studies have gradually increased in both complexity and fidelity and can roughly be divided into three generations.  The first generation chemical model solved the chemical rate equations at equilibrium, or steady-state, and demonstrated the importance of ion-molecule reactions in driving the gas-phase chemistry (\\cite{HK73}).  The second generation of models utilized the later availability of increased computing power to study the time evolution of chemical abundances (\\cite{PH80}; \\cite{GLF82}). These models, labeled as ``pseudo-time dependent'' -- ``pseudo'' because the chemistry evolves with fixed physical conditions -- used observed physical properties of dense molecular cores ($T_k \\sim 10 - 30$ K, \\nhtwo\\ $= 10^{4-6}$ \\cc ) to show that chemical equilibrium was reached in $\\sim 10^{6-8}$ years. Pseudo-time dependent models represent the most prevalent chemical model and have been quite successful in describing the chemistry of quiescent regions in both dark and giant molecular cloud cores (\\cite{Lee_etal96}; \\cite{BGSL97}).  Second generation chemical models operate under one simplifying assumption: that the molecular gas undergoes no dynamical evolution as it chemically evolves. However, molecular clouds are certainly dynamically evolving objects. The widespread occurrence of superthermal widths in molecular lines suggests that, overall, the evolution of molecular clouds is not simple quiescent evolution at a single gas temperature.  Moreover, the formation of a low- or high-mass star from a molecular condensation involves a collapse that increases the density by many orders of magnitude.  It has also been recognized that the birth of a protostar is associated with a period of intense mass loss, which manifests itself in energetic winds, bipolar jets, and large-scale outflows (c.f. \\cite{Lada85}; \\cite{Bachiller96}).  The impact of energetic flows on surrounding quiescent gas can compress and heat the gas, in some cases providing enough energy to overcome endothermic barriers of chemical reactions, sputter or destroy grains, or even dissociate molecules.  The result can be a considerably different chemical composition in the shocked gas than observed in quiescent material. Since molecules are the primary coolants of the gas, any chemical change induced by collapse or shocks can also alter the ensuing physical evolution of the cloud. To address such concerns, a third generation of chemical models has been constructed to build on the successes of previous generations by combining chemistry with dynamics.  These models have been principally directed towards investigating the chemistry of core and star formation (c.f. \\cite{PTVH87}; \\cite{RHMW92}; \\cite{BL97}), the physical and chemical structure of shocked gas (c.f. \\cite{DM93}), or even complex scenarios in which gas is cycled between low and high density through recurrent episodes of low-mass star formation (\\cite{CDHW88a}).   One goal of coupled astrochemical models is to search for and isolate specific molecular species that serve as signposts of particular dynamical events, such as shocks. Coupled models of shocked molecular gas have isolated one molecule in particular, \\HtwoO , which is predicted to form in large abundance in shock-heated gas (\\cite{DRD83}; \\cite{KN96a},b); if temperatures in the shocked gas exceed 400~K, the endothermic barriers of a few key chemical reactions are exceeded and all of the available oxygen not locked up in CO will be driven into \\HtwoO . This led to the assertion that water emission is a ubiquitous tracer of shock-heated gas (\\cite{NM87}).  Unfortunately, due to constraints imposed by the earth's atmosphere, the detection of \\HtwoO\\ in interstellar clouds from ground-based observatories is a challenging task. Nevertheless, several studies have observed, and detected, transitions of isotopic water (\\HtwoeiO), and even \\HtwoO\\ (c.f. \\cite{Jacq_etal88}; \\cite{KL91}; \\cite{Zmuidzinas_etal94}). Quite recently, absorption from warm ($T_{gas} > 200$ K) water vapor has been unambiguously detected by the Infrared Space Observatory (ISO) towards hot stars (\\cite{Helmich_etal96}; \\cite{vDH96}), and in emission in HH54 (\\cite{Liseau_etal96}).  The inferred water abundance in these sources is $\\sim 1 - 6 \\times 10^{-5}$, much larger than predicted by chemical models run for cold quiescent conditions ($x$(\\HtwoO ) $\\sim$10$^{-7}$), and is consistent with water production in warm gas either though high-temperature chemistry -- as would apply in shock-heated gas or in the near vicinity of embedded sources -- or via evaporation of water-ice mantles.  An important question yet to be addressed by coupled dynamical and chemical models of shocked gas in molecular clouds is how long the high abundances of water and other molecules persist following the passage of a shock.   As we will demonstrate, the time needed for shock-heated gas to return to its pre-shock temperature is several orders of magnitude less than the time required for the gas to return to its pre-shock chemical composition.  Thus, the enhanced abundances of water and other species could potentially exist long after the dynamical effects of a shock passage are dissipated. Therefore, the chemistry inside a cloud is reflective of, and can be used to probe, the physical shock history of the molecular gas.  Motivated by these questions, we present the results of a coupled dynamical and chemical study of molecular gas that is subjected to shocks with velocities greater than 10 km s$^{-1}$. In particular, we follow the time dependence of the chemistry in a shock-heated gas layer as it cools and the quiescent time-dependent chemistry is ultimately re-established. We present models examining this evolution using pure gas-phase chemistry as well as chemistry which includes the interaction of molecules with grain surfaces.  In Section 2 we discuss the combined chemical and dynamical model. In Section 3 we use published surveys of molecular outflows to investigate the average rate at which shocks with a minimally sufficient velocity to affect the water chemistry -- 10 km s$^{-1}$ -- pass a given region of a cloud. Section 4 presents the results from our combined models along with one example of quiescent chemical evolution.  Two models are presented: (1) a three-stage model of pre-shock, shocked, and post-shock gas; and, (2) a Monte-Carlo cloud simulation in which the results from Section 3 are used to examine the effects of stochastic shock activity on molecular gas over a cloud lifetime. Section 5 discusses the importance of these results on chemical models of molecular gas and also reviews the observations required to test these models.   In Section 6 we summarize our results. ",
        "conclusions": "\\subsection{Comparison with Observations}  Our results demonstrate that the {\\em time-averaged} abundance of \\HtwoO\\ is a sensitive function of the rate at which molecular gas is subjected to shocks. In order to properly interpret these results it is useful to consider how the time-averaged abundances relate to observations of water in a giant molecular cloud (GMC).   The gas inside a star-forming cloud can roughly be divided into three categories:  (1) quiescent material that has been relatively unaffected by current or previous epochs of star formation; (2) gas that is being physically and/or chemically affected by current star formation; and, (3) as suggested here, gas that has been affected by a previous generation of star formation and is in the process of chemically and dynamically evolving back to a quiescent state.  Observations of \\HtwoO\\ toward molecular material associated with each of these categories should find abundances varying from as low as $x(\\rm{H_2O}) \\sim 10^{-6}$ in (1), and ranging upwards to $\\le 10^{-4}$ in (2) and (3).   {\\em A single pointed detection of H$_2$O emission in a GMC can therefore be represented by a single stage of the three-stage or Monte-Carlo models (quiescent, shock, post-shock).  Because star formation will be spread throughout a cloud or core, the time-averaged abundance therefore refers to water abundances averaged over an area that contains both quiescent material and any associated gas currently being affected by local star formation. Thus the average abundance over a cloud (cloud-average) is comparable to the time-averaged abundances in the Monte-Carlo cloud model.}  This average abundance might not necessarily apply to an entire GMC complex (e.g. Orion, Gem OB1), which can extend for several square degrees on the sky, but may apply to several dense cores within a single cloud. Regions with high star-formation rates, such as OMC-1 which is associated with the Trapezium cluster (see Section 3), might be expected to have a higher probability for strong shocks,  and therefore a high cloud-averaged water abundance.  Cores with lower star-formation rates, such as L1641 also in Orion, would have low cloud-averaged water abundances.  {\\em To test our model's ability to constrain the history of molecular clouds it is important to obtain maps of the water emission over large spatial scales in GMC cores.} Because of the strong absorption by atmospheric water vapor, observations of \\HtwoO\\ in the ISM are extremely difficult. As a result, the small number of detections of water in molecular clouds have typically been taken towards single lines of sight mostly containing luminous protostars. The very convincing detections of water by ISO provide the greatest evidence for high water abundances in hot gas, with $x(\\rm{H_2O}) \\sim 1 - 6 \\times 10^{-5}$ inferred towards hot stars (\\cite{Helmich_etal96}; \\cite{vDH96}) and HH54 (\\cite{Liseau_etal96}).  Towards Sgr B2 Zmuidzinas et al. (1994) observed \\HtwoeiO\\ in absorption with the Kuiper Airbourne Observatory (KAO), while Neufeld et al. (1997) observed \\HtwoO\\ in emission using ISO. Combining these observations with previous ground based detections Neufeld et al. (1997) estimate an abundance of  $x(\\rm{H_2O}) = 3.3 \\times 10^{-7}$ for the cooler outer parts of Sgr B2 and  $x(\\rm{H_2O}) \\sim 5 \\times 10^{-6}$ in the hot core. These observations are important not only because \\HtwoO\\ was detected in dense gas, but also because there appears to be a range of water abundances.  However, it is difficult to discern whether the enhanced water abundances are the result of high-temperature chemistry that occurs behind shocks, high-temperature chemistry appropriate to gas near young stars (c.f. Doty \\& Neufeld 1997), or evaporation of grain mantle species (see discussion in Section 5.2) and there is little information on the spatial distribution of the water emission.  There have been a few attempts to map the extended emission of water. Gensheimer, Mauersberger, \\& Wilson (1996) mapped emission of the quasi-thermal $3_{13} \\rightarrow 2_{20}$ transition of \\HtwoeiO\\ in both the Orion Hot Core and Sgr B2, and found that the emission originated in a compact region ($< 10''$).  Cernicharo et al. (1994) and Gonzalez-Alfonso et al. (1995) find evidence for widespread water emission in Orion and W49N using the $3_{13} \\rightarrow 2_{20}$ masing transition of \\HtwoO .   In Orion they argue that the water abundance is quite high, $x(\\rm{H_2O}) > 10^{-5}$, over an extended $50'' \\times 50''$ region centered on BN-KL and the Orion Nebular Cluster. From these results, the cloud-averaged water abundance for the central regions of the Orion core near the Orion Nebular Cluster is $^{>}_{\\sim} 10^{-5}$ which, using Figures 11 and 12 (\\nhtwo\\ $= 10^6$ \\cc ; \\cite{BSG96}), is consistent with $\\tau_s \\sim 10^6$ years, as suggested in Section 3. However, the determination of abundances from masing transitions is a difficult task and these results must be viewed as suggestive until they are confirmed by mapping data obtained in other \\HtwoO\\ lines.  For molecular oxygen, which also suffers from strong absorption due to atmospheric \\Otwo , the situation is even more uncertain.  There exists only one tentative detection of $^{16}$O$^{18}$O towards L134N by \\cite{PLC93}, which implies $x(\\rm{O_2}) \\sim 4 - 8 \\times 10^{-5}$.  However, a search for $^{16}$O$^{18}$O emission in different positions in L134N and other galactic sources by \\cite{MPLC97} did not confirm this detection and provides only upper limits of \\Otwo /CO $< 0.1$, which, assuming $x$(CO) $= 2.7 \\times 10^{-4}$ (Lacy et al 1994), gives $x(\\rm{O_2}) < 3 \\times 10^{-5}$.  Searches, with similar negative results, have also been performed for molecular oxygen in extra-galactic sources with favorable redshifts (Goldsmith \\& Young 1989; Combes et al. 1991; Liszt 1992) finding $x(\\rm{O_2}) <  10^{-6}$. A recent study by Combes, Wiklind, \\& Nakai (1997) towards a z = 0.685 object provides the lowest limit to date of \\Otwo /CO $< 2 \\times 10^{-3}$. For our most realistic case, the Monte Carlo gas-grain model, we find the time-averaged \\Otwo\\ abundance is $\\sim 2 \\times 10^{-5}$.  Thus, the combined effects of shocks and gas-grain chemistry could lower \\Otwo\\ abundances below the observed limits in Galactic sources, but another explanation is required to account for extra-galactic observations.  The low time-averaged \\Otwo\\ abundance suggests that these results could have some bearing on the high abundances of neutral carbon relative to CO (C/CO $\\sim$ 0.1) that are observed towards a variety of star forming regions (c.f. Plume 1995, Schilke et al 1996).  However, because the number of dissociative shocks which destroy CO is not high enough, the time-averaged carbon abundance in the models is well below the observed value.  The best opportunity to test these results will be offered by two spaceborne observatories, {\\em The Submillimeter Wave Astronomy Satellite (SWAS)} (\\cite{Melnick_etal95}) and {\\em ODIN} (\\cite{Hjalmarson95}), both of which should launch within the next year.   {\\em SWAS} and {\\em ODIN} are capable of observing and mapping the fundamental transition $1_{10} \\rightarrow 1_{01}$ of \\HtwoO\\ at 557 GHz and the $3,3 \\rightarrow 1,2$ transition of \\Otwo\\ at 487 GHz. These transitions have low upper state energies ($\\sim 26$ K) and should be readily excited in hot gas near star-forming sites and, more importantly, in the colder more extended material such as the ridge of dense gas in Orion (c.f. \\cite{Ungerechts_etal97}).  \\begin{figure*} \\figurenum{16} \\plotfiddle{fig16.ps}{2.5in}{0}{50}{50}{-160}{-75} \\caption{ Points defining the plane of O$_2$ and H$_2$O abundances for the gas-grain Monte-Carlo model with $\\tau_s = 10^{6}$ years and $n_{H_2} = 10^5$ cm$^{-3}$.  The spread of plotted points has been artificially increased by smearing them by 0.2 dex so as to give a better sense of the density of points in different regions of the plot.  The cross represents the mean (time-averaged) abundances. } \\label{shockspec} \\end{figure*} Besides the computation of averaged abundances from mapping observations, combined observations of \\HtwoO\\ and \\Otwo\\ should be a powerful tool in examining the shock history of gas. This is demonstrated in Figure 16 where the water abundance is plotted versus the molecular oxygen abundance for the Monte-Carlo model, including grain depletion and desorption (Section 4.4.2).  In this plot, the majority of points trace an area of roughly constant water abundance of $x(\\rm{H_2O}) \\sim 10^{-6}$, with the \\Otwo\\ abundance ranging between $x(\\rm{O_2}) = 3 - 30 \\times 10^{-6}$. This area has, by far, the largest number of points and defines the ``main-sequence'' of quiescent chemistry.  From this main sequence a shock will trace a line extending almost horizontally to the right, with constant \\Otwo\\ abundance.  Very strong shocks continue at the end of the horizontal line and extend down almost vertically, lowering the \\Otwo\\ abundance for constant \\HtwoO\\ abundance. This dependence of horizontal and vertical lines occurs because the \\HtwoO\\ is created more efficiently than \\Otwo\\ destruction (see Figures 3 and 4). This plot is a different way of examining the models, as opposed to average abundances, because it presents the evolution in a continuous fashion; the various possible solutions define a plane that traces evolutionary tracks including quiescent chemistry through a broad range of shock strengths.   \\subsection{Post-Shock Chemistry and Water Ice Mantles}  For comparison with previous modeling efforts, our three-stage model is similar to the episodic models presented in Charnley et al. (1988a,b).  However, our model is simpler in that we do not model the formation process of a dense clump nor do we account for any mixing between shocked and non-shocked layers.  Like their models, we find that the chemistry converges to well defined abundance values even when numerous cycles are repeated. Charnley et al. (1988a) also find that the water abundance varies greatly between shock and quiescent cycles, but they do not examine in detail the chemistry in the post-shock layer.  An interesting result in the gas-grain models is that the abundance of water on grain surfaces in post-shock gas is quite large $x(\\rm{H_2O})_{gr} \\sim 10^{-4}$. This abundance is quite close to that inferred for water ice along lines of sight towards background stars in Taurus, where $x(\\rm{H_2O})_{gr} \\sim 8 \\times 10^{-5}$ (\\cite{Whittet93}).   Thus, these results offer an alternative explanation for the large abundance of water in grain mantles -- one that requires no grain surface chemistry! We stress that this mechanism may not be responsible for all \\HtwoO\\ observed on grains.  Grain surface chemistry formation of \\HtwoO\\ must be considered -- especially during cloud formation stages when the abundance of atomic hydrogen and oxygen is higher. There are also other potential desorption mechanisms that could be active and are not included in these models, such as grain mantle explosions induced by UV radiation (\\cite{SG91}) or water desorption via the infrared radiation field (\\cite{WHW92}). However, since most molecular clouds are very active star-forming sites, we suggest that caution must be applied when interpreting observations of ice mantles in molecular clouds as the sole result of grain surface chemistry. The interpretation of high gas-phase water abundances towards hot stars as the result of either high-temperature chemistry or grain mantle evaporation of \\HtwoO\\ is further obscured because water on grains may have been produced in an earlier shock episode.  It is possible that the HDO/H$_2$O ratio could discriminate between water mantles created in shocks or in grain mantles and  we are in the process of investigating this question (Bergin, Neufeld, \\& Melnick 1997).  The large abundance of water on grains, with CO remaining in the gas phase, will also alter the ratio of total carbon-to-oxygen (C/O) in the gas phase.  Comparison of chemical theory with observations of molecular abundances in GMC cores associated with massive star formation has found that the chemical abundances of many species are best reproduced with C/O ratios greater than the solar value ($>$ 0.4; \\cite{BSMP87}; \\cite{BGSL97}). For chemical modeling this is the result of depletion of the initial abundance of oxygen relative to carbon, which reduces the abundances of small oxygen-bearing species that are major carbon destroyers in the gas phase.    In the three-stage gas-grain model shown in Figure 7, the C/O ratio in the post-shock gas at $t \\sim 10^{5}$ years is C/O $\\sim 0.7$.  It is difficult to gauge the effect of this on other species, because of the lack of certainty with regard to high-temperature reactions. It is therefore possible that the C/O ratios inferred in Blake et al. (1987) and Bergin et al. (1997) are indicative of the core formation process, which could involve gas temperatures rising high enough to convert atomic oxygen to water. In this case, these results suggest a simple mechanism to place large amounts of oxygen on grain surfaces while still keeping most of the carbon in the gas-phase.  Other molecular species, such as SO, SiO (\\cite{MPBF92}), and \\CHthreeOH\\ (\\cite{BLWC95}) have been observed with enhanced abundances in energetic outflows (see also \\cite{Bachiller96}; \\cite{Bachiller_Perez97}). These species are included in the three-stage chemical model presented in Section 4.3, but we do not present the results in detail here because of uncertainties in the high-temperature reaction rates.  However, these results do have some bearing on the chemistry of these species in the post-shock layer.   The formation pathway of \\CHthreeOH\\ in the gas phase is linked to a reaction of CH$_3^+$ with \\HtwoO\\ (\\cite{MHC91}).  Thus,  it is possible that the enhanced abundance of water through high-temperature reactions could lead to larger \\CHthreeOH\\ abundances through this reaction.  In the three-stage model, the \\CHthreeOH\\ abundance in the post-shock layer is $x(\\rm{CH_3OH}) = 1 - 5 \\times 10^{-8}$ when the shock temperature ranges from 1000 to 2000 K. These abundances are at least an order of magnitude below values inferred in molecular outflows, which can be as high as $x(\\rm{CH_3OH}) \\sim 10^{-6}$ (\\cite{BLWC95}). The mechanism of methanol enhancements may therefore be the result of grain mantle evaporation or an unidentified high- or low-temperature pathway.  The abundance of \\CHthreeOH\\ accreted onto grain surfaces provides another constraint.  In our model the abundance of \\CHthreeOH\\ ice in the grain mantle is only 0.1 percent of the water-ice abundance. This ratio is below that observed towards NGC 7538 or W33A where $x(\\rm{CH_3OH})$/$x(\\rm{H_2O}) \\sim 10-40$ percent (\\cite{ASTH92}).  Our models do not set constraints on the chemistry of SO and SiO because the large abundances of these species in outflows may be the result of sputtering of grain refractory and/or mantle material (c.f. Schilke et al. 1997; Caselli et al. 1997).  When the abundances of these species are enhanced, through any mechanism, the high abundances will persist until the timescale for the individual molecule to deplete onto grain surfaces. If the dust temperature is high enough to keep a given species in the gas-phase, or for pure gas-phase chemistry, the lifetime will be $\\tau_{ra} = 4 -7 \\times 10^{5}$ years. In these models all SiO depletes onto grain surfaces at $\\tau_{dep} \\sim 10^4$ years (for \\nhtwo\\ = 10$^{5}$ \\cc\\ and assuming a sticking coefficient of unity).  For \\CHthreeOH\\ abundance enhancements we find that the depletion timescale for gas-grain chemistry is equal to that of \\HtwoO\\ ($\\tau_{dep}(H_2O) \\sim 10^5$ years). The disparity in post-shock lifetimes between \\CHthreeOH , \\HtwoO , and SiO suggests that differences might exist between younger outflows, perhaps those associated with so-called Class~0 sources, and the more evolved sources (e.g. Class I). A survey of such sources in these tracers might prove to be useful in probing the links between evolutionary effects observed in outflows and the relationship to the driving source. Another possibility is that differences could also be found within different components inside a single outflow  (see \\cite{Bachiller_Perez97}).  The abundance of some molecular species, notably \\HCOp\\ and \\NtwoHp , are adversely affected by the high water abundance.  In Figure 7, for the three-stage model, the abundance of \\HCOp\\ is decreased through reaction R6 when the water abundance is raised. Thus, the abundances of these two important molecular species should be anticorrelated. Because \\NtwoHp\\ also reacts with \\HtwoO , a similar anticorrelated behavior would be found  between the abundances of \\NtwoHp\\ and \\HtwoO\\ and may account for the low \\HCOp\\ and \\NtwoHp\\ abundances in hot regions such as the Orion Hot Core (c.f. Blake et al. 1987) and in the L1157 outflow (\\cite{Bachiller_Perez97})."
    },
    "9803/astro-ph9803324_arXiv.txt": {
        "abstract": "The IRAS 100 micron image of the GRB 970228 field shows that the amount of galactic dust in this direction is substantial and varies on arcminute angular scales.  From an analysis of the observed surface density of galaxies in the $2.6' \\times 2.6'$ HST WFPC image of the GRB 970228 field, we find $A_V = 1.1 \\pm 0.10$.  From an analysis of the observed spectra of three stars in the GRB 970228 field, we find $A_V = 1.71^{+0.20}_{-0.40}$. This value may represent the best estimate of the extinction in the direction of GRB 970228, since these three stars lie only $2.7''$, $16''$, and $42''$ away from the optical transient.  If instead we combine the two results, we obtain a conservative value $A_V = 1.3 \\pm 0.2$.  This value is significantly larger than the values $A_V = 0.4 - 0.8$ used in papers to date.  The value of $A_V$ that we find implies that, if the extended source in the burst error circle is extragalactic and therefore lies beyond the dust in our own galaxy, its optical spectrum is very blue: its observed color $(V-I_c)_{\\rm obs} \\approx 0.65^{+0.74}_{-0.94}$ is consistent only with a starburst galaxy, an irregular galaxy at $z > 1.5$, or a spiral galaxy at $z > 2$.  On the other hand, its observed color and surface brightness $\\mu_V \\approx 24.5$ arcsec$^{-2}$ are similar to those expected for the reflected light from a dust cloud in our own galaxy, if the cloud lies in front of most of the dust in this direction. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803262_arXiv.txt": {
        "abstract": "We present accurate measurements of the central wavelengths of 4947 atomic absorption lines in the solar optical spectrum. The wavelengths, precise to a level $\\sim 50-150$ m s$^{-1}$, are given for both flux and disc-centre spectra, as measured in relatively recent FTS solar atlases. This catalogue modernizes existing sources based on photographic measurements and provides a benchmark to test and perform wavelength calibrations of astronomical spectra. It will also permit observers to improve the absolute wavelength calibration of solar optical spectra when lamps are not available at the telescope. ",
        "introduction": "\\label{sec1}  Wavelength calibration is almost always needed in the process of producing useful astronomical spectra. To calibrate accurately is a non-trivial problem, in particular when working at high or very-high spectral resolution. Fourier transform spectrographs (FTS) are specially well-suited to this task, but they are not readily applied in conditions requiring high spatial or time resolution, so grating spectrometers are much more commonly used for astronomical observations. In this case, it becomes necessary to set reference positions corresponding to known wavelengths on the detector.  This can be achieved by using very sharp observed telluric lines, but their location in the spectrum cannot be chosen by the astronomer. It is very usual to find spectral calibration lamps available for use with an astronomical spectrograph. The emission lines produced in the lamps have been previously measured at the laboratory, and this method usually provides a valid reference frame. However, it is often impractical to expose the calibration lamp simultaneously with the astronomical target and, unless the spectrograph is installed at a very stable focal station, the position of the spectrum on the  detector varies depending on the telescope position. Accuracy is then limited by the instrument characteristics and observations of the calibration lamps are required between successive astronomical exposures. Nonetheless, calibrations via arc or hollow cathode spectra are normally accurate enough for most purposes.  Ingenious techniques have been used to improve the accuracy of wavelength calibrations, such as placing gas cells at the entrance of the spectrograph (e.g., Deming \\& Plymate 1994), but it is rare to find such systems available and convenient for regular observations.  On occasion the available lamps are not very rich in lines in the spectral range of observation. In some circumstances, an external check of the final precision in the translation into wavelengths would be desirable. One method for tackling problems such as these is to use solar spectra as templates.  Changes in the wavelengths of the lines in the  integrated sunlight spectrum around the solar cycle have been proved to be very small, bellow some 15 m s$^{-1}$ (Jim\\'enez et al. 1980; Wallace et al. 1988; McMillan et al. 1993; Deming \\& Plymate 1994).  At 5000 \\AA, this translates into $\\sim$ 0.3 m\\AA, so the solar spectrum does offer a  very stable source. In most practical cases, the accuracy will be imposed by the spectral resolution achieved. During night-time observations, the solar flux spectrum is observable after reflection from the Moon.  Measurements of solar wavelengths in the integrated solar optical spectrum were published in 1929 by Burns and collaborators (Burns 1929; Burns \\& Kiess 1929; Burns \\& Meggers 1929), using photographic detectors and a grating spectrograph. The relatively recent solar flux FTS atlases offer a much higher quality spectrum of the Sun seen as a star.  As the solar spectrum is so intense, on some solar telescopes no calibration lamps are deemed necessary, and the wavelength scale is set using the solar spectrum itself. Reasonable precision can be reached using the spectrum at the centre of the disc to compare with previously measured disc centre wavelengths, thus avoiding differential shifts due to the limb effect. In this case, small scale motions have to be averaged out, integrating in time and/or space, in order to minimize errors.  The {\\it Kitt Peak Table of Photographic Solar Spectrum Wavelengths} (Pierce \\& Breckinridge 1973) has been extensively used by solar observers to set up the wavelength scale on their spectra. These observations, made on photographic plates, have been superseded in quality by the more recent FTS observations at the centre of the disc.  To improve on the various sets of photographically based measurements (which date back to 1930 in the case of the solar flux spectrum), provide them in a homogeneous machine-readable format, use them to test spectral calibrations of very high resolution stellar spectra (e.g., Allende Prieto et al. 1995), and improve the accuracy of our own solar observations, we have determined the position of the central wavelengths of 4947 atomic lines in the optical solar spectrum. The employed source solar atlases, prepared from FTS data, and the fashion in which we performed the measurements is described in the succeeding sections. ",
        "conclusions": ""
    },
    "9803/astro-ph9803218_arXiv.txt": {
        "abstract": " ",
        "introduction": "The goal of modern cosmology is to find the large scale matter distribution and spacetime structure of the Universe from astronomical observations, and it dates back from the early days of observational cosmology the realization that in order to achieve this aim it is essential that an accurate empirical description of galaxy clustering be derived from the systematic observations of distant galaxies. As time has passed, this realization has become a program, which in the last decade or so took a great impulse forward due to improvements in astronomical data acquisition techniques and data analysis. As a result of that an enormous amount of data about the observable universe was accumulated in the form of the now well-known {\\it redshift surveys}, and some widely accepted conclusions drawn from these data created a certain confidence in many researchers that such an accurate description of the distribution of galaxies was just about to being achieved. However, those conclusions are mainly based in a standard statistical analysis derived from a scenario provided by the standard Friedmannian cosmological models, which assume homogeneity and isotropy of the matter distribution, scenario which is still thought by many to be the best theoretical framework capable of explaining the large scale matter distribution and spacetime structure of the Universe.  The view outlined above, which now has become the orthodox homogeneous universe view, has, however, never been able to fully overcome some of its objections. In particular, many researchers felt in the past, and others still feel today, that the relativistic derived idea of an eventual homogenization of the {\\it observed} matter distribution of the Universe is flawed, since, in their view, the empirical evidence collected from the systematic observation of distant cosmological sources also supports the claim that the universal distribution of matter will not eventually homogenize. Therefore, the critical voice claims that the large-scale distribution of matter in the Universe is intrinsically inhomogeneously distributed, from the smallest to the largest observed scales and, perhaps, indefinitively.  Despite this, it is a historical fact that the inhomogeneous view has never been as developed as the orthodox view, and perhaps the major cause for this situation was the lack of workable models supporting this inhomogeneous claim. There has been, however, one major exception, in the form of a hierarchical cosmological model advanced by Wertz (1970, 1971), although, for reasons that will be explained below, it has unfortunately remained largely ignored so far.  Nevertheless, by the mid 1980's those objections took a new vigour with the arrival of a new method for describing galaxy clustering based on ideas of a radically new geometrical perspective for the description of irregular patterns in nature: {\\it the fractal geometry}.  In this review we intend to show the basic ideas behind this new approach for the galaxy clustering problem. We will not present the orthodox traditional view since it can be easily found, for instance, in Peebles (1980, 1993) and Davis (1997). Therefore, we shall concentrate ourselves in the challenging voice based on a new viewpoint about the statistical characterization of galaxy clustering, whose results go against many traditional expectations, and which keep open the possibility that the universe never becomes observationally homogeneous. The basic papers where this fractal view for the distribution of galaxies can be found are relatively recent. Most of what will be presented here is based on Pietronero (1987), Pietronero, Montuori and Sylos Labini (1997), and on the comprehensive reviews by Coleman and Pietronero (1992), and Sylos Labini, Montuori and Pietronero (1998).  The plan of the paper is as follows. In section 2 we present a brief, but general, introduction to fractals, which emphasizes their empirical side and applications, but without neglecting their basic mathematical concepts. Section \\ref{Distr Galaxies} briefly presents the basic current analysis of the large scale distribution of galaxies, its difficulties and, finally, Pietronero-Wertz's single fractal (hierarchical) model that proposes an alternative point of view for describing and analysing this distribution, as well as some of the consequences of such an approach. The paper finishes with a discussion on some aspects of the current controversy about the fractal approach for describing the distribution of galaxies. ",
        "conclusions": ""
    },
    "9803/astro-ph9803133_arXiv.txt": {
        "abstract": "\\input abstract ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803305_arXiv.txt": {
        "abstract": "We present an $R$-band and $J$-band photometry of an optical transient which is likely to be associated with the gamma-ray burst event GRB\\,971214.  Our first measurement took place 13 hours after the gamma-ray event.  The brightness decayed with a power-law exponent $\\alpha = -1.20 \\pm 0.02$, which is similar to those of GRB\\,970228 and GRB\\,970508 which had exponents of $\\alpha = -1.10 \\pm 0.04$ and $\\alpha = -1.141\\pm 0.014$ respectively.  The transient decayed monotonically during the first four days following the gamma-ray event in contrast with the optical transient associated with GRB\\,970508 which increased in brightness, peaking two days after the burst, before settling to a power-law decay. ",
        "introduction": "The launch of the BeppoSAX satellite \\cite{boel97} in 1996 has led to a recent breakthrough in the study of gamma-ray bursts by detecting fading X-ray counterparts and localizing them to a few arc-minutes on the sky.  This has allowed subsequent identification of optical counterparts in three cases, GRB\\,970228 \\cite{groo97}, GRB\\,970508 \\cite{bond97}, and GRB\\,971214 \\cite{halp97}.  GRB\\,971214 triggered the BeppoSAX Gamma Ray Burst Monitor on December 14.97272 UT with a peak flux of 650 counts s$^{-1}$.  In addition, the burst was localized to an error radius of \\decmin{3}{9} (99\\% confidence) with the Wide Field Camera \\cite{heis97}. The gamma-ray event was also detected by BATSE which measured a total fluence above 20 keV of $1.09 \\pm 0.07 \\times 10^{-5}$ erg cm$^{-2}$ and by one RXTE-ASM camera yielding a peak intensity of $470 \\pm 140$ mCrab \\cite{kipp97}. Shortly thereafter, a fading optical source was detected within the BeppoSAX error circle at $\\alpha(J2000)={\\rm11^h56^m}$\\decsectim{26}{4}, $\\delta(J2000)={\\rm65^\\circ12'}$\\decsec{00}{5} with I-band magnitudes of $21.2 \\pm 0.3$ on Dec 15.47 UT and $\\sim 22.6$ on Dec 16.47 UT \\cite{halp97}. Further observations by BeppoSAX detected a previously unknown fading X-ray source, later designated 1SAX J1156.4+6513, within the initial error circle at $\\alpha(J2000)={\\rm11^h56^m}25^{\\rm s}$, $\\delta(J2000)={\\rm65^\\circ13'}11''$ with an error radius of about $1'$ \\cite{anto97}. Since this second X-ray detection is consistent with the position of the fading optical source identified by Halpern \\etal\\ (1997) it is quite likely that these objects are the X-ray and optical afterglow of GRB\\,971214.  We report here $R$-band and $J$-band observations of this optical transient (OT) made with the Apache Point Observatory(APO) 3.5 m telescope.  This work extends the preliminary photometry reported in Diercks \\etal\\ (1997), Castander \\etal\\ (1997), and Tanvir \\etal\\ (1997). ",
        "conclusions": "The well-observed light curve of the GRB\\,970508 OT, the brightest observed thus far, shows a dramatic rise, peaking nearly two days after the initial burst, before beginning a power-law decay. GRB\\,970228 was not observed nearly as often through a consistent filter, but there is also evidence (depending on spectral assumptions) \\cite{guar97} that the transient increased in brightness until $\\sim20$ hr after the burst, after which it also began fading with a power-law slope essentially identical to that of GRB\\,970508.  Despite the same decay slope, the GRB\\,970228 OT was $\\sim1.5$ mag fainter than the GRB\\,970508 OT.  A detailed analysis of the difference between these two light curves is presented in Pedersen et al. (1998).  A power-law of the form $F = F_{0}t^{\\alpha}$ was fit to the four $R$-band data points yielding $\\alpha = -1.20 \\pm 0.02$.  This rate of decline is similar to the two previously identified bursts although there is no evidence of the OT brightening over the course of the observations.  \\placefigure{fig-2}  The observations from all four nights are combined into one deep image totaling 3.25 hrs of integration with a limiting $R$-band magnitude $\\sim25.4$ (Figure~\\ref{fig-3}).  The two brightest objects within 20$''$ of the OT are an extended source (A) \\decsec{4}{6} southwest of the transient with $R = 22.7 \\pm 0.1$, and an unresolved source (B) \\decsec{4}{9} north of the transient with $R = 23.43 \\pm 0.08$.  The resolution of these images is insufficient to identify any structure in the extended object. \\footnote{To obtain the images discussed in this work, contact A. Diercks or E. Deutsch}"
    },
    "9803/astro-ph9803075_arXiv.txt": {
        "abstract": "We report on two observations of the Seyfert galaxy IRAS18325-5926 made in 1997 December and 1998 February with the Rossi X-ray Timing Explorer (RXTE). We find evidence for periodicities in the resulting X-ray lightcurves which are shorter than the 58~ks period found from data of the source taken in 1997 March with the imaging satellite ASCA. It is therefore likely that IRAS18325-5926 has a quasi-periodic oscillation (QPO) similar to, but at a much longer period than, the QPO seen in some Galactic Black Hole Candidates. The power spectrum of the February data has several peaks, the second highest of which is consistent with a monotonic decrease in the X-ray period. The period change is then consistent with that expected from two massive black holes spiralling together due to the emission of gravitational radiation. This possibility is very unlikely but mentioned because of its potential importance. ",
        "introduction": "We recently discovered a 16~hr X-ray periodicity in the Seyfert galaxy IRAS~18325-5926 using data from ASCA (Iwasawa et al 1998). Nearly 9 cycles of the oscillation were observed, with a total amplitude of about 15 per cent. The active nucleus has similar properties to that of Seyfert 1 galaxies, except that the power-law continuum is slightly steeper than most (photon index $\\Gamma\\sim 2.1$) and there is moderate absorption by a column density of $\\sim 10^{22}\\pcmsq$. There is also a broad iron line in the X-ray spectrum (Iwasawa et al 1996) which indicates the presence of an accretion disk in the X-ray emission region, viewed at moderate inclination (40--50 deg). The periodicity is plausible for Keplerian motion at 10--20 gravitational radii around a black hole of mass $2\\times 10^8 - 2\\times 10^7\\Msun$. It also scales well with quasi-periodic oscillations (QPO) seen from some Galactic Black Hole Candidates (BHC; Belloni et al 1997). The cause of the oscillation in IRAS~18325-5926, or indeed in any BHC, is unknown.  In order to determine whether the variation is exactly periodic or a QPO, we observed IRAS~18325-5926 with the Rossi X-ray Timing Explorer in 1997 December 25--27 and 1998 February 21--23. We report here the light curves and power spectra of those observations, which also show oscillations, but of different periods. We conclude that the AGN has a clear QPO signal and note that our results are consistent with the exciting possibility that 2 black holes are rapidly spiralling together. ",
        "conclusions": "\\begin{table} \\begin{center} \\caption{Mean flux and period for X-ray observations of IRAS~18325-5926. The Ginga data are reported by Iwasawa et al (1995). The error bars indicate the period change over which the power drops by a factor of two from the peak. 1998a and b refer to the highest and next highest peaks in the February power spectrum, respectively.} \\begin{tabular}{ccr} Detector & Flux (4--10~keV) & Period  \\\\ & $\\ergpcmsqps$ & ks \\\\[5pt] Ginga 1989  & $1.5\\times 10^{-11}$ & $>30$ \\\\ ASCA\\ 1994 & $5.6\\times 10^{-12}$ & $>80$ \\\\ ASCA\\ 1997 & $1.0\\times 10^{-11}$ & $58.0^{+2.6}_{-1.8}$ \\\\ RXTE\\ 1997 & $2.6\\times10^{-11}$ & $38.7^{+4.2}_{-2.9}$  \\\\ RXTE\\ 1998a & $2.1\\times 10^{-11}$ & $40.3^{+4.5}_{-2.7}$ \\\\ RXTE\\ 1998b & $2.1\\times 10^{-11}$ & $28.2^{+1.9}_{-1.4}$ \\\\ \\end{tabular} \\end{center} \\end{table}  The X-ray emission from IRAS~18325-5926 oscillates in a manner similar to that seen in BHC QPO. Such a large clear oscillation from matter around a black hole may suggest that we are detecting the fundamental frequency of some space-filling corona above the disk. This may be possible if the corona has a proton temperature close to the virial value and a lower electron temperature (see e.g. Di Matteo, Blackman \\& Fabian 1997). Of course this must happen over a very restricted range of radii in order that a single dominant oscillation is seen. Perhaps it corresponds to the radius where the surface emission from the disk peaks (i.e. $\\sim$7 Schwarzschild radii for a non-spinning hole). The variation of period with flux (Fig.~5) is similar to that seen in some QPO.  Otherwise, as mentioned by Iwasawa et al (1998), it could be due to the Bardeen-Petterson effect if the angular momentum vectors of the black hole and accreting matter are not aligned. A range of radii are then selected over which the accretion disk tilts over to match the equatorial plane of the black hole. (The disk is not actually precessing in this case; Markovic \\& Lamb 1998.) Reflection/obscuration by blobs of gas in the tilt zone might then lead to observed flux variations, especially if the inclination is fairly high (as the broad iron line may indicate; Iwasawa et al 1996), but we consider it unlikely that 50 per cent variations can be obtained in this way.  \\begin{figure} \\centerline{\\psfig{figure=pflux.ps,width=0.48\\textwidth,angle=270}} \\caption{Variation of period with 4--10~keV flux. The two highest power spectrum peaks are shown for February.} \\end{figure}  The data are consistent with a continuous decrease of period with time from GINGA observations in 1989 to December 1997, and continuing if the second highest peak in the February power spectrum is used. This raises the possibility that the period is due to a massive object orbiting the black hole (see Cunningham \\& Bardeen 1973). The separation of the objects requires that this second object is compact. The period decrease is then explained as due to orbital energy loss by the emission of gravitational waves (as for the binary pulsar). The rate of period change $$\\dot P ={{-96}\\over5}{{G^{5/3}}\\over c^5}{{M^{2/3}\\mu(4\\pi^2)^{4/3}}\\over P^{5/3}},$$ where $M=M_1+M_2$, the sum of the separate masses, and $\\mu=M_1M_2/(M_1+M_2).$ This integrates to give $P\\propto(t_0-t)^{3/8}$ where $t_0$ is the time when the masses finally merge.  We have fitted this last relation to the data and find an acceptable fit (Fig.~6). It suggests that the merger may occur in late April 1998. The rate of spiral-in enables us to estimate that the product $M^{2/3}\\mu\\approx 1.5\\times 10^{10}$ where the masses are in units of $\\Msun$. This means that for $M_1\\sim 2\\times 10^6\\Msun - 10^8\\Msun$, $M_2>10^5\\Msun$. There is no solution for $M_1<1.5\\times 10^6\\Msun$. Much of the energy from such a merger would emerge as gravitational waves, but some is likely throughout the electromagnetic spectrum.  We note that it is {\\it a priori} unlikely to find an object close in time to a spiral-in merger event. Possibilities that can enhance that probability are if a) black holes are built out of many merger events where the basic unit has a mass of only a few $10^5\\Msun$ (say from dwarf galaxies) and b) the inspiralling black hole switches on, or significantly enhances, an otherwise quiescent active nucleus. If we make the extreme argument that this process happens in all galaxies of space density $3\\times10^{-3}n_{-2.5}\\Mpc^{-3}$, all of which have a central black hole of mass $5\\times10^7M_{7.5}\\Msun$ growing by the addition of smaller black holes of $10^5m_5\\Msun$, then a merger will take place within $120\\Mpc$ (the distance to IRAS18325-5926) every $1000m_5 M_{7.5}^{-1} n_{-2.5}^{-1}\\yr$. We therefore see that such an event is not completely improbable but is unlikely. The probability that we should find the signal of a merger in its last year (it was the ASCA result of an observation in 1997 March which alerted us to make the RXTE observations) is $\\sim 10^{-3}$. Such mergers would of course be more common at fainter flux levels ($\\sim 1\\yr^{-1}$ within 1 Gpc), and the prospects for space-based gravitational wave astronomy (which will be sensitive to events in massive black holes) could be very positive.  A lower mass black hole captured by a central black hole is likely initially to have a highly eccentric orbit (Sigurdsson 1997). When the eccentricity $e$ is high, the above estimates for $\\dot P$ are dramatically increased (Shapiro \\& Teukolsky 1983), for example by a factor of 100--1000 when $e\\sim 0.8-0.9$, respectively. This allows $M_2$ to be much lower than estimated above for a circular orbit (it can be as low as about $100\\Msun$). The overall temporal behaviour of the system is then more complex ($e$ decreases with time).  It may be difficult for a low mass black hole to modulate the observed X-ray emission greatly; this problem is offset by the enhanced probability of such an event occurring.  We stress that this last, exciting, interpretation involving an in-spiralling black hole is unlikely and depends upon the identification of the second highest power peak in the February 1998 data with the orbital period of the second black hole. It does not explain the other peaks in the power spectrum nor why some peaks expected in the lightcurve do not occur. We discuss it here only because of its potential importance. Clearly further observations are urgently required. Even if such observations fail to show a consistently decreasing period, IRAS1832-5926 has a convincing, high-amplitude, QPO, which requires explanation.  \\begin{figure} \\centerline{\\psfig{figure=spiral_in_new.ps,width=0.48\\textwidth,angle=270}} \\caption{The peak period displayed against time in days, measured backwards from the February 1998 observation. Error bars show where the peak power has decreased by about a factor of 2. The line shows the best-fitting relation $P=6203.(t+63.5)^{3/8}\\s$. For the February data we plot both the period with the peak power and the longer one which had slightly less power.} \\end{figure}"
    },
    "9803/astro-ph9803243_arXiv.txt": {
        "abstract": "We report the discovery of two low redshift HI 21cm absorbers, one at $z = 0.2212$ towards the $z_{em} = 0.630$ quasar OI 363 (B0738+313), and the other at $z = 0.3127$ towards PKS B1127-145 ($z_{em} = 1.187$).  Both were found during a survey of MgII selected systems at redshifts $0.2 < z < 1$ using the new UHF-high system at the Westerbork Synthesis Radio Telescope (WSRT).  New HST/FOS observations also identify both systems as damped Ly$\\alpha$ (DLa) absorbers.  By comparing the column density from the DLa line with that from the HI 21cm line, we calculate the spin temperature, T$_s$ of the two systems.  We find $T_s \\approx 1000$ K for both of these low redshift absorbers.  For the $z = 0.3127$ system towards PKS B1127-145, two galaxies have been previously identified with emission lines at the absorber redshift (Bergeron \\& Boiss\\'e, 1991), with the galaxy at a closer projected distance to the quasar assumed to be responsible for the absorption system.  An ESO-NTT/EFOSC2 spectrum of a 3rd, fainter companion at 3.9 arcsec or 11 $h^{-1}_{100}kpc$ from the line of sight of PKS 1127-145 reveals [OIII]4958 and 5007 at $z = 0.3121 \\pm 0.0003$.  We consider this object the most likely to be responsible for the 21cm absorption, as it is much closer to the QSO sightline than the two galaxies identified by Bergeron \\& Boiss\\'e.   \\end {abstract} ",
        "introduction": "The study of low redshift examples of the quasar absorption line systems responsible for the damped Ly$\\alpha$ (DLa) and HI 21cm absorption lines is important to help bridge our understanding of neutral gas-rich systems between those at redshift $z = 0$ and those at high $z$.  Our knowledge of the neutral gas in nearby spiral galaxies is mainly based on observations of the HI 21cm line in emission.  At high redshift, however, we observe the HI 21cm line in absorption, which can only be seen along a limited number of lines of sight through the intervening absorber, making detailed knowledge of the gas characteristics difficult.  The low redshift ($z < 1$) neutral absorbers are still close enough that both optical and radio data of reasonable quality can be obtained in order to understand their kinematics and physical gas characteristics.  Such information is necessary to build a framework for a correct interpretation of the higher z counterparts to these systems.  Searches for redshifted HI 21cm absorption can be time-consuming since radio spectrometers typically observe only relatively narrow instantaneous bandwidths and only the highest column density QSO absorption line systems have measurable optical depths in the 21cm line.  Since DLa systems have high column densities of neutral HI, they are the most likely objects to have HI 21cm absorption. Unfortunately for low redshift work, the Ly$\\alpha$ line is not shifted into the optical window until $z \\simeq 1.65$, so finding these lines requires UV spectra to be taken with space telescopes. This, combined with the small cross section for DLa absorption, means that only a small number of DLa systems have been identified at low redshift.  Therefore, alternative selection criteria which are reasonably effective at finding an HI 21cm absorber must be used.  All known DLa and HI 21cm absorbers have associated low-ionization metal lines (cf. Lu and Wolfe, 1994) such as the MgII $\\lambda\\lambda$2796, 2803 doublet, which can be observed easily in ground-based spectra down to about $z = 0.1$.  A study of MgII selected systems using previously existing UV data yielded about 1 DLa system per 10 MgII systems observed (Rao, et al., 1995).  A similar statistic exists for HI 21cm absorption in MgII systems (Briggs \\& Wolfe, 1983).  This suggests that MgII can be used to optically select systems likely to have high column densities of neutral gas observed either as DLa or HI 21cm absorption.  We are conducting a survey for HI 21cm absorption in low redshift MgII selected systems using the Westerbork Synthesis Radio Telescope (WSRT).  In this paper we present two new HI 21cm absorbers from our survey, one towards PKS B1127-145 at $z = 0.3127$ and the other towards OI 363 at $z = 0.2212$.  Recent HST/FOS spectra have identified DLa absorption in both of these objects as well. ",
        "conclusions": "\\subsection{OI 363} OI 363 (0738+313), is a core dominated quasar at $z_{em} = 0.630$. Observations at 1640 MHz (Murphy, et al., 1993) show that the lobes extending $\\approx 30\\arcsec$ from the core contain only $3\\%$ of the total flux of the quasar.  The quasar is slightly variable at low frequencies (Bondi, et al., 1996b).  The metal line absorption system was originally reported by Boulade, et al. (1987) at a redshift of $z = 0.2213$, and subsequently by Boiss\\'e, et al. (1992) at a redshift of $z = 0.2216 \\pm 0.0003$. The only identified lines in the spectrum are the MgII $\\lambda\\lambda2796,2803$ doublet and a possible MgI line.  Deep optical imaging was made by LeBrun et al. (1993), who identified what they considered a likely absorber at a projected separation from the quasar of 5.70'', or $1.24R_H$ if it lies at $z = 0.221$.  They identified three additional galaxies, at smaller angular separations from the quasar, which are very faint and would be dwarf galaxies at the absorption redshift.  Unfortunately, they do not report a confirmed redshift for any object in the field near the quasar.  The HI 21cm absorption line, shown in fig. 1, has a narrow width of only two channels over most of its depth, and may be unresolved by this observation.  This implies an upper limit to the line width of $\\sim8$ km s$^{-1}$ at the 6 km s$^{-1}$ resolution of the data.  The HI column density from the DLa profile is $N(HI) = 7.9\\pm1.4 \\times 10^{20}$ cm$^{-2}$, and the calculated mean harmonic spin temperature is T$_s = 1230 \\pm 335$ K.  The termal kinetic temperature of the gas for a line of width 8 km s$^{-1}$ is T$_k = 1400$ K.  This means that within the errors, T$_s = $T$_k$.  \\subsection{PKS B1127-145}  PKS B1127-145 is a compact, gigahertz peaked radio source at $z = 1.187$.  VLBI observations at 1670 MHz show a slightly elongated structure with an extent of approximately 20 mas. Observations at 408 MHz give a flux variability of 1.2 Jy/year (Bondi, et al., 1996a), and indicate structural variations as well.  Fig. 2 shows the neutral HI 21cm absorption for this system.  The observed flux of the quasar is 5.25 Jy in this observation.  This is somewhat lower than reported fluxes in NED\\footnote{The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.}, however, as discussed above the quasar is a known variable at low frequency.  If the absorption line is fit by one gaussian component, the optical depth of the absorption is $6.2\\%$, and the FWHM is approximately 42 km s$^{-1}$.  However, there is some evidence for structure in the line.  In particular, the split in the middle appears to be real, suggesting at least two components, and the asymmetric low frequency side of the profile may result from a third component.  There is some low level interference in the spectrum adjacent to the low frequency side of the absorption line, at $\\sim$1081.75 MHz, which makes interpretation of the real shape difficult.  The HI column density from the DLa profile is $N(HI) = 5.1\\pm0.9 \\times 10^{21}$ cm$^{-2}$, and the calculated T$_s = 1000 \\pm 200$ K.  Bergeron and Boiss\\'e (1991) studied this $z_{abs}=0.313$ metal-line system in some detail.  They identify MgII, FeII and MgI absorption in the spectrum.  The average redshift of the metal lines is $z = 0.3127 \\pm 0.0002$.  In addition to the absorption spectrum, and deep images of the field, they present emission spectra for two of the bigger bright galaxies near the quasar (different from the two faint companions discussed in Sect. 4), both of which have redshifts similar to that of the metal line system. They identify the galaxy at the smaller projected distance from the quasar sightline as the absorber.  The ESO/EFOSC2 spectral observations (discussed in Sect. 3) identify a new candidate absorber galaxy for this system. The new emission object is a close companion to the west of the quasar and lies at a projected distance of $3.9\\arcsec$ or 11 $h^{-1}_{100}$kpc from the quasar sightline.  Based on comparison with other objects in our image for which apparent magnitudes are given in Bergeron and Boiss\\'e (1991), it has an apparent magnitude of $m_r = 22.3$.  The galaxy which Bergeron and Boiss\\'e identify as the absorber is at a projected separation of $9.6\\arcsec$, or 37 $h^{-1}_{100}$kpc from the quasar at the redshift of the galaxy.  This corresponds to a galactic radius of $2.7 R_H$.  A column density of neutral gas of $10^{21}$ cm$^{-2}$ is unlikely at this galactic radius, and we consider the new emission object, although smaller and fainter, to be the more likely absorber given its proximity to the quasar sightline.  It is also possible that the three galaxies at similar redshift have undergone strong interaction with each other, in which case the absorbing gas could be tidal debris.  \\subsection{Spin Temperature}  Fig. 4 shows T$_s$ vs. redshift for all of the known HI 21cm and DLa absorber systems, calculated from the literature (as summarized by Carilli, et al., 1996), with the two new data points marked by open symbols.  The shaded region shows the range of Galactic T$_s$ values at optical depths comparable to those found in the DLa systems, using numbers from Braun and Walterbos (1992).  The large errorbars in any given measurement are dominated by uncertainties in the true optical continuum level, due to confusion from the Ly$\\alpha$ forest lines, which make fitting the damped profile difficult.  As noted in previous studies (cf. de Bruyn, et al.  1996), all of the redshifted absorbers except the highest optical depth (lowest T$_s$) system have T$_s$ values which are roughly two or more times greater than the Galactic values at similar optical depths.  Most estimates of T$_s$ for the present epoch have been based on studies of the Milky Way or Andromeda (cf. Dickey and Brinks (1988); Braun and Walterbos (1992), and references therein).  These estimates use column densities derived from HI 21cm emission lines rather than from DLa absorption lines, but studies have shown that N(HI)$_{Ly\\alpha} \\approx$ N(HI)$_{21cm}$ (Dickey and Lockman, 1990) . This suggests that T$_s$ values for the galaxy and the redshifted DLa systems, although calculated from different quantities, can be compared.  On the other hand, values for T$_s$ have a strong correlation with the column density, N(HI), in each of the clouds along the line of sight in the Galaxy, and are thus sensitive to the location of individual gas clouds.  A single cloud usually has a greater angular size on the sky than the beam with which it is observed, so it is possible to calculate the spin temperature for just one cloud.  This is not the case in redshifted systems, where a line of sight includes more of the galaxy and possibly many clouds.  Thus it is difficult to usefully compare present epoch spin temperature values to those calculated at higher redshifts.  The new systems in this study were observed at low redshift with the same line of sight limitations as the higher redshift systems, allowing a meaningful comparison.  There is no clear trend for a change in T$_s$ with increasing redshift among the eight systems shown in Fig.  4.  If the gap between T$_s$ values in the galaxy and those in the DLa systems is an effect of evolution over time, then all of that evolution must have occured between $z = 0.2$ and the present epoch.  Instead, we consider it more likely that the gap arises from the presence of many clouds in one line of sight at high redshift, or from differences in the radio and optical sightlines.  The sizes of the radio emission regions of quasars are much larger than the optical regions, and usually larger than an average cloud as well. It is therefore likely that the optical and radio lines of sight actually sense different clouds in a redshifted galaxy, and hence have different column densities of neutral gas.  This implies that the spin temperatures derived by assuming both column densities are equal may be meaningless.  Unfortunately, the resolution obtained in most radio survey observations (with single dishes or synthesized beams from arrays like the WSRT) is at least as large on the sky as compact background radio sources, and gives no spatial information about the clouds in front of the quasar.  Future use of VLBI techniques to pinpoint the radio sightlines more accurately may help to clarify this problem."
    },
    "9803/hep-ph9803295_arXiv.txt": {
        "abstract": "The effects of a possible rotation of the galactic dark halo on the calculation of the direct detection rates for particle dark matter are analyzed, with special attention to the extraction of the upper limits on the WIMP--nucleon scalar cross section from the experimental data. We employ a model of dark halo rotation which describes the maximal possible effects. For WIMP masses above 50 GeV, the upper limit exclusion plot is modified by less than a factor of two when rotation is included. For lighter masses the effect can be stronger, suggesting the necessity to develop specific models of halo rotation in order to provide more accurate conclusions. ",
        "introduction": "The possibility to detect Weakly Interacting Massive Particles (WIMPs) distributed in the halo of our Galaxy has been a major issue in the last years, since these particles could provide the amount of dark matter necessary to explain many observed dynamical properties of galaxies, clusters and of the Universe itself. Different kinds of possible signals have been identified and looked for, in order to outline the presence of WIMPs in our Galaxy. These signals are usually referred to as ``direct\" and ``indirect\" detection rates. Direct detection refers to the possibility to measure a WIMP--nucleus interaction in a low--background detector, while indirect detection relies on the measurement of WIMP annihilation products: photons, antiprotons and positrons produced by the annihilation in the galactic halo, or neutrinos coming out of the Earth or the Sun where WIMPs may have been accumulated as a consequence of gravitational capture. It is remarkable that the present sensitivity of the different experiments is already at the level of the predicted rates for specific WIMP candidates, like the neutralino, which represents one of the most interesting and studied cold relic particles \\cite{noi}.  The calculation of the different detection rates depends not only on the particle physics properties of the WIMPs interactions, but also on the characteristics of the galactic halo where the WIMPs are distributed. Direct detection rates and upgoing-muon fluxes at neutrino telescopes, which both rely on the WIMP elastic scattering off nuclei, depend on the WIMP matter density $\\rho_\\odot$ and velocity distribution $f_\\odot(v)$ at the Earth position $r_\\odot$ in the Galaxy. In particular, the dependence of the signals on $\\rho_\\odot$ is a linear one. The other indirect signals (photon, antiproton and positron fluxes) have a stronger dependence on the matter distribution, since they are proportional to the square of the matter distribution function (DF) $\\rho(\\vec r)$ integrated over the effective region of production and propagation of the annihilation products. On the contrary, this kind of signals are essentially independent on the details of the velocity DF, since the annihilating WIMPs are almost at rest and corrections due to their velocity dispersion are negligible.  Detailed estimates of the detection rates would require specific and accurate models of the galactic halo able to provide a reliable WIMPs DF $g(\\vec x, \\vec v)$ (not necessarily separable in phase space, i.e. $g(\\vec r, \\vec v) = \\rho(\\vec r) f(\\vec v)$). Unfortunately, detailed halo models are not available at present, mainly because the constraints obtained from astrophysical observations are not stringent enough to restrict different possibilities. The most important observational constraint is provided by the flatness of the rotation curves at large radii. Although the available data on our Galaxy do not provide a compelling evidence of a flat rotation curve, this feature is observed in a large number of spiral galaxies and therefore it looks reasonable to assume its validity also for our Galaxy.  The standard and simplest model of the dark galactic halo, which is compatible with a flat rotation curve, is the so--called isothermal sphere. This model relies on the two basic assumptions of spherical symmetry and thermal equilibrium, which find a strong support in the argument of ``violent relaxation\" introduced by Lynden--Bell 30 years ago\\cite{lynden-bell}. In this model, the DF is separable into a matter density distribution $\\rho(r)$, which has a $r^{-2}$ behaviour at large radii, and into a Maxwell--Boltzmann (MB) velocity DF $f(v)$ \\cite{binney}. Although such a model gives a divergent total mass and therefore an appropriate cut-off has to be introduced at large radii, its range of validity has been tested at least in the inner parts of many galactic systems. Moreover, since it represents a simple and reasonable approximation, in the absence of a more detailed model it is widely adopted to describe the dark halo of our Galaxy. However, many different models are known to be consistent with flat rotational curves.  For  instance, models  which  describe non--spherically  symmetric or flattened halo  distributions have been discussed  \\cite{binney}. In these models, the specific form of $\\rho(\\vec r)$ differs from the standard isothermal sphere matter DF, especially at  small radii,  entailing quite large uncertainties  on the  local value  $\\rho_\\odot$. A comprehensive numerical study which takes into account a large number of models  indicates that the local  value of the   non--baryonic dark  matter  density falls in  the (rather conservative) range 0.1 $\\lsim  \\rho_\\odot  \\lsim$  0.7 GeV cm$^{-3}$\\cite{turner}. Contrary to the matter DF, the specific form of the velocity DF $f(v)$ has been much less investigated. Modifications to the standard MB velocity DF are known \\cite{binney,evans}, but the  problem of determining  the correct form of the distribution of the WIMP velocities in the halo has  no clear  and simple   solution at  present,  both  theoretically and observationally.   The   velocity DF is required to be consistent  with a given  $\\rho(\\vec r)$ but this,  in general,  does not determine $f(v)$ in a unique way.  The calculation of the  WIMP detection rates is  usually performed by using the standard  isothermal sphere model. However, modifications in the isothermal  model  can  affect the detection rates, introducing uncertainties in the theoretical predictions and in the extraction of the experimental limits on the WIMPs parameters. The effects induced on the detection rates by a modification in the matter DF are simple to take into account, since the dependence of the detection rates on $\\rho(\\vec r)$ can be factorized. Specifically, the physical range of $\\rho_\\odot$ quoted above implies an uncertainty of about a factor of 7 in the evaluation of the direct detection rates and in the neutrino fluxes \\cite{noi} (it has to be remarked that this large factor reflects a rather conservative attitude). Even larger uncertainties affect the indirect rates from WIMP annihilation in the halo, since in this case a modification in the matter density profile can strongly affect the integral of $\\rho^2(\\vec r)$ over the effective production region of the signal \\cite{pbar_to,pbar_japan,gamma}. Contrary to the case of the matter DF, a modification of the standard MB velocity  DF would affect the direct detection  rates and the indirect rates at neutrino telescopes in a much  more  involved way. This is because the  dependence  of these  rates on  $f(v)$ is through a  convolution of   $f(v)$  with the  differential  WIMP--nucleus cross section. Since the  WIMP--nucleus scattering depends on the relative velocity of the WIMPs  with  respect to the  detector  nuclei, a potentially significant effect could be due to a bulk rotation of the halo. This would necessarily  modify the  WIMP phase--space  DF with respect to the standard MB form.  In this paper we wish to discuss the possible effects  induced by a halo rotation on the direct detection  rates, with special attention to the ensuing consequences on the determination   of the  upper  limits  on the   WIMP--nucleus  cross  section from the experimental data. A calculation of  the direct  detection rates in  the case of a  rotating halo has been  addressed in  Ref.  \\cite{kamion},  where it has  been  concluded that the maximal effect of rotation leads to a 30\\% effect on the total detection rates for a Ge nucleus, in the case of an ideal detector with no threshold. However,  when  considering a real  detector the  behaviour of  the differential rates at  threshold and the detector characteristics are crucial in  determining the  experimental  limits on the WIMP--nucleus  cross  section \\cite{noi}.  Therefore, we  explicitly take into account the  features of running detectors,  such as thresholds, quenching factors and energy resolution, in order to estimate the largest uncertainties induced by a possible halo rotation in a confident way. To this aim, following Ref.\\cite{kamion} we model the galactic rotation as described by Lynden--Bell in Ref.\\cite{lynden-bell2}, where the maximally rotating velocity DF compatible with a given mass distribution has been derived, on the ground of purely kinematical arguments. Even if Lynden--Bell's model of halo rotation may not represent a situation which is realized in a physical halo, we consider it useful to bracket the size of the effect of halo rotation on the direct detection rates.  The plan of our paper is the following. In Sect.II we briefly describe the calculation of the direct detection rates in the presence of halo rotation. In Sect.III we discuss our results for Ge, NaI and Xe detectors, taking into account the most recent experimental data of the different Collaborations. Finally, in Sect.IV we draw our conclusions. An Appendix is added, where we report the analytical expressions of the relevant part of the direct detection rates which contain the details of the velocity DF in the case of the standard non--rotating, maximally co--rotating and maximally counter--rotating haloes. ",
        "conclusions": "In this paper we have investigated the effect induced by a possible rotation of the galactic halo on the rates of WIMP direct detection. In particular, we have discussed the implication of halo rotation on the determination of the exclusion plots on the WIMP--nucleon cross section for different detectors, namely Ge, NaI and Xe ones. The rotation of the halo has been described by using a model \\cite{lynden-bell2} which corresponds to a situation where the halo possesses the maximal rotation compatible with a given mass DF, which for simplicity we have chosen to be that of the isothermal sphere.  We found that the exclusion plots  obtained from the data are affected by  less than a factor of 2 in the  case of  counter--rotating models. The same  size of uncertainty  occurs also for  the co--rotating  models, when  the WIMP  mass is larger than  about 50  GeV. For  lighter WIMPs and  co--rotation, the exclusion plots are  modified  by a larger  amount. We  have to  remind  that, due to the particular model of halo  rotation which we have  employed here,  these are expected to be maximal  effects. For  specific physical  rotation  models, the effect of halo rotation  will be  plausibly smaller. We  can therefore  conclude that, at least for  WIMP  masses greater  than about  50 GeV,  the  determination of the exclusion  plots from  the  experimental  data are  affected by an  uncertainty smaller  than  a factor of two  due to  the possibility that the  galactic halo rotates,  independently on the  specific model of halo  rotation. We notice that recent preliminary data from accelerators indicate that the lower limit on the mass of the most plausible WIMP candidate, the neutralino, is $m_\\chi\\simeq 30$ GeV for low value of the susy parameter $\\tan\\beta$, and $m_\\chi\\simeq 45$ GeV for $\\tan\\beta\\gsim 3$ \\cite{lep183}. Therefore, for this dark matter candidate, the uncertainty on the exclusion plot due to a possible rotation of the halo is expected to be relatively small. The situation  is different for lighter WIMPs. In this  case,  it would  be required to develop  specific models  of halo  rotation in  order to obtain  more accurate conclusions."
    },
    "9803/astro-ph9803239_arXiv.txt": {
        "abstract": "We report the discovery of a 2.1hr optical modulation in the transient source GS1826-24, based on two independent high time-resolution photometric observing runs.  There is additional irregular variability on shorter timescales.  The source also exhibited an optical burst during each observation, with peak fluxes consistent with those of the three X-ray bursts so far detected by {\\it Beppo}SAX.  We compare the low-amplitude variation ($\\sim 0.06^m$) to that seen on the orbital periods of the short period X-ray bursters, X1636-536 and X1735-444, as well as the similarity in their non-periodic fluctuations.  Other transient neutron star LMXBs possess short periods in the range 3.8-7.1 hrs.  However, if confirmed as the orbital, a 2.1 hr modulation would make GS1826-24 unique and therefore of great interest within the context of their formation and evolution. ",
        "introduction": "GS 1826-24 was discovered serendipitously in 1988 by the {\\it Ginga} LAC during a satellite manoeuvre (Makino {\\it et al.} 1988). The source had an average flux level of 26 mCrab (1-40 keV), and a power law spectrum  with $\\alpha = 1.7$. Observations both a month before and after by the {\\it Ginga} ASM, by TTM in 1989 \\cite{intZ92} and by ROSAT in 1990 and 1992 \\cite{barr95} found comparable flux levels.  Temporal analyses of both the {\\it Ginga} detection and ROSAT data yielded a featureless $f^{-1}$ power spectrum extending from $10^{-4} - 500$ Hz \\cite{tan95,barr95}, with neither QPO nor pulses being detected.    Despite its detection by {\\it Ginga}, the source had not been previously catalogued.  Neither were X-ray bursts detected by {\\it Ginga}.  Together with its similarities to Cyg X-1 (hard X-ray spectrum, strong flickering), this led to an early suggestion by \\scite{tana89} that it was a soft X-ray transient with a possible black-hole primary.  Later, \\scite{stri96} called this suggestion into doubt, following examination of data from CGRO/OSSE observations. They found that fitting both the Ginga and OSSE spectra produced a model with an exponentially cutoff power law plus reflection term.  The observed cut-off energy of $\\sim$58 keV is typical of the cooler neutron star hard X-ray spectra.  The recent report of three X-ray bursts detected by {\\it Beppo}SAX \\cite{uber97} and our detection of optical bursts here confirms the presence of a neutron star accretor.  Following the first ROSAT/PSPC all-sky survey observations in September 1990, and the determination of a preliminary X-ray position, a search for the counterpart yielded a time variable, UV-excess, emission line star \\cite{motc94,barr95}. This source had $B= 19.7 \\pm 0.1$, and an uncertain V magnitude of $V \\simeq 19.3$, due to contamination by a nearby star.  There was also evidence for $\\simeq 0.3^m$ variations on a one hour timescale, but the time sampling was fairly poor. For this reason, we included this object in our target list for a high-speed photometry run at the South African Astronomical Observatory (SAAO). Further time-series photometry was also obtained on the William Herschel Telescope, La Palma (WHT) to confirm the variability that was seen. ",
        "conclusions": "\\begin{table*} \\begin{minipage}{150mm}  \\caption{Properties of GS1826-24 and the short-period transient bursters$^a$\\label{tab:sptb}} \\footnotetext[1]{Data taken from \\scite{vP95} and references therein, unless cited elsewhere.} \\begin{tabular}{c c c l c l l l } \\hline Source \t& $P_orb$ & V \t\t& B-V, U-B \t&\tE${\\rm _{B-V}}$\t\t& F$_X$\t&\tActive &Quiescent \t\t\t\\\\ &\t(hr)\t&\t\t\t&\t& & ($\\mu$Jy)&\t&  \\\\ \\hline GS1826-24&\t2\\footnote[2]{This work.}\t\t&19.3\t\\footnote[3]{\\scite{motc94}.}\t& 0.4$^c$, -0.5$^c$\t&\t0.4$^c$\t& 30$^c$& 1988-present\\footnote[4]{\\scite{mak88}.} & before 1988\\\\ &\t\t\t&\t\t\t&\t&\t\t\t&\t\t&&\\\\ X0748-678&\t3.82\t&$>$23\\footnote[5]{\\scite{wad85}.}-16.9&0.1, -0.9\t&0.42\t\t& 0.1-60\t& 1985-present\\footnote[6]{\\scite{whi95}.} &before 1985\\\\ &\t\t\t&\t\t\t&\t\t&\t\t&\t&\t\\\\ X2129+470 &\t5.24\t&16.4-17.5&\t0.65,-0.3 &0.5\t\t&\t9\t& $<$1979 - 1983$^f$ & 1983-present\\\\ &\t\t\t&\t\t\t&\t&\t\t&\t\t\t&\t\t&\\\\ X1658-298 & 7.11&  18.3\t\t&0.45,-0.4\t\t& 0.3 & $<$5-80& pre-1980s $^f$ & 1980-present\t\\\\ &\t\t\t&\t&\t&\t\t\t&\t\t&&\\\\ \\end{tabular}   \\end{minipage} \\end{table*}   The period distribution of low-mass X-ray binaries (LMXB) has shown a scarcity of systems below 3 hr presumably in part due to their faintness, but there is a notable absence of any systems with periods between \\til 1 and 3 hr \\cite{whi85a,whi85b}.  Recently, \\scite{king97} have investigated the formation of neutron star LMXBs.  They found that in order to produce the relatively large fraction of soft X-ray transients in the $\\simlt$ 1-2 day range, the secondaries must have 1.3\\msun $\\simlt M_2\\simlt$ 1.5\\msun\\ at the onset of mass transfer and be significantly nuclear evolved (provided that the SN kick-velocity is small compared to the pre-SN orbital velocity).  This mass range ensures that the mass transfer rates driven by angular momentum loss are below the critical rate needed for SXT behaviour in the required fraction of neutron star LMXBs.  The large initial masses then account for the rarity of any short $P \\simlt 3$ hr systems.  If indeed the 2 hr modulation of GS1826-24 is confirmed as orbital in origin, this would make it of great interest within the context of transient neutron star LMXB formation and evolution.  However, in many respects GS1826-24 does show similarities to other neutron star binaries with comparable periods.  The X-ray bursters X1636-536 and X1735-444 exhibit optical modulations of a similar amplitude on their orbital periods (3.80 and 4.65 hr respectively), plus irregular variability on somewhat shorter timescales (\\pcite{whi95} and references therein), although unlike GS1826-24 these systems are not transient on a timescale of decades. The optical emission of LMXBs is dominated by the reprocessed X-rays from the accretion disc and companion.  The origin of the underlying sinusoidal modulation is interpreted as the varying contribution from the X-ray heated face of the companion star \\cite{vP95a}. As for the irregular variability, this is probably caused by changes in the spatial distribution of the reprocessing material in the accretion disc and/or fluctuations in the central X-ray luminosity, as suggested in the case of X1636-536 \\cite{vP90b}.  Moreover, GS1826-24 has a high $L_X/L_{opt}(\\sim500)$, very similar to that of the compact 41 min binary X1627-673, but lower than the $L_X/L_{opt}\\sim700$ of the 50 min binary X1916-053, which is a higher inclination dipping source (once again these are persistent sources).  Since this ratio is related in part to the physical size of the system and the disc  reprocessing area available, the similarity supports the hypothesis that GS1826-24 is also a relatively compact system.   Furthermore, the non-detection of GS1826-24 prior to 1988 is characteristic of the observed variability of the transient bursting systems X0748-678, X1658-298 and X2129+470 \\cite{whi95}, which have either been detected over many years and then gone into quiescence for a similar period of time or vice-versa.  These transients all have known periods in the range 3.82-7.1hrs (see Table \\ref{tab:sptb}).  Clearly, further high-speed optical photometric monitoring is required in order to confirm the stability of the 2hr modulation and hence its orbital origin.   With only a short {\\it ROSAT} lightcurve published to date \\cite{barr95}, confirmation might be possible from a longer X-ray observation. However, the low-amplitude of the observed modulation ($0.06^m$) implies a low-inclination system ($< 70 \\degree$), and hence X-ray dipping behaviour is unlikely to be seen."
    },
    "9803/astro-ph9803149_arXiv.txt": {
        "abstract": "The local luminosity function at 25 $\\mu$m provides the basis for interpreting the results of deep mid-infrared surveys planned or in progress with space astrophysics missions including ISO, WIRE and SIRTF. We have selected a sample of 1458 galaxies from the IRAS Faint Source Survey with a flux density limit of 250 mJy at 25 $\\mu$m.  The local luminosity function is derived using both parametric and non-parametric maximum-likelihood techniques, and the classical $1/V_{max}$ estimator. Comparison of these results shows that the $1/V_{max}$ estimate of the luminosity function is significantly affected by the Local Supercluster.  A maximum-likelihood fit to the radial density shows no systematic increase that would be caused by density evolution of the galaxy population. The density fit is used to correct the $1/V_{max}$ estimate. We also demonstrate the high quality and completeness of our sample by a variety of methods.  The luminosity function derived from this sample is compared to previously published estimates, showing the prior estimates to have been strongly affected by the Local Supercluster.  Our new luminosity function leads to lower estimates of mid-infrared backgrounds and number counts. ",
        "introduction": "\\label{sec:intro}  Much of the effort to study infrared-luminous galaxies has centered on wavelengths greater than 50 $\\mu$m.  Modeling work is focused on the near-IR (e.g. \\markcite{chok94}Chokshi et al. 1994) and far-IR (e.g. \\markcite{hac87}Hacking et al. 1987; \\markcite{rrbinson96}Rowan-Robinson et al. 1996) portions of the galaxian spectrum. However, the mid-infrared is well-suited for studying starburst and ultraluminous galaxies.  About 40\\% of the luminosity from starburst galaxies is radiated from 8-40 $\\mu$m (\\markcite{soi87}Soifer et al. 1987).  Extinction effects are small, and problems due to infrared cirrus are minimized.  Most importantly for space astrophysics, for a fixed telescope aperture, the spatial resolution is higher at shorter wavelengths, and the confusion limit lies at higher redshifts.  Recent work using the {\\it Infrared Space Observatory (ISO)} (e.g. \\markcite{knapp96}Knapp et al.  1996; \\markcite{boul96}Boulade et al. 1996; \\markcite{rrobinson96}Rowan-Robinson et al. 1996) shows the relative importance of the 7 $\\mu$m and 15 $\\mu$m bands for galaxy studies.  The {\\it Wide-Field Infrared Explorer (WIRE)}, a Small Explorer mission due to launch in late 1998 (\\markcite{hac96}Hacking et al.\\ 1996; \\markcite{schemb96}Schember et al. 1996), will conduct a very deep survey at 24 $\\mu$m to study starburst galaxy evolution.  The {\\it Space Infrared Telescope Facility (SIRTF)} is also expected to conduct surveys in mid-infrared bands. To interpret the results of these surveys now in progress or soon to commence, it is necessary to better understand the mid-infrared properties of galaxies in the local Universe.  The 25 $\\mu$m luminosity function provides the basis for predicting the faint source counts in the mid-infrared.  The empirical model of \\markcite{hac91}Hacking \\& Soifer (1991) uses an analytic fit to the luminosity function derived by Soifer \\& Neugebauer \\markcite{soi91} (1991).  This function was estimated from a complete subsample of the Bright Galaxy Sample (\\markcite{soi87}Soifer et al.\\ 1987) containing 135 galaxies to a flux density limit of 1.26 Jy.  The availability of many more redshifts of IRAS galaxies (principally from the 1.2 Jy Survey (\\markcite{str90}Strauss et al.\\ 1990; \\markcite{fis95} Fisher et al.\\ 1995)) enables a much larger sample to be studied, reducing uncertainties at high and low luminosities.  In this paper we present the selection of a large galaxy sample that is flux-limited at 25 $\\mu$m, and derive the local luminosity function based on this sample.  The sample selection is described in the next section.  In Section \\ref{sec:lfresults} we describe the $1/V_{max}$ and the maximum-likelihood estimators for deriving the local luminosity function, and present the results. The completeness of the sample is discussed in Section \\ref{sec:compdisc}. In Section \\ref{sec:correct} we calculate the radial density distribution of the sample using a maximum-likelihood method. The radial density fit is used to correct the $1/V_{max}$ estimate of the local luminosity function, as well as the redshift distribution with which the luminosity function can be compared.  Section 6 includes discussions of the different luminosity function estimators, a comparison of our newly derived luminosity function with previous estimates, the implications for mid-infrared backgrounds and number counts, and the effects of evolution on the derivation of the luminosity function. The color properties of the sample are treated in another paper (\\markcite{fang98}Fang et al.\\ 1998). ",
        "conclusions": "\\label{sec:conclusions}  The following are the results of this paper:  1. We have selected a sample of 1458 galaxies with redshifts from the IRAS Faint Source Survey with a flux density limit of 250 mJy at 25 $\\mu$m.  An additional 17 galaxies do not have redshifts available.  2. The local luminosity function is derived using the $1/V_{max}$ estimator and both parametric and non-parametric maximum likelihood methods. The $1/V_{max}$ estimate is significantly affected by the Local Supercluster.  The maximum likelihood methods are independent of density variations, and we consider the parametric fit with parameters in Table \\ref{tab:fitpars} to be the best estimate of the local luminosity function at 25 $\\mu$m.  3. A maximum likelihood fit to the radial density in this sample is used to correct the $1/V_{max}$ estimate.  The fit shows no sign of a systematic increase with redshift of the density, as would result from density evolution of the galaxy population.  4. The $1/V_{max}$ luminosity function derived from a smaller sample by Soifer \\& Neugebauer (1991) is significantly contaminated by the Local Supercluster.  Predictions of number counts and local luminosity density based on that function are 15-20\\% higher than those indicated by our improved luminosity function.  The new function also leads to lower predictions of the mid-infrared background due to galaxies."
    },
    "9803/astro-ph9803180_arXiv.txt": {
        "abstract": "We report radial velocities for 99 galaxies with projected positions within $30\\arcmin$ of the center of the cluster A3733 obtained with the MEFOS multifiber spectrograph at the 3.6-m ESO telescope. These measurements are combined with 39 redshifts previously published by Stein (1996) to built a collection of 112 galaxy redshifts in the field of A3733, which is used to examine the kinematics and structure of this cluster. We assign cluster membership to 74 galaxies with heliocentric velocities in the interval $10\\,500$--$13\\,000$ \\kms. From this sample of cluster members, we infer a heliocentric systemic velocity for A3733 of $11\\,653^{+74}_{-76}$ \\kms, which implies a mean cosmological redshift of 0.0380, and a velocity dispersion of $614^{+42}_{-30}$ \\kms. The application of statistical substructure tests to a magnitude-limited subset of the latter sample reveals evidence of non-Gaussianity in the distribution of ordered velocities in the form of lighter tails and possible multimodality. Spatial substructure tests do not find, however, any significant clumpiness in the plane of the sky, although the existence of subclustering along the line-of-sight cannot be excluded. ",
        "introduction": "The rapid development of multifiber spectroscopy in recent years has made possible the simultaneous acquisition of large numbers of galaxy spectra. The obtention of extensive and complete redshift data bases for clusters of galaxies has hastened the investigation of the physical properties of their visual component which, in turn, is allowing for a better understanding of the characteristics of the dark matter distribution on Mpc scales. Here, we report a total of 104 redshift measurements for 99 galaxies in the field of A3733 and use these data, in combination with a previously published sample of 39 redshifts, to perform a kinematic and spatial analysis of the central regions of this cluster.  A3733 is a southern galaxy cluster listed in the ACO catalog (Abell, Corwin, \\& Olowin 1989)\\cite{ACO89} as of intermediate Abell's morphological type and richness class $R=1$. This cluster hosts a central cD galaxy, included in the Wall \\& Peacock (1985)\\cite{WP85} all-sky catalog of brightest extragalactic radio sources at 2.7 GHz, which has led to its classification as of Bautz-Morgan type I--II (Bautz \\& Morgan 1970)\\cite{BM70}. A3733 is also one of the 107 nearby rich ACO clusters ($R\\ge 1$, $z\\le 0.1$) included by Katgert et al. (1996)\\cite{Ka96} in the ESO Nearby Cluster Survey (ENACS), as well as a one of the X-ray-brightest Abell clusters detected in the ROSAT All-Sky Survey by Ebeling et al. (1996)\\cite{Eb96}.  The only major kinematical study of A3733 done so far is that of Stein (1997)\\cite{St97}. From a sample of 27 cluster members located within $r\\la 16\\arcmin$ from the cluster center, this author has found no evidence of significant substructure in the cluster core. This study of A3733, which is part of a more general investigation of the frequency of substructure in the cluster cores from an optical spectroscopic survey conducted on a sample of 15 nearby ($0.01\\la z\\la 0.05$) galaxy clusters (Stein 1996)\\cite{St96}, is based on a dataset that has many characteristics in common with the ENACS data gathered for the same field. Indeed, the two datasets have been obtained with the OPTOPUS multifiber spectrograph at the ESO 3.6-m telescope and cover essentially the same area on the sky. Besides, they have also a very similar number of galaxies: 39 and 44, respectively (28 of which are shared).  The MEFOS redshift dataset for A3733 reported in this paper contains two and a half times the number of galaxy radial velocities reported by Stein (1996)\\cite{St96}, including 26 reobservations, while it covers a circular region around the center of A3733 four times larger. Furthermore, its high degree of completeness offers the possibility of extracting a complete magnitude-limited subset with a number of galaxies large enough for its use on statistical analysis. The plan of the paper is as follows. In Sect.~2 we discuss the MEFOS spectroscopic observations and data reduction, and present a final sample with 112 entries built by the combination of the MEFOS and Stein's (1996)\\cite{St96} data. Section~3 begins with a brief description of the tools which will be used for the analysis of the data. Next, we identify the galaxies in our sample that belong to A3733, and use this dataset and a nearly complete magnitude-limited subset of it to examine the kinematical properties and structure of the central regions of the cluster. Section~4 summarizes the results of our study. ",
        "conclusions": "We have reported 104 radial velocity measurements performed with the MEFOS multifiber spectrograph at the 3.6-m ESO telescope for 99 galaxies in a region of $30\\arcmin$ around the center of the cluster A3733. To augment this data, we have combined the MEFOS measurements with 39 redshifts measured by Stein (1996)\\cite{St96} with the OPTOPUS instrument at the same telescope. This has given a final dataset with a total of 112 entries in the field of A3733. Radial velocities have been then supplemented by COSMOS \\bj\\ magnitudes and accurate sky positions in order to investigate the kinematics and structure of the central regions of the cluster.  From a sample containing 74 strict cluster members, we have derived a heliocentric systemic velocity for A3733 of $11\\,653^{+74}_{-76}$ \\kms, resulting in a $\\overline{z}_{\\mathrm CMB}$ of 0.0380, and a velocity dispersion of $614^{+42}_{-30}$ \\kms, in good agreement with the estimates by Stein (1997)\\cite{St97} from the OPTOPUS data alone. Statistical tests relying exclusively on the distribution of observed velocities have yield suggestive indication of the possible kinematical complexity of A3733, especially when applied to a nearly complete magnitude-limited (\\bj\\ $\\leq 18$) sample of cluster members. Despite this result, two powerful substructure tests that incorporate spatial information have failed to detect in this latter sample any statistically significant evidence of clumpiness in the galaxy component, in agreement with the findings of a previous study based on a spatially less extended and less complete dataset. Given that the sensitivity of the spatial substructure tests we have used is reduced when the subunits are seen with small projected separations, the results of the present study cannot exclude, however, the possibility that the signs of kinematical complexity detected in the velocity histogram of A3733 might be due to the existence of galaxy subcondensations superposed along the line-of-sight."
    },
    "9803/astro-ph9803194_arXiv.txt": {
        "abstract": "An analysis of the observed characteristics of the Galactic Cepheid variables is carried out in the framework of their \\plr\\ being used as a standard candle for the  distance measurement. The variation of the observed number density of Galactic Cepheids as function of their period and amplitude along with stellar pulsation characteristics is used to divide the population into two groups: one with low periods, probably multi-mode or higher mode oscillators, and another of high period variables which should be dominantly fundamental mode radial pulsators. Methods to obtain extinction-corrected colors from multi-wavelength observations of the second group of variables are described and templates of the $\\vi $ \\lig s are obtained from the $V$ \\lig s. Colors computed from the model atmospheres are compared with the extinction-corrected colors to determine the Cepheid instability strip in the {\\em mean surface gravity--effective temperature diagram}, and relations are derived between mean colors {\\em $\\BV$ vs period of pulsation}, {\\em $\\vi$ vs period}, and {\\em $\\vi$ at the brightest phase vs amplitude of pulsation}. The strength of the $\\kappa$-mechanism in the envelope models is used to estimate the metal dependency of the instability strip from which an idea of the sensitivity of the \\plr\\ to the helium and metal abundance is given. Some estimate of the mass of Cepheids along the instability strip is provided. ",
        "introduction": "\\label{sec:intro}  The classical Cepheid variables provide an important standard candle to measure distances to galaxies up to $\\sim 30$ Mpc. The Cepheid Distance Scale is considered to be among the most reliable methods because the physics of Cepheid pulsation is well-understood and the relation between the pulsation period and luminosity of the star is observationally well-established. The Cepheids are luminous, have a narrow range of surface temperatures; their pulsation is very stable and exhibit large amplitude. The intrinsic scatter in their \\plr\\ is believed to be only around $0.3$ mag. However, the Cepheid distance scale cannot be directly calibrated from the observation of nearby stars and consequently, several systematic effects still undermine its effectiveness as a standard primary candle to determine extragalactic distances beyond a few Mpc. Some of the questions which have direct bearing on the problem of distance calibration, but whose answers remain inconclusive in spite of extensive research, are listed below. \\bei \\item Are the preferential pulsation modes of Cepheids period-dependent? \\item Is a single \\plr\\ applicable to the entire instability strip? \\item Is the \\plr\\ modified appreciably due to metallicity dependency of the stellar structure? \\item Is the \\plc\\ relation a better indicator of distance than the \\plr\\ only? \\eni  \\noindent Iben and Tuggle (1975) numerically computed the period and luminosity of Cepheids for a range of masses and obtained a relation between metallicity, surface temperature, period and luminosity. The \\plc\\ relation is found to be dependent on metallicity due to extreme sensitivity of the color--temperature relation on chemical composition. However, according to Becker, Iben and Tuggle (1977), within the uncertainties, the relation between period and luminosity for the first and second crossings of the Cepheid instability strip by a particular star does not crucially depend on the chemical composition. Since the time spent in traversing the strip is largest for the second crossing, most of the observed Cepheids are in this stage of evolution. So, although conversion of the period into a $V$-magnitude will introduce small effects due to surface temperature and metallicity, a \\plr\\ derived from observations should not be affected by small variations in the chemical composition. Indeed, the robustness of the \\plr\\ against changes in the chemical composition is borne out by theoretical as well as observational studies. The theoretical models of Bressan \\ea (1993) produce the same period (within an error of $2 \\% $) for a given luminosity, irrespective of the chemical composition. From observations of Cepheids in M31, it appears that there is no significant dependence of the \\pl\\ zero point on metallicity gradients (\\cite{fm:90}). The recent review on the metallicity dependence of the Cepheid Distance Scale in the context of the HST Key Project on Extragalactic Distance Scale (\\cite{kenn:98}) also leads to the same conclusion. Similarly, even though color-color diagram of Cepheids can be used to determine their metallicity, it is not an improvement in terms of its application to the estimation of distances, given that after extinction correction the color has larger error than the $V$-magnitude (e.g., \\cite{fg:93}). Clearly, in order to address the above question, systematic work is required on the pulsation properties, evolution as well as  stellar atmospheric structure, taking into account the onset of convection in the atmosphere during the pulsation cycle.  In the present study we shall adopt the working hypothesis that the luminosity of the star {\\it as a function of period} is not directly altered by metallicity (subject to the star being a classical Cepheid), and that the color of the star provides a better diagnostic for the estimation of extinction than for the determination of the distance. We devise methods to determine the extinction by the interstellar medium, and particularly emphasize the importance of observations in multi-wavelength bands, and also address the question of pulsation modes of Cepheids.  The Cepheids in our own Milky Way galaxy have been observed by several astronomers over the years, and it is possible to obtain multi-wavelength data as well as accurate periods for a large number of them. A careful analysis of these Galactic Cepheids will naturally provide a useful template for identifying and estimating the various errors in the calibration of the Cepheid Distance Scale. A robust calibration of this distance scale is particularly important for extending it to extragalactic domains, as Cepheids are being observed in several far away galaxies, including those in the Virgo Cluster, by the Hubble Space Telescope. The hope is that distances based on these observations will ultimately lead to an accurate determination of the Hubble Constant. In this context, we attempt to provide a new calibration of the Cepheid Distance Scale, which is free from many of the systematic errors. In an accompanying communication (which we will refer to as Paper II), we apply these results for estimating the distance to the Virgo Cluster, based on the HST data for the Cepheids in the spiral M100.  This paper is organized as follows. In Section~\\ref{sec:number} we discuss the number distribution of Cepheids against their periods. In Section~\\ref{sec:ligcur}, we demonstrate the feasibility of obtaining accurate $\\vi $ \\lig s from the $V$ \\lig\\ of a Cepheid and limited number of observations in the I band. This method is particularly useful for analyzing Cepheid data with few observations in one band (as in HST observations). In Section~\\ref{sec:extcor}, we devise a formalism for extinction correction for each individual Cepheid, based on model atmospheres and an $\\rv$-dependent extinction law. Several useful period--color and amplitude--color relationships are also derived. Section~\\ref{sec:modes} is concerned about the different modes of pulsation of Cepheid variables and their manifestations in the observed properties like period and amplitude of pulsation. We argue that it is necessary to choose the correct lower cutoff period in the \\plr\\ in order to prevent contamination from multi-mode pulsators. In Section~\\ref{sec:mass}, we give an estimate for Cepheid masses at different periods, based on our results about the instability strip in the surface gravity versus effective temperature plane. Some discussion on the metallicity effects are presented in Section~\\ref{sec:metal} and the major limitations of the present work are listed in Section~\\ref{sec:limit}. The main conclusions from this work are summarized in Section~\\ref{sec:concl}. ",
        "conclusions": ""
    },
    "9803/hep-ph9803412_arXiv.txt": {
        "abstract": "We show that the decaying magnetohydrodynamic turbulence leads to a more rapid growth of the correlation length of a primordial magnetic field than that caused by the expansion of the Universe.  As an example, we consider the magnetic fields created during the electroweak phase transition.  The expansion of the universe alone would yield a correlation length at the present epoch of 1 AU, whereas we find that the correlation length is likely of order 100 AU, and cannot possibly be longer than $10^4$ AU for non-helical fields.  If the primordial field is strongly helical, the correlation length can be much larger, but we show that even in this case it cannot exceed 100 pc.  All these estimates make it hard to believe that the observed galactic magnetic fields can result from the amplification of seed fields generated at the electroweak phase transition by the standard galactic dynamo. ",
        "introduction": "Recently, considerable interest has been focused on the possibility that a primordial magnetic field may have been created at some early stage of the evolution of the Universe.  While the existence of a weak widespread extragalactic chaotic magnetic field cannot be ruled out, almost the only place where such a primordial field may have left an observable imprint is in galaxies, many of which possess microgauss, kiloparsec scale fields that are thought to be the result of dynamo amplification of a weak seed field \\cite{Kronberg}.  The idea that the seed for the galactic dynamo may be the field created in the very early Universe has inspired a number of works.  For example, such fields may have appeared at the electroweak phase transition \\cite{BaymMcLerran}. There are numerous other proposals involving physics at various scales (for a brief overview and further references, see \\cite{Olesen-rev}).  One feature shared by most particle-physics scenarios is the smallness of the correlation length of the magnetic field which results.  Indeed, at the moment of creation, the correlation length is limited by the horizon radius.  By pushing this moment to a very early stage in the history of the Universe, one makes the correlation length much smaller than it would be if the magnetic fields were created more recently, say, during proto-galactic contraction \\cite{Ostriker}.  For example, the fields generated at the electroweak phase transition, when the horizon radius is about 1 cm, would have a correlation length of at most about $10^{15}$ cm in the present Universe.  Contraction of proto-galaxies is likely to reduce this length scale by 2 orders of magnitude, which gives $10^{13}$ cm, or 1 AU, as the characteristic length scale of the seed field which the galactic dynamo is supposed to amplify.  Realistically, this length scale is likely much smaller, since the correlation length at creation is typically much less than the horizon size\\footnote{A magnetic field created during or before inflation \\cite{TurnerWidrow} may have a large correlation length, however the proposed inflationary models that could possibly generate a large-scale magnetic field may seem rather contrived \\cite{Ratra}.  We will not consider inflation-induced magnetic fields in this paper, but one should keep in mind this alternative.}.  Because any primordial magnetic fields generated at early times (via particle physics occurring at, say, the electroweak scale) have such small correlation lengths, these fields are not attractive candidates to play the role of seed fields for the galactic dynamo, unless the correlation length is somehow increased. The best developed theory of magnetic amplification in galaxies, the mean-field dynamo \\cite{Ruzmaikin}, operates under the assumption that the magnetic field is smooth on the scale of turbulent motion of the interstellar gas, which is of order 100 pc.  One may attempt to apply the mean-field dynamo theory for large-scale Fourier components of the chaotic magnetic field, neglecting the small-scale ones, however it is not obviously the correct procedure.  That the magnetic field at very small scales exponentiates rapidly is a well known fact that poses a serious problem for the mean-field dynamo theory \\cite{KurlsrudAnderson}.  In the situation when the seed field itself resides at small scales, this problem is likely to become more severe.  In this paper, we will not address the question of how a seed field is created, whether at the electroweak phase transition or during some other early epoch.  Our goal is to investigate the possibility that magnetohydrodynamic (MHD) effects can lead to a substantial increase in the length scale of the magnetic field at the present epoch.  We show that decaying MHD turbulence typically leads to a faster growth of the magnetic correlation length than one would expect from the expansion of the Universe alone.  We estimate that in the case of magnetic fields generated at the electroweak phase transition, the enhancement factor is $10^2$, and the correlation length may reach 100 AU.  If the primordial field has a large Chern-Simons number, the enhancement factor may be much larger, but the correlation length cannot possibly exceed 100 pc today if the magnetic field comes from the electroweak epoch.  Thus, although the magnetohydrodynamic effects we consider are of some help, they cannot increase the correlation length enough to make the electroweak generation of a primordial seed field a viable option.  We also consider generating the seed field at the QCD phase transition.  This may be a possibility, but is only viable if the bubble separation at the phase transition is very large.  This paper is organized as follows.  In Sec.\\ \\ref{sec:MHDeq} we write the basic equation governing the MHD of the Universe.  Sec.\\ \\ref{sec:non-hel} is devoted to the decay of non-helical MHD turbulence, whereas Sec.\\ \\ref{sec:hel} describes the scaling laws of the decay of helical turbulence.  Sec.\\ \\ref{sec:concl} contains concluding remarks.  The Appendix contains details about the EDQNM approximation used in our numerical simulation. ",
        "conclusions": "\\label{sec:concl}  In this paper we have considered the behavior of the correlation length of the primordial magnetic field, if such a field is generated in the early Universe.  We found the decay law of the MHD turbulence, which is responsible for a substantial increase of the correlation length.  We also observed a qualitative difference between the cases of non-helical and helical initial magnetic field.  Taking the case in which the magnetic field is generated during the electroweak phase transition as an example, we find that if the magnetic field is not helical the factor gained from MHD is about $10^2$ and today the field may be correlated at scales as large as 100 AU.  If by some chance the field is helical, the enhancement factor is much larger, but the final correlation length cannot exceed 100 pc, corresponding to 1 pc after proto-galactic collapse.  Let us note some uncertainties remaining in our estimation.  (In general, addressing these issues would lead to less optimistic values of the enhancement factor.)  First, it is not at all clear whether in MHD turbulence the magnetic energy reaches equipartition with the kinetic energy of bulk fluid motion.  While semi-analytical calculations (including our simulation) favor equipartition \\cite{Pouquet}, some numerical results indicate that the mean magnetic energy density remains small and concentrated at scales shorter than the largest turbulence scale \\cite{dns}.  If the latter remains true at very high Reynolds number, the magnetic field will be smaller after the electroweak phase transition and the correlation length at present will be smaller than we have estimated.  The second factor is neutrino diffusion.  After the electroweak phase transition, neutrinos are the particles with the longest mean free path. There is a certain time interval when neutrinos are still not decoupled from the physics at the turbulence scale, but their diffusion length is so large that the neutrino contribution to viscosity makes the Reynolds number to drop below 1.  During this time, there is no turbulence and the magnetic field is frozen.  After the neutrinos decouple from the fluid motion at the turbulence scale, magnetic stress leads to restoration of turbulence, and the magnetic length may continue to grow.  The overall effect is some reduction in the estimate of the final magnetic correlation length.  It would be nice to end this paper on a more positive note, and to this end we turn now to investigating the QCD phase transition, which occurs later than the electroweak transition and so may yield a longer correlation length.  If the QCD phase transition is first order, it will introduce fresh turbulence stronger than the decaying electroweak turbulence.  The amplitude and correlation length of the magnetic field will be determined by the QCD turbulence.  If fields were regenerated at the QCD transition on length scales of order of the horizon, the length scale it would have today is very large, since the horizon size at the QCD phase transition is much larger than at the electroweak epoch.  Even with no enhancement from MHD effects, the expansion of the universe yields a magnetic correlation length of order 1 pc. It can be estimated that the turbulence will survive till at least matter-radiation decoupling, at which the correlation length is is of order 1 kpc in the non-helical case and 100 kpc if the field is helical.  However, a magnetic field correlated on the horizon size is not expected to be produced at the QCD phase transition, and these estimates are too optimistic. Instead, the initial length scale must be of the order of the bubble spacing at the end of the phase transition (which is the natural scale of turbulence).  Bubble spacings larger than 100 cm are unlikely \\cite{Ignatius} and may have problem with standard big bang nucleosynthesis \\cite{bbn}.  A bubble spacing smaller than 100 cm again implies a very small correlation length at the present epoch.  If this constraint is respected, our analysis of the QCD case must end pessimistically, as in the electroweak case. However, one cannot completely rule out the possibility of a very large bubble spacing, which leads to a magnetic field correlated on a long enough length scale to be of interest.  A nonstandard evolution in which mixing after the phase transition occurs rapidly by hydrodynamic flows instead of slowly by diffusion or nonstandard nucleosynthesis may be required in this case.  Finally, although in this paper we presented the smallness of the magnetic correlation length as undesirable and tried to overcome it by invoking MHD turbulence, there may exist a non-standard dynamo mechanism where a small-scale seed gives rise to a large-scale magnetic field (see, e.g., \\cite{Chandran}).  This possibility is, however, outside the scope of the present paper."
    },
    "9803/astro-ph9803257_arXiv.txt": {
        "abstract": "We present a systematic search for OVI(1032\\AA,1037\\AA) absorption in a Keck HIRES spectrum of the $z=3.62$ quasar Q1422+231, with the goal of constraining the metallicity and ionization state of the low density intergalactic medium (IGM).  Comparison of CIV absorption measurements to models of the \\lya forest based on cosmological simulations shows that absorbers with $\\nh \\ga 10^{14.5}\\cdunits$ have a mean carbon abundance [C/H]~$\\approx -2.5$, assuming a metagalactic photoionizing background with the spectral shape predicted by Haardt \\& Madau (1996, HM).  In these models, lower column density absorption arises in lower density gas where most CIV is photoionized to CV.  Therefore, OVI should be the most sensitive tracer of metallicity in \\lya absorbers with $\\nh \\la 10^{14.5}\\cdunits$.  OVI lines lie at wavelengths heavily contaminated by Lyman series absorption, so we interpret the search results by comparing to carefully constructed, mock Q1422 spectra drawn from a hydrodynamic simulation of a $\\Lambda$-dominated cold dark matter model.  A search for deep, narrow absorption features yields only a few candidate OVI lines in the spectrum of Q1422.  HI absorption blankets the position of the doublet companion line in each case, and the total number of narrow lines is statistically consistent with that in zero-metallicity artificial spectra.  Artificial spectra generated with the HM background and [O/H]~$\\ga -2.5$ predict too many narrow lines and are statistically inconsistent with the data.  We also search for OVI associated with CIV systems, using the optical depth ratio technique of Songaila (1998).  With this method we {\\it do} find significant OVI absorption; matching the data requires [O/C]~$\\approx +0.5$ and corresponding [O/H]~$\\approx -2.0$.  Taken together, the narrow line and optical depth ratio results imply that (a) the metallicity in the low density regions of the IGM is at least a factor of three below that in the overdense regions where CIV absorption is detectable, and (b) oxygen is overabundant in the CIV regions, consistent with the predictions of Type II supernova enrichment models and the observed abundance pattern in old halo stars.  The photoionizing background spectrum would be truncated above 4~Ry in regions that have not undergone helium reionization (HeII$\\longrightarrow$HeIII), and in this case matching the Q1422 data requires lower [C/H] but higher [O/H]. Taking [O/C]$\\approx +1$ as the maximum plausible overabundance of oxygen, we conclude that helium must have been reionized through at least 50\\% of the volume from $z \\sim 3 - 3.6$. ",
        "introduction": "The Lyman alpha (Ly$\\alpha$) ``forest\" (\\cite{lyn71}; \\cite{sar80}) of spectral features caused by HI absorption along the line of sight to a quasar probes the state of the intergalactic medium over a wide range of physical conditions.  During the past few years, high-precision observations made using the HIRES spectrograph (\\cite{vog94}) on the 10m Keck telescope have quantified the statistics of these low column density absorbers to unprecedented accuracy (e.g., \\cite{hu95}; \\cite{lu96b}; \\cite{kim97}; \\cite{kir97}).  During the same time span, cosmological simulations that incorporate gas dynamics, radiative cooling, and photoionization have been able to reproduce many of the observed properties of quasar absorption spectra (\\cite{cen94}; \\cite{zha95}; \\cite{her96}; \\cite{mir96}; \\cite{dav97a}).  The rapid progress on theoretical and observational fronts has led to the emergence of a new paradigm for the origin of the high-redshift ($z\\ga 2$) \\lya forest, in which most \\lya forest lines are produced by regions of low to moderate overdensity in hierarchically collapsing structures that are not in dynamical or thermal equilibrium.  \\lya lines of lower column density generally arise in gas of lower physical density, which has a lower neutral hydrogen fraction because of the reduced recombination rate.  In this paper, we use a matched comparison between a cosmological hydrodynamic simulation and a Keck HIRES spectrum of the quasar Q1422+231 to constrain the metal abundance of this low density gas.  The recent detection of metal lines associated with \\lya forest absorbers having column densities $\\nh \\la 10^{15} \\cdunits$ (\\cite{cow95}; \\cite{tyt95}; Songaila \\& Cowie 1996, hereafter \\cite{son96}) has provided a new avenue for investigating the ionization state and enrichment history of the high-redshift intergalactic medium (IGM).  \\cite{son96} showed that 75\\% of \\lya absorbers with $\\nh > 10^{14.5} \\cdunits$ have associated CIV absorption.  Using simple photoionization models they estimate that the mean metallicity of these absorbers is between\\footnote{We use the standard notation of brackets to denote the relative abundance, in logarithm, versus solar.} [C/H]$\\sim -2$ and $-3$. Studies that use cosmological simulations to model the density and temperature of the absorbing gas obtain a better match to the data for a mean metallicity of [C/H]$\\sim -2.5$ with around one dex of scatter, assuming either a power-law ionizing background (\\cite{hae96}) or a reprocessed quasar ionizing background (Hellsten \\etal 1997, hereafter \\cite{hel97}).  The question of how metals came to reside in these intermediate-density \\lya absorbers remains unanswered.  Since \\lya absorbers with $\\nh \\la 10^{15}\\cdunits$ are optically thin and associated with density peaks of overdensity $\\la 20$, it is unlikely that they contain star-forming regions that produce {\\it in situ} enrichment.  Thus some transport mechanism must be invoked to explain the presence of metals in these regions.  One possibility is that the metals are ejected from nearby, forming galaxies, by some combination of supernova blowout (\\cite{mir97}) and tidal stripping (\\cite{gne97}; \\cite{gne98}). Since the enriching proto-galaxies are likely to form more efficiently in high density environments, these scenarios predict a significant correlation between metallicity and density (see, e.g., figure~6 of Gnedin \\& Ostriker 1997). Alternatively, the IGM may have been enriched at a very early epoch by more ubiquitous Population III objects (e.g., \\cite{hai97}), in which case all \\lya absorbers might be expected to have roughly the same metallicity. Because CIV lines probe only a small range of HI column densities (and hence physical densities) with adequate statistics, it is difficult to distinguish between these enrichment models using only CIV data. Metal abundances can be studied at higher densities in Lyman limit systems ($\\nh \\ga 10^{17}\\cdunits$) and damped \\lya systems ($\\nh \\ga 10^{20}\\cdunits$), but these are objects where {\\it in situ} enrichment is likely and where (especially for Lyman limit systems) uncertain radiative transfer effects complicate the inference of metal abundances from line strengths.   In this paper, we attempt to extend metallicity constraints to the low column density ($\\nh \\la 10^{14.5}\\cdunits$), and hence low physical density, \\lya forest. Carbon is difficult to detect in this regime because the low density reduces $N_{C}$ and ionizes more CIV to CV. However, \\cite{lu98} have used composite spectra to attempt to detect CIV in such systems, and we briefly discuss their results in \\S\\ref{sec: disc}.  In this paper, we focus instead on OVI, which is expected to be the one detectable metal absorption feature tracing \\lya absorbers with $\\nh \\la 10^{14.5}\\cdunits$ because of its high ionization state and large oscillator strength (Hellsten \\etal 1998, hereafter \\cite{hel98}). The difficulty with this approach, and the reason that it has not been previously attempted, is that the OVI absorption features lie embedded within the \\lya forest, which is quite crowded at these redshifts.  We overcome this problem by using a line identification scheme specifically designed to select candidate OVI features and by using artificial spectra extracted from a realistic cosmological simulation to calibrate the efficiency of OVI detection and the contamination from narrow \\lya lines. With these procedures, we can test whether the low density regions of the \\lya forest are consistent with a uniform metallicity extrapolated from the CIV data at higher densities.  The low density IGM is highly photoionized by the metagalactic ultraviolet (UV) background.  For our standard spectral shape, we assume that the UV background is produced by quasar emission reprocessed by \\lya forest absorption (Haardt \\& Madau 1996, hereafter \\cite{haa96}). We find that if the IGM metallicity is uniform at [C/H]$\\sim -2.5$, the UV background has the spectral shape given by \\cite{haa96}, and oxygen has the factor of three overabundance (relative to solar) predicted by Type II supernova enrichment models, then OVI should be readily detectable in the spectrum of Q1422+231. However, our detection algorithm finds very few candidate OVI lines in the spectrum.  The absence of detectable OVI features has several possible interpretations: (1) oxygen is not overabundant relative to carbon in the low density IGM, (2) the metallicity of low density regions as traced by $\\nh \\la 10^{14.5}\\cdunits$ \\lya absorption systems is lower than the metallicity of intermediate-density regions traced by higher column density absorption, or (3) there are many fewer high-energy photons capable of photoionizing OV to OVI than are predicted by the \\cite{haa96} ionizing background.  To distinguish between these interpretations, we apply a second algorithm, the optical depth ratio technique of Songaila (1998, hereafter \\cite{son98}) designed to detect OVI in regions where significant CIV absorption is found.  Since we have independently determined [C/H] in these regions, we can use this technique to discriminate between different ionization conditions and abundance patterns.  By applying this technique to the spectrum of Q1422+231 and calibrating the results using artificial spectra, we find that: (1) a significant metallicity gradient (declining metallicity with declining density) must exist regardless of whether helium has mostly reionized or not, (2) if helium has reionized by $z\\sim 3.6$, our results are consistent with [O/C]~$\\approx +0.5$, and (3) if the epoch of helium reionization does not begin until $z\\sim 3$, a highly implausible oxygen overabundance of [O/C]~$\\ga +2$ is required.  Combining the results from these two analysis techniques, our favored scenario for the low density IGM at $z\\ga 3$ is one in which more than half of the volume of the universe has helium predominantly reionized by $z\\sim 3.6$, oxygen is overabundant relative to carbon by a factor $\\ga 3$, and spatial regions with overdensities $\\sim 10$ have an average metallicity {\\it at least} a factor of 3 higher than regions near the mean baryonic density.  Section~\\ref{sec: modeling} describes our cosmological simulation, the Q1422+231 data, and our procedure for constructing artificial absorption spectra.  Section~\\ref{sec: OVIsearch} discusses previous OVI searches and reviews \\cite{hel98}'s argument that OVI should be the most effective tracer of metallicity in \\lya absorbers with $\\nh \\la 10^{14}\\cdunits$.  Section~\\ref{sec: civ} examines CIV absorption at intermediate column densities, repeating the general arguments of \\cite{hel97} and Rauch et al.\\ (1997a) but with a much closer match between the theoretical and observational analysis procedures. Section~\\ref{sec: search} is the heart of the paper.  It describes our algorithm for identifying candidate OVI lines, presents the results of the OVI searches in the real and artificial spectra, and discusses the properties of candidate OVI absorbers in Q1422+231 and the simulations. Figures~\\ref{fig: auto5}--\\ref{fig: OVIsel} demonstrate the paper's central results.  Section~\\ref{sec: assumptions} discusses the impact of varying the assumptions of our standard theoretical model. Section~\\ref{sec: pixmet} describes the optical depth ratio technique for quantifying OVI absorption, shows the results from this technique applied to Q1422+231 and artificial spectra, and discusses these results in conjunction with the results from Section~\\ref{sec: search}. Section~\\ref{sec: disc} summarizes our conclusions and discusses them in light of other recent observational and theoretical developments. ",
        "conclusions": "\\label{sec: disc}  We present a systematic search for OVI absorption in the spectrum of Q1422+231 ($z=3.62$), using a narrow line detection algorithm proven effective at identifying OVI absorption in artificial spectra, and an optical depth ratio technique introduced by \\cite{son98}.  The first technique traces OVI predominantly in systems with $10^{13.5}\\la \\nh\\la 10^{15}\\cdunits$, whereas the second technique traces only OVI associated with CIV absorption, \\ie in $10^{14.5}\\la \\nh\\la 10^{16}\\cdunits$ systems.  By comparing Q1422 and artificial spectra having varying metallicities, we determine that \\begin{enumerate}  \\item{[O/H] must be lower in lower density regions, for either an \\cite{haa96} ionizing background or an ionizing background significantly truncated above 4~Ry.  If [O/C] is constant in systems up to $\\nh\\sim 10^{16}\\cdunits$, our results imply that regions traced by $\\nh\\la 10^{14}\\cdunits$ systems (corresponding to gas at roughly the mean baryonic density) have a mean metallicity lower by at least a factor of 3 compared to regions traced by $\\nh\\sim 10^{15}\\cdunits$ systems (corresponding to a baryonic overdensity of $\\sim 10$).}  \\item{More than half the universe must have helium reionized by $z\\sim 3$. If helium has completely reionized by $z\\sim 3.6$ (the highest redshift probed by the Q1422 data), then our analysis implies [O/C]~$\\approx +0.5$, in good agreement with overabundance measurements of Type II supernovae enriched systems.  If a significant portion of the universe has not reionized helium by $z\\sim 3$ and therefore has a softer ionizing background spectrum, then the required oxygen overabundance is higher. For example, if half of the volume has not reionized helium, then [O/C]~$\\approx +1.2$, already greater than the observed overabundance of any class of Type II supernovae enriched systems. If the spectrum were soft throughout the universe at $z\\ga 3$ then an implausibly high overabundance, [O/C]~$\\approx +2.3$, would be required. } \\end{enumerate}  These conclusions are in good agreement with the recent study of \\cite{lu98}, who used composite spectra to investigate CIV absorption in systems with $10^{13.5}< \\nh < 10^{14} \\cdunits$ and found that the metallicity of these absorbers must be [C/H]~$\\la -3.5$. Cosmological simulations show that $\\nh \\sim 10^{14} \\cdunits$ roughly corresponds to the dividing line between overdense and underdense regions of the universe (though the value of $\\nh$ that marks this division depends on redshift and, to a lesser extent, on cosmological parameters).  Our results therefore imply that mildly overdense regions such as filaments and sheets have been enriched, while underdense regions are virtually chemically pristine.  Simulations that self-consistently enrich the IGM by tracking metal production and transport find that a strong metallicity gradient is predicted between the mildly overdense and underdense regions (see figure~3 in \\cite{gne98}); this predicted gradient is in good agreement with the \\cite{lu98} data and with the scenario we present above.  Recent measurements of a jump in the SiIV/CIV ratio around $z\\sim 3$ (\\cite{son98}) may be difficult to reconcile with conclusion~(2) above.  While we have yet to conduct a systematic comparison of SiIV in observed and artificial spectra, primarily because of the small numbers of SiIV systems detectable in our one available quasar spectrum, we expect our results will be in agreement with \\cite{son98}, who argues that such a jump requires a much softer ionizing background at $z\\ga 3$.  One way to reconcile these results may be to invoke patchy helium reionization at that epoch, as suggested by \\cite{rei97}; such a model may be tested in greater detail by searching for an anti-correlation between OVI and SiIV detections.  If helium reionization occurs around $z\\sim 3$, our narrow line algorithm should yield many more OVI detections at redshifts $z\\la 3$.  Such searches are difficult because of the poor blue sensitivity of the HIRES spectrograph (and complete loss of sensitivity at $\\lambda\\la 3800$\\AA), but quasar spectra do exist that could provide constraints down to $z\\sim 2.7$.  The presence of a substantial number of OVI lines in this regime would strongly favor the late helium reionization scenario; there is already some weak evidence that OVI is more abundant at $z\\la 3$ (\\cite{son98}).  If OVI features continue to be virtually undetectable down to $z\\sim 2.7$, this would be compelling evidence against the late helium reionization scenario, since the HeII absorption measurements of Davidsen, Kriss, \\& Zheng (1996) imply that helium has been reionized by this redshift.  We hope to work with observers to attempt this search in the near future.  In a broader context, our work illustrates the power of combining cosmological hydrodynamic simulations of structure formation with high-quality quasar spectra to infer the ionization state and the enrichment history of the high-redshift IGM.  Future observations and simulations promise a wealth of information, which, when combined, will help us to better understand the evolution of the IGM and its connection to early star formation and the epoch of primeval galaxies."
    },
    "9803/gr-qc9803026_arXiv.txt": {
        "abstract": "Clues as to the geometry of the universe are encoded in the cosmic background radiation. Hot and cold spots in the primordial radiation may be randomly distributed in an infinite universe while in a universe with compact topology distinctive patterns can be generated. With improved vision, we could actually see if the universe is wrapped into a hexagonal prism or a hyperbolic horn.  We discuss the search for such geometric patterns in predictive maps of the microwave sky. ",
        "introduction": "\\label{compf}  We want  to predict a map of the temperature fluctuations. In a homogeneous and isotropic space, an angular average over the fluctuations contains all of the essential information. Of the six compact, orientable flat spaces, all destroy global isotropy and all except for the hypertorus destroy global homogeneity. As a result, there is more information in a map of temperature fluctuations than just the angular power spectrum. Although we argue that the angular average overlooks conspicuous features in general, for the equal sided flat cases angular spectra do provide a reasonable bound.     Four of the six orientable, compact topologies of $\\e$ are constructed from a parallelepiped as the  fundamental domain.  The other two are built from a hexagonal prism. The hypertorus is the simplest and has been studied by many authors \\cite{{flat},{moreflat},{sss},{add}}. Stevens, Scott, and Silk \\cite{sss} pointed out that in a flat $3$-torus, the spectrum of temperature fluctuations was truncated at long wavelengths in order to fit within the finite box. Contrary to standard lore, we find all of the equal sided compact flat manifolds show a truncation in the power of fluctuations on wavelengths comparable to the size of the fundamental domain \\cite{{lss},{tarun}}. The longest wavelength fluctuation observed, namely the quadrupole, is in fact low. Some might even take this as evidence for topology \\cite{workshop}. Cosmic variance is also large on large scales. Consequently, a fundamental domain the size of the observable universe is actually consistent with the COBE data \\cite{lss}.    A very small universe however is incompatible with the data. The cutoff in long wavelength perturbations is accompanied by gaps in power at wavelengths that do not correspond to integer windings through the fundamental domain. {\\it All} compact spaces show discrete harmonics and as such the sharp harmonics may be a more generic sign of compact topology. The jaggy spectra of such small compact flat spaces are tens of times less likely than the smooth spectrum of infinite $\\e$.  We conclude, quite conservatively, that the universe, if finite and flat and equal-sided, must be at least $80$\\% the radius of the surface of last scatter and so $40$\\% of the diameter of the observable universe. There could still be as many as eight copies of our universe within the observable horizon.    \\begin{figure} \\centerline{{\\psfig{file=hexglue.eps,width=3in}}} \\centerline{{\\psfig{file=hex0_1_.1.ps,width=3in}}} \\caption{A hexagonal prism with a $2\\pi/3$ twist.  The observer is at the center of the universe in the map of $\\delta T(\\hat n)/T$. The fundamental domain is half the diameter of the observable universe in two directions and one-tenth that in the twisted direction. \\label{fighex}} \\end{figure}   If instead of an equal-sided space we consider a fundamental domain with disparate length scales, the angular power spectrum is in general a poor discriminant. The averaging over the sky fails to recognize the strong features in the cosmos. Fig.\\ \\ref{fighex} shows a predictive map of the hot and cold fluctuations in a $2\\pi/3$-twisted hexagonal prism. We have set the length of the fundamental domain to be ten times smaller in the twisted direction than along the face of the hexagon. The average large angle power in fluctuations is actually consistent with the data, although clearly this anisotropic space does not look like the sky we observe. A better statistic to discern patterns and correlations is badly needed \\cite{{lss},{lssb}}. The promising suggestions of \\cite{{bps},{css},{lssb}} may be the key and are discussed more in \\S \\ref{comph}.   \\begin{figure} \\centerline{{\\psfig{file=tile2.eps,width=1.5in}}} \\centerline{{\\psfig{file=hexmodes.eps,width=1.75in}}} \\caption{Top: A guess at one mode.  Bottom: A contour plot of the temperature fluctuation for a similar mode. \\label{onemode}} \\end{figure}    We could have predicted certain features of the map of Fig.\\ \\ref{fighex}, even if we had not known the eigenmodes explicitly. A 2D slice through the 3D tiling of space is represented in Fig.\\ \\ref{onemode}.  If we draw bands connecting opposite sides of the hexagons and highlight any overlaps, we can predict the imprint of one mode as shown on the top of Fig. \\ref{onemode}. Given that we do know the eigenmodes, we can show the actual contour plot of the hot and cold fluctuations for a similar mode on the bottom of Fig.\\ \\ref{onemode}.  Comparing the guess with the actual contours shows our guess did quite well. The hexagonal shape of the universe is clearly seen. In actuality, there are many modes competing to imprint a pattern on the sky which blurs the signature hexagons. In Fig.\\ \\ref{mode2}, is another contour plot of hot and cold spots for a different mode which exhibits the $2\\pi/3$  twist through the prism.  \\begin{figure} \\centerline{{\\psfig{file=hexmodes2.eps,width=1.75in}}} \\caption{A contour plot of the temperature fluctuation for a mode that winds through the twisted prism. \\label{mode2}} \\end{figure}  The competition between fluctuations obscurs some features while enhancing others. The surface of last scatter cuts a sphere out of the full 3D space, an elliptic projection of which is given in the map of Fig. \\ref{fighex}. Can you see hexagons in the map? Almost. Is the $2\\pi/3$ twist in the space visible? We are currently developing ways of looking at the sky that pull the underlying patterns out of the noise \\cite{lssb}. ",
        "conclusions": ""
    },
    "9803/astro-ph9803311_arXiv.txt": {
        "abstract": "We present the results of deep spectropolarimetry of two powerful radio galaxies at $z\\sim2.5$ (4C 00.54 and 4C 23.56) obtained with the W.M. Keck II 10m telescope, aimed at studying the relative contribution of the stellar and non-stellar components to the ultraviolet continuum. Both galaxies show strong linear polarization of the continuum between rest-frame $\\sim$1300-2000~\\AA, and the orientation of the electric vector is perpendicular to the main axis of the UV continuum. In this sense, our objects are like most 3C radio galaxies at $z\\sim1$. The total flux spectra of 4C 00.54 and 4C 23.56 do not show the strong P-Cygni absorption features or the photospheric absorption lines expected when the UV continuum is dominated by young and massive stars. The only features detected can be ascribed to interstellar absorptions by SiII, CII and OI. Our results are similar to those for 3C radio galaxies at lower $z$, suggesting that the UV continuum of powerful radio galaxies at $z\\sim2.5$ is still dominated by non-stellar radiation, and that young massive stars do not contribute more than $\\approx$50\\% to the total continuum flux at 1500~\\AA. ",
        "introduction": "High-$z$ radio galaxies (H$z$RGs) are observable to very high redshifts and can be used to study the formation and evolution of massive elliptical galaxies (see McCarthy 1993 for a review). One of the most controversial issues is the physical cause of the alignment between the radio source and UV continuum axes of the H$z$RGs (the so called `alignment effect', Chambers, Miley \\& van Breugel 1987, McCarthy et al. 1987). Two main competing scenarios have been proposed. The first is star formation induced by the propagation of the radio source through the ambient gas (see McCarthy 1993 and references therein); the second  explains the alignment effect as the result of a hidden quasar whose radiation is emitted anisotropically and scattered towards the observer, producing strong linear polarization perpendicular to the radio-UV axis (Tadhunter et al. 1988; di Serego Alighieri et al. 1989). The latter scenario is closely related to the unification of powerful radio-loud AGN, and provides a way of testing it directly (see Antonucci 1993 and references therein). After the first detections of strong UV polarization in H$z$RGs obtained with 4m-class telescopes (di Serego Alighieri et al. 1989; Jannuzi \\& Elston 1991; Tadhunter et al. 1992; Cimatti et al. 1993), recent observations made with the Keck I 10m telescope have demonstrated the presence of spatially extended UV continuum polarization and of hidden quasar nuclei in some of the 3C radio galaxies at $0.7<z<1.8$, favoring the beaming and scattering scenario (Cohen et al. 1996; Cimatti et al. 1996,1997; Dey et al. 1996; Tran et al. 1998). On the other hand, Dey et al. (1997) have recently shown that the UV continuum of 4C 41.17 ($z=3.8$) is unpolarized and consistent with that of a typical starburst galaxy. The most stringent comparison between the starburst and the scattering scenarios can be performed at $\\lambda_{rest} \\sim 1000-2000$~\\AA, where most of the strongest spectral features of O and B stars are located. This spectral window can be covered from the ground by observing radio galaxies at $z>2$. We have started a program of observations of these galaxies using spectropolarimetry at the Keck II 10m telescope, and in this Letter we report on the first two objects we have studied, concentrating on their continuum and absorption line properties. Throughout this paper we assume $H_0=50$ kms$^{-1}$ Mpc$^{-1}$ and $q_0=0$. ",
        "conclusions": "Our observations suggest that the UV spectra of 4C 23.56 and 4C 00.54 are not dominated by young massive stars, whereas the strong perpendicular polarization indicates the presence of a relevant scattered continuum, making 4C 23.56 and 4C 00.54 similar to the polarized 3C radio galaxies at 0.7$<z<$2. Adopting the prescriptions of Dickson et al. (1995) and Manzini \\& di Serego Alighieri (1996) and assuming the average HeII$\\lambda$1640/H$\\beta$ ratio (3.18) observed in radio galaxies (McCarthy 1993), we estimate that the nebular continuum contributes only $\\sim$8\\% and $\\sim$13\\% to the total flux at 1500~\\AA~ for 4C 23.56$a$ and 4C 00.54 respectively.  If we assume that all H$z$RGs have an obscured quasar nucleus which feeds the powerful radio source and whose light is scattered by dust and/or electrons, we can interpret the low (or null) polarization of 4C 41.17 (Dey et al. 1997) as due to dilution of the scattered radiation by the unpolarized light of young stars. A limit on the amount of stellar light in the UV continuum of 4C 23.56$a$ and 4C 00.54 can be derived by assuming that the observed polarization is diluted by the unpolarized stellar and nebular continua. The ratio between the stellar light and the total flux at 1500~\\AA~ can be roughly estimated as $F_{stars}/F_{total}=\\{[1-(P_{obs}/P_0)]-\\kappa\\}$, where $P_{obs}$ and $P_0$ are the observed and intrinsic degree of polarization and $\\kappa$ is the ratio between the nebular and the total continuum. Adopting a half-cone opening angle of 45$^{\\circ}$ and an angle of 90$^{\\circ}$ between the cone axis and the line of sight, we derive $P_0\\sim$30\\% and $P_0\\sim$50\\% for dust (Manzini \\& di Serego Alighieri 1996) and electron scattering (Miller, Goodrich \\& Mathews 1991) respectively. Thus, adopting $P_{obs}$(1500~\\AA)=13.1\\% and 14.4\\% for 4C 00.54 and 4c 23.56$a$ respectively, for dust scattering we obtain that $F_{stars}/F_{total} \\leq$43\\% and $\\leq$44\\% for 4C 00.54 and 4C 23.56$a$ respectively, whereas $F_{stars}/ F_{total}$ increases to $\\leq$61\\% and $\\leq$63\\% for electron scattering. These ratios can be considered upper limits because we do not know if the observed scattered light is really diluted by a stellar continuum, and they imply that stellar light cannot contribute more than about half of the UV continuum at 1500~\\AA.  The properties of 4C 23.56, 4C 00.54 and 4C 41.17 can be interpreted in a evolutionary scenario where H$z$RGs at $z>3$ have a major episode of star formation, and their AGN scattered component is diluted by the stellar light, but it becomes observable at lower $z$ when the starburst ceases. However, given the rapid evolution of the UV light from a starburst, it is also possible that 4C 41.17 simply represents a case dominated by the starburst rather than an evolutionary sequence. Future observation of a complete sample of H$z$RGs will help us to understand the nature of the alignment effect and the evolution of the host galaxies of powerful radio sources."
    },
    "9803/astro-ph9803127_arXiv.txt": {
        "abstract": "We present high-quality HST/GHRS spectra in the Hydrogen L$\\alpha$ spectral region of Vega and Sirius-A. Thanks to the signal-to-noise ratio achieved in these observations and to the similarity of the two spectra, we found clear evidence of emission features in the low flux region, $\\lambda\\lambda$1190-1222\\,\\AA. These emission lines can be attributed unambiguously to \\ion{Fe}{ii} and \\ion{Cr}{ii} transitions. In this spectral range, silicon lines are observed in absorption.  We built a series of non-LTE model atmospheres with different, prescribed temperature stratification in the upper atmosphere and treating \\ion{Fe}{ii} with various degrees of sophistication in non-LTE. Emission lines are produced by the combined effect of the Schuster mechanism and radiative interlocking, and can be explained without the presence of a chromosphere. Silicon absorption lines and the L$\\alpha$ profile set constraints on the presence of a chromosphere, excluding a strong temperature rise in layers deeper than $\\tau_{\\rm R} \\approx 10^{-4}$. ",
        "introduction": "On the main sequence, A-type stars are at a juncture point between hot and cool stars. While hot, massive stars undergo strong mass loss in fast winds ($\\dot{M}\\ge 10^{-9}$\\,M$_\\odot$/yr), cool stars show chromospheric activity connected to their subsurface convective layers. Both phenomena apparently disappear or become much weaker at spectral type A. Many studies have thus been devoted to the outer layers of A-type stars to search for indications of a wind or of stellar activity. Several attempts to detect signatures of weak winds in main-sequence A stars have been unsuccessful (e.g. Lanz~\\& Catala \\cite{lanz92}). Recently, however, a quite weak, blue-shifted absorption was detected in the \\ion{Mg}{ii} resonance lines of Sirius, and interpreted as a wind signature (Bertin et~al. \\cite{sirius3}). A mass loss rate of $\\dot{M}\\approx 10^{-12}$\\,M$_\\odot$/yr was derived, consistent with the idea that A-type star winds are radiatively-driven like the winds of hotter stars. On the cool side, a limit to chromospheric activity has been set at A7 (B\\\"ohm-Vitense~\\& Dettmann \\cite{bohmvitense80}, Marilli et~al. \\cite{marilli97}, Simon~\\& Landsman \\cite{simon97}). The most common diagnostics of chromospheres and winds are emission features. Therefore, we are not expecting emission lines in A stars, except cases where such lines arise from the circumstellar environment.  High-quality ultraviolet spectra have become available with the {\\em Goddard High Resolution Spectrograph} (GHRS) aboard the {\\em Hubble Space Telescope} (HST). Even the core of strong resonance lines, including \\ion{H}{i} L$\\alpha$, can be observed with a reasonably good signal-to-noise ratio. This makes it possible to investigate in greater detail the line profile of strong resonance lines. They are the best tool to probe the outer layers of stars, being indeed formed very high in the atmosphere. In this respect, L$\\alpha$ is most interesting because it spans the largest range of depth of formation, from the far wing to the line core. This large variation in opacity also affects the formation of lines of other elements, especially close to L$\\alpha$ core. Such lines see a much lower local pseudo-continuum than lines outside L$\\alpha$, and will be formed much higher in the atmosphere than weak lines in other regions.   \\begin{table*} \\caption[]{Observation log.} \\label{TabObs} \\begin{tabular}{lccccr} \\hline &&&&& \\\\ [-3mm] Target & Spectral Range & Date of Observation & GHRS Grating & Aperture & Exposure time \\\\ \\hline &&&&& \\\\ [-3mm] Sirius-A & 1188 \\AA\\ -- 1218 \\AA & 1996 Nov 20 & G140M & SSA & 1632.0 s \\\\ Sirius-A & 1278 \\AA\\ -- 1307 \\AA & 1996 Nov 20 & G140M & SSA & 217.6 s \\\\ Sirius-A & 1308 \\AA\\ -- 1337 \\AA & 1996 Nov 20 & G140M & SSA & 217.6 s \\\\ Vega & 1185 \\AA\\ -- 1222 \\AA & 1996 Dec 23 & G160M & SSA & 435.2 s \\\\ Vega & 1274 \\AA\\ -- 1311 \\AA & 1996 Dec 23 & G160M & SSA & 108.8 s \\\\ Vega & 1303 \\AA\\ -- 1341 \\AA & 1996 Dec 23 & G160M & SSA & 108.8 s \\\\ \\hline \\end{tabular} \\end{table*}  In Sect. 2 and 3, we will describe our GHRS observations around L$\\alpha$ of two bright A stars, Vega and Sirius-A. We will in particular point out the presence of \\ion{Fe}{ii} and \\ion{Cr}{ii} emission lines between 1190 and 1222\\,\\AA. Bertin et~al. (\\cite{sirius2}) noticed the presence of emission features around L$\\alpha$ in a Cycle~1 GHRS spectrum of Sirius-A, originally recorded to derive the D/H abundance ratio in the local interstellar medium. This prompted us to repeat and extend these observations to investigate their origin. An explanation of these emission features is given in the second half of the paper. In Sect.~4, we describe our new non-LTE model atmospheres. We explore and set limits on a chromosphere (Sect.~5), and investigate non-LTE effects in \\ion{Fe}{ii} line formation (Sect.~6). ",
        "conclusions": "We have reported emission features in the L$\\alpha$ profile of Vega and Sirius-A. These emission lines have been attributed to \\ion{Fe}{ii} and \\ion{Cr}{ii} transitions. The identification appears quite secure because all the lines of several multiplets appear in emission.  We have built non-LTE model atmospheres with different assumed temperature structures in the outer layers and incorporating \\ion{Fe}{ii} with different degrees of sophistication. We found that the emission features cannot be explained by a chromospheric temperature rise. To produce the observed \\ion{Fe}{ii} emissions, the temperature would have to increase in relatively deep layers, turning other lines into emission (e.g. \\ion{Si}{ii}, \\ion{Si}{iii} lines). However, we cannot exclude a chromospheric rise in shallower layers ($\\tau_{\\rm R} \\le 10^{-4}$) based on our present observations, in agreement with earlier results (Freire-Ferrero et~al. \\cite{freire83}).  Non-LTE \\ion{Fe}{ii} line formation calculations with different model atoms have demonstrated that some \\ion{Fe}{ii} lines can turn into emission in the wavelength range between 1190 and 1240\\,\\AA. We stress that emission lines are predicted {\\em only} in this very low flux, central region of L$\\alpha$. This results from the combined effect of the Schuster mechanism and radiative interlocking. Some highly-excited levels are overpopulated by transitions occurring in a high-flux region, and preferentially de-excite in this region near L$\\alpha$. This mechanism explains the similarity of Vega and Sirius spectra.  Differences between the two stars can also be understood with this mechanism. The higher heavy-element content in Sirius' photosphere results in depressing the flux, in particular near the flux maximum. The efficiency of the pumping is thus reduced, yielding generally weaker emissions in Sirius than in Vega. The details depend on the exact wavelength of the pumping transitions. The flatter L$\\alpha$ profile in Vega is also a consequence of the different metallicity. Lyman continuum heating must be more efficient in Vega's case (less heavy-element line opacity) yielding a somewhat higher temperature in the outer layers and a higher flux in the central region of L$\\alpha$.  We believe that the origin of the \\ion{Cr}{ii} emission lines may be explained by similar mechanisms. While we cannot rule out that radiative interlocking is also effective in \\ion{Cr}{ii}, the difference between Vega and Sirius points to the Schuster mechanism being the major cause of the emission in this case. The lines are thus stronger in Sirius due to the larger chromium abundance, and are simply too weak in Vega to stand out of the noise.  Although our model atmosphere calculations provide an explanation to an unexpected observation of emission lines in the spectrum of early A-type stars, we did not achieve a good fit to the L$\\alpha$ profile at this stage. It seems however likely that the flux observed in the central region of L$\\alpha$ may be explained by increasing somewhat the fraction of non-coherent scattering in the PRD approximation that we have used. Matching these observations would require (at least) non-LTE line-blanketed model atmospheres, treatment of L$\\alpha$ in partial redistribution tuning the ratio between coherent and non-coherent scattering, and improved, non-hydrogenic \\ion{Fe}{ii} photoionization cross-sections. Such an approach is necessary to gain a deeper insight into the outer layers of Vega and Sirius, and this certainly deserves further study.  Finally, we did not find emission lines very close to the L$\\alpha$ core, especially in the 0.5\\,\\AA\\  blueward of the central wavelength. This implies fortunately that we have so far no reason to question the previous results on the local interstellar cloud (Bertin et~al. \\cite{sirius2}), and on the wind absorption feature (Bertin et~al. \\cite{sirius3})."
    },
    "9803/astro-ph9803037_arXiv.txt": {
        "abstract": "We have preliminary results on the parallelization of a Tree-Code for evaluating  gravitational forces in N-body astrophysical systems. For our T3D CRAFT implementation, we have obtained an encouraging speed-up behavior, which reaches a value of 37 with 64 processor elements (PEs). According to the Amdahl'law, this means that about 99\\% of the code is actually parallelized. The speed-up tests regarded the evaluation of the forces among $N = 130,369$ particles distributed scaling the actual distribution of a sample of galaxies seen in the Northern sky hemisphere. Parallelization of the time integration of the trajectories, which has not yet been taken into account, is both easier to implement and not as fundamental. ",
        "introduction": "Super computers are allowing a rapid development of numerical simulations of large N--body systems in Astrophysics. These systems are generally composed by both collisionless matter (such as: stars, galaxies, ...) and collisional matter (i.e. gas). Both phases are usually characterized by being self--gravitating, that is the dynamics of the bodies (stars or fluid elements) is strongly influenced by the gravitational field produced by the bodies themselves.  This {\\it self-influence} is what makes the evaluation of the long--range gravitational force the heaviest computational task to perform in a dynamical simulation. In fact, the number of terms which has to be considered in a direct and trivial evaluation of all the interactions between bodies grows like $N^2$, and since many astrophysically realistic simulations require very large $N$ (greater than $10^5$), such a direct numerical evaluation seems hard to face with presently available computers.  To overcome this problem various approximate techniques to compute gravitational interactions have been proposed. Among them, the Tree--code algorithm proposed by Barnes \\& Hut\\footnote{ Barnes J., Hut P. ``A hierachical $O(N\\log N)$ force calculation algorithm''. {\\it Nature}, vol. 324, p. 446 (1986).} is now widely used in Astrophysics because it does not require any spatial fixed grid (like, for example, methods based on the solution of Poisson's equation). This makes it particularly suitable to follow very inhomogeneous and variable (in time) situations, typical of self-gravitating systems out of equilibrium. In fact its intrinsic capability to give a rapid evaluation of forces allows spending more CPU-time to follow fast dynamical evolution, in contrast to other higher accuracy methods that are more suitable for other physical situations, e.g. dynamics of polar fluids, where the Coulomb term is present.  With the help of the parallelization of our codes, we intend to increase by one or two order of magnitude the number of particles we can use to represent physical systems, in respect to that generally adopted on serial computers ($\\sim 10^4$). In particular our first scientifical aim is the study of close encounters between massive black holes and globular clusters. These latter are systems formed by more than $10^5$ stars gravitationally bounded in a spherical peaked distribution. Such a problem is important in the effort to understand better the nature and formation mechanisms of the {\\it Active Galactic Nuclei}\\ \\footnote{ Capuzzo-Dolcetta R., Miocchi P., ``Galactic Nuclei Activity Sustained by Globular Cluster Mass Accretion'', {\\it PaSS} (1998) in press.}. We hope parallelization makes possible to represent each star with a single particle, in a one--to--one correspondence. This fact clearly will make simulations much more physically meaningful. ",
        "conclusions": ""
    },
    "9803/astro-ph9803171_arXiv.txt": {
        "abstract": "The isothermal gravitational collapse and fragmentation of a molecular cloud region and the subsequent formation of a protostellar cluster is investigated numerically. The clump mass spectrum which forms during the fragmentation phase can be well approximated by a power law distribution $dN/dM \\propto M^{-1.5}$. In contrast, the mass spectrum of protostellar cores that form in the centers of Jeans unstable clumps and evolve through accretion and $N$-body interaction is best described by a log-normal distribution. Assuming a star formation efficiency of $\\sim\\!10\\;\\!$\\%, it is in excellent agreement with the IMF of multiple stellar systems. ",
        "introduction": "\\label{sec:intro}  Understanding the processes leading to the formation of stars is one of the fundamental challenges in astronomy and astrophysics.  However, theoretical models considerably lag behind the recent observational progress. The analytical description of the star formation process is restricted to the collapse of isolated, idealized objects (Whitworth \\& Summers 1985). Much the same applies to numerical studies (e.g.~Boss 1997, Burkert et al.~1997 and reference therein).  Previous numerical models that treated cloud fragmentation on scales larger than single, isolated clumps were strongly constrained by numerical resolution.  Larson (1978), for example, used just 150 particles in an SPH-like simulation.  Whitworth et al.~(1995) were the first who addressed star formation in an entire cloud region using high-resolution numerical models. However, they studied a different problem: fragmentation and star formation in the shocked interface of colliding molecular clumps.  While clump-clump interactions are expected to be abundant in molecular clouds, the rapid formation of a whole star cluster requires gravitational collapse on a size scale which contains many clumps and dense filaments.  Here, we present a high-resolution numerical model describing the dynamical evolution of an entire {\\em region} embedded in the interior of a molecular cloud. We follow the fragmentation into dense protostellar cores which form a hierarchically structured cluster. ",
        "conclusions": "\\label{sec:summary}  Large-scale collapse and fragmentation in molecular clouds leads to a hierarchical cluster of condensed objects whose further dynamical evolution is extremely complex. The agreement between the numerically-calculated mass function and the observations strongly suggests that gravitational fragmentation and accretion processes dominate the origin of stellar masses.  The final mass distribution of protostellar cores in isothermal models is a consequence of the chaotic kinematical evolution during the accretion phase.  Our simulations give evidence, that the star formation process can best be understood in the frame work of a probabilistic theory. A sequence of statistical events may naturally lead to a log-normal IMF (see e.g.~Zinnecker 1984, Adams \\& Fatuzzo 1996; also Price \\& Podsiadlowski 1995, Murray \\& Lin 1996, Elmegreen 1997)."
    },
    "9803/chao-dyn9803019_arXiv.txt": {
        "abstract": "We study the dynamical and statistical behavior of the Hamiltonian Mean Field (HMF) model in order to investigate the relation between microscopic chaos and phase transitions. HMF is a simple toy model of $N$ fully-coupled rotators which shows a second order phase transition. The canonical thermodynamical solution is briefly recalled and its predictions are tested numerically at finite $N$. The Vlasov stationary solution is shown to give the same consistency equation of the canonical solution and its predictions for rotator angle and momenta distribution functions agree very well with numerical simulations. A link is established between the behavior of the maximal Lyapunov exponent and that of thermodynamical fluctuations, expressed by kinetic energy fluctuations or specific heat. The extensivity of chaos in the $N \\to \\infty$ limit is tested through the scaling properties of Lyapunov spectra and of the Kolmogorov-Sinai entropy. Chaotic dynamics provides the mixing property in phase space necessary for obtaining equilibration; however, the relaxation time to equilibrium grows with $N$, at least near the critical point. Our results constitute an interesting bridge between Hamiltonian chaos in many degrees of freedom systems and equilibrium thermodynamics. ",
        "introduction": "Many-particle systems  can show collective behavior when the average kinetic energy is small enough. This collective macroscopic behavior can coexist with chaos at the microscopic level. Such a behavior is particularly evident for systems that have a phase transition, for which a nonvanishing order parameter measures the degree of macroscopic organization, while at the microscopic level chaotic motion is a source of disorder. The latter  can induce non trivial time dependence in the macroscopic quantities, and it would be desirable to relate the time behavior of such quantities and their fluctuations to the chaotic properties of microscopic motion, measured through the Lyapunov spectrum. A naive idea is that an increase of chaos as the energy (temperature) is increased should be accompanied with a growth of fluctuations of some macroscopic quantity. These should be maximal at the critical point and then drop again at high energy. In this paper we study a model of $N$ fully-coupled Hamiltonian rotators which realizes such a behavior, it has been called Hamiltonian Mean Field (HMF) model~\\cite{antoni,latora}. It can also be considered as a system of interacting particles moving on a circle. This system has a second order phase transition and in the ordered phase the rotators are clustered; the high temperature phase is a gaseous one, with the particles uniformly distributed on the circle. It has been shown in ref.~\\cite{latora} that the maximal Lyapunov exponent grows up to the critical energy density $U_c$ and then drop to zero in the whole high temperature phase in the $N \\to \\infty$ limit. Correspondingly one observes a growth of kinetic energy fluctuations up to the critical point and then a phase of vanishing fluctuations. Finite $N$ effects complicate this simple picture. In the high temperature phase the maximal Lyapunov exponent vanishes quite slowly (with $N^{-1/3}$) and  finite size effects influence the first region below the critical point. In this region the system displays metastability: starting far from equilibrium, this is reached in a time $\\tau_r$ which grows with $N$. On the contrary, the extremely low energy phase is characterized by a weak $N$ dependence, with the maximal Lyapunov exponent $\\lambda_1$ which behaves as $\\lambda_1 \\sim \\sqrt{U}$.  Although the model is extremely simplified, it shares many features with more complex models, for which the relation between chaotic motion at the microscopic level and collective macroscopic properties has been studied. Let us mention studies in solid state physics and lattice field theory~\\cite{solid,nayak,dellago,yama,lapo}. However, it has been actually in nuclear physics~\\cite{ata,cmd}, where there is presently a lively debate on multifragmentation phase transition~\\cite{ata,cmd,gsi,eos,bondgro,bond,perco,cmd1,mastinu}, that the interest in the connection between chaos and phase transitions has been revived. In this case in fact, an energy/temperature relation quite close to the HMF model has been observed~\\cite{gsi} and critical exponents have been measured experimentally~\\cite{eos}. Statistical thermodynamical models~\\cite{bondgro} and percolation approaches ~\\cite{perco} have been proved to give a good description of the experimental data, though  the dynamics is missing. On the other hand classical molecular dynamics models \\cite{cmd,bond,cmd1} seem to contain all the main ingredients, but have the disadvantage that a detailed understanding of the dynamics  can be too complicated. In this respect, the HMF model can be very useful in clarifying some general dynamical features which could be eventually compared with real experimental data. In fact, when studying nuclear multifragmentation, one deals with excited clusters of  100-200 particles interacting via long-range (nuclear and Coulomb) forces. Quantum effects are relevant only at very low energy. In fact in the nuclear case, at very low energy, $T$ is not linear in $U$, but grows as $\\sqrt{U}$ because nucleons are fermions~\\cite{gsi,bondgro}. However, a classical picture should be quite realistic in the critical region where the excitation energy is substantial ~\\cite{mastinu}.  In this paper we present new numerical data concerning both statistical quantities, like specific heat and distribution functions, and chaotic probes, like Lyapunov spectra and Kolmogorov-Sinai entropy. Moreover, we add to the theoretical analysis of the model a thorough treatment of differences in the fluctuating quantities between the canonical and microcanonical ensembles. We also investigate in detail the relaxation to equilibrium and compare numerical results with a complete self-consistent Vlasov calculation of distribution functions. Finally a  comparison of numerically obtained maximal Lyapunov exponents with theoretical formulas is attempted.  The paper is organized as follows. In Sec. 2 we  briefly discuss the details of the HMF model. The equilibrium statistical mechanics and the continuum Vlasov solution are described in Sects. 3 and 4 respectively. In Sec. 5 we discuss the relaxation to equilibrium and in Sec. 6 we present the numerical calculations of the Lyapunov spectra and Kolmogorov-Sinai entropy as a function of the energy and $N$. Analytical estimates are discussed in Sec. 7 and conclusions are drawn in Sec. 8. ",
        "conclusions": "We have investigated the dynamical and statistical behavior of  a system with long-range forces showing a second order phase transition. Both the maximal Lyapunov exponent $\\lambda_1$ and the Kolmogorov-Sinai entropy density $S_{KS}/N$ are peaked at  the phase transition point, where kinetic energy fluctuations and specific heat are maximal. There is actually a small shift to lower energies due to finite size effects. The latter are present also in the Lyapunov spectra and in the Kolmogorov-Sinai entropy. Above the phase transition point, both $\\lambda_1$ and $S_{KS}$ vanish as $N \\to \\infty$. We think that this toy model contains some important ingredients to understand the behavior of macroscopic order parameters when dynamical chaos is present at the microscopic level. Most of our findings are probably common to other Hamiltonian systems showing second order phase transitions. In particular our results could be very important in order to understand the relaxation to the equilibrium solution and the success of statistical approaches in describing the  nuclear multifragmentation phase transition.  \\begin{ack} We thank M.C. Firpo for communicating us her results before publication and P. Holdsworth for interesting suggestions. We thank A. Torcini for many useful discussions and a careful reading of the text. A.R. thanks the Centre for Theoretical Physics of MIT for the kind hospitality and M. Robnik for stimulating discussions during his visits  at CAMTP in Maribor, Slovenia. V.L. and S.R. thank INFN for financial support. S.R. thanks CIC, Cuernavaca, Mexico for financial support. This work is also part of the European contract No. ERBCHRXCT940460 on ``Stability and universality in classical mechanics\".  {\\it In the 60's, Boris Chirikov was also an explorer of the  (no man's land at that time) relation between chaotic motion and statistical behavior in classical systems with many degrees of freedom. We hope that he will be interested by this work.} \\end{ack}"
    },
    "9803/astro-ph9803084_arXiv.txt": {
        "abstract": "Using direct N-body simulations which include both the evolution of single stars and the tidal field of the parent galaxy, we study the dynamical evolution of globular clusters and rich open clusters. We compare our results with other N-body simulations and Fokker-Planck calculations.  Our simulations, performed on the GRAPE-4, employ up to 32,768 stars.  The results are not in agreement with Fokker-Planck models, in the sense that the lifetimes of stellar systems derived using the latter are an order of magnitude smaller than those obtained in our simulations.  For our standard run, Fokker-Plank calculations obtained a lifetime of 0.28 Gyr, while our equivalent $N$-body calculations find $\\sim4$ Gyr.  The principal reason for the discrepancy is that a basic assumption of the Fokker-Plank approach is not valid for typical cluster parameters.  The stellar evolution timescale is comparable to the dynamical timescale, and therefore the assumption of dynamical equilibrium leads to an overestimate of the dynamical effects of mass loss. Our results suggest that the region in parameter space for which Fokker-Planck studies of globular cluster evolution, including the effects of both stellar evolution and the galactic tidal field, are valid is limited. The discrepancy is largest for clusters with short lifetimes. ",
        "introduction": "Theoretical models for the evolution of star clusters are generally too idealized for comparison with observations.  However, detailed model calculations with direct $N$-body methods are not feasible for real globular clusters, even with fast special-purpose computers such as GRAPE-4 (Makino et al.\\ 1997)\\nocite{1997ApJ...480..432M} or advanced parallel computers (Spurzem \\& Aarseth 1996).\\nocite{1996MNRAS.282...19S}  If we could scale the results of $N$-body simulations with relatively small numbers of particles (such as $\\sim30,000$) to real globular clusters, then it would become feasible to perform computations with relatively small numbers of particles and still derive useful qualitative conclusions about larger, more massive systems.  However, to determine the proper scaling is difficult because the ratio between two fundamental time scales, the relaxation times and the dynamical time, is proportional to $N$.  In typical globular clusters, this ratio exceeds $10^3$ and the two time scales are well separated.  In $N$-body simulations, the ratio is generally much smaller.  The inclusion of realistic effects such as mass loss due to stellar evolution and the effect of galactic tidal fields (with the galaxy approximated as a point mass, but also with the inclusion of disc shocking) further complicate the scaling problem (see, e.g., Heggie 1996).\\nocite{heg96} A proper treatment of stellar evolution is particularly problematic, since its characteristic timescale changes as stars evolve.  Chernoff \\& Weinberg (1990, CW90)\\nocite{cw90} performed an extensive study of the survival of star clusters using Fokker-Planck calculations which included 2-body relaxation and some rudimentary form of mass loss from the evolving stellar population.  In their simulations the number of particles is not specified.  Their models are defined by the initial half-mass relaxation time and by the initial mass function of the cluster.  Since their models do not specify the number of stars per cluster, each of their model calculations corresponds to a one-dimensional series of models, when plotted in a plane of observational values, such as total mass versus distance to the galactic center (Fig.\\ 1).  All points of the solid line in that figure correspond to a single calculations by CW90, since they have an identical relaxation time.  As we will see later, it is useful to consider other series of models, for which the crossing time is held constant while varying the mass.  An example of such a series is indicated by the dashed line in Fig.\\ 1.  The shapes of these lines are derived under the assumption of a flat rotation curve for the parent galaxy.  The main conclusion of CW90 was that the majority of the simulated star clusters dissolve in the tidal field of the galaxy within a few hundred million years. Fukushige \\& Heggie (1995, FH95)\\nocite{fh95} studied the evolution of globular clusters using direct $N$-body simulation, using the same stellar evolution model as used by CW90. They used a maximum of 16k particles and a scaling in which the dynamical timescale of the simulated cluster was the same as that of a typical globular cluster, corresponding to one of vertical lines in Fig. 1.  FH95 found lifetimes much longer than those in CW90's Fokker-Planck calculations, for the majority of the models used in CW90.  However, the reason for the discrepancy is rather unclear, because the calculations of FH95 and those of CW90 differ in several important respects.  The relaxation times differ because FH95 held the cluster crossing time fixed in scaling from the model to the real system. However, the crossing times themselves are also different, since the crossing time is by definition zero in a Fokker-Planck calculation. Finally, the implementation of the galactic tidal field is also quite different.  CW90 used a simple boundary condition in energy space (spherically symmetric in physical space), in which stars were removed once they acquired positive energy, but the underlying equations of motion included no tidal term.  FH95 adopted a much more physically correct treatment, including tidal acceleration terms in the stellar equations of motion and a proper treatment of centrifugal and coriolis forces in the cluster's rotating frame of reference (see FH95).  \\begin{figure} \\centerline{ \\psfig{file=fig_isomodels.ps,bbllx=570pt,bblly=40pt,bburx=110pt,bbury=690pt,height=5cm,angle=-90}} \\caption {Cluster mass versus the distance to the galactic center.  The solid line indicates the model parameters for which the relaxation time is constant (iso relaxation time); the dashed line indicates the initial conditions for which the crossing time of the star cluster is constant (iso crossing time) } \\label{isomodels}\\end{figure}  In order to study the behavior of star clusters with limited numbers of stars, and to compare with the results of the Fokker-Planck simulations of CW90, we selected one of their models and perform a series of collisional N-body simulations in which the evolution of the individual stars is taken into account.  According to CW90 the results should not depend on the number of stars in the simulation as long as the relaxation time is taken to be the same for all models. It is, among others, this statement which we intend to study.  We find that for this set of initial conditions Fokker-Planck models do not provide a qualitatively correct picture of the evolution of star clusters. The effects of the finite dynamical time scale are large, even for models whose lifetime is several hundred times longer than the dynamical time.  The main purpose of this paper is to study the survival probabilities of star clusters containing up to a few tens of thousands of single stars, in order to gain a deeper understanding of the influence of the galactic tidal field and the fundamental scaling of small $N$ clusters to larger systems.  Only single stars are followed; primordial binaries are not included.  The computation of gravitational forces is performed using the GRAPE-4 (GRAvity PipE, see \\cite{emf+93}, a special-purpose computer for the integration of large collisional $N$-body systems).  Hardware limitations (speed as well as storage) restrict our studies to $\\aplt 32$k particles.  The dynamical model, stellar evolution, initial conditions and scaling are discussed in Sect. 2.  Section 3 reviews the software environment and the GRAPE-4 hardware, and discusses the numerical methods used. The results are presented in Sect. 4 and discussed in Sect. 5; Sect. 6 sums up. ",
        "conclusions": "We have followed the evolution of a star cluster, to the point of dissolution in the tidal field of the parent galaxy, taking into account both the effects of stellar dynamics and of stellar evolution. Our calculations are based on direct $N$-body integration, coupled to approximate treatments of stellar evolution.  Our results differ greatly from those obtained with Fokker-Planck calculations, as presented by CW90: their model clusters dissolve after a few times $10^8$ years, whereas our equivalent model clusters live at least ten times longer.  As we discussed in the previous section, a number of different reason conspire to produce such a drastic difference.  Our hope was that we would be able to find a way to bridge our $N$-body results and previous results based on Fokker-Planck approximations.  The fact that the GRAPE-4 special-purpose hardware allowed us to model much larger numbers of particles, reaching to within an order of magnitude of that of real globular clusters, seemed to indicate that it would finally be possible to make a firm connection between the two types of simulations.  However, our results indicate that no clear process of extrapolation has emerged yet. Even within the different runs we have studied, extrapolation from the smaller to the larger number of stars would have resulted in rather large errors.  This suggests that further extrapolation will suffer from the same fate.  In the present paper, we have studied in detail a single model. However, the way the result depends on the time scaling might be different for other models.  In a subsequent paper, we plan to carry out a systematic study, similar to the one we have presented here, for a much wider range of initial models."
    },
    "9803/hep-ph9803270_arXiv.txt": {
        "abstract": "The idea that the universe might be open is an old one, and the possibility of having an open universe arise form inflation is not new either. However, a concrete realization of a consistent single-bubble open inflation model is known only recently. There has been great progress in the last two years in the development of models of inflation consistent with observations in such an open universe.  In this overview I will describe the basic features and the phenomenological consequences of such models, making emphasis in the predictions of the CMB temperature anisotropies that differ from ordinary inflation. ",
        "introduction": "The idea that the universe might be open is an old one, see e.g. \\cite{Peebles}. Early attempts to accomodate standard inflation in an open universe~\\cite{ratra} failed to realize that in usual inflation homogeneity implies flatness~\\cite{turner}, due to the Grishchuck-Zel'dovich effect~\\cite{GZE}. The possibility of having a truly open universe arise form inflation is not new either, see \\cite{Gott}, via the nucleation of a single bubble in de Sitter space. However, a concrete realization of a consistent model is known only recently, the single-bubble open inflation model~\\cite{singlebubble,LM}. Soon afterwards there was great progress in determining the precise primordial spectra of perturbations~[8-19], most of it based on quantum field theory in spatially open spaces. Simultaneously there has been a large effort in model building~\\cite{LM,Green,induced,GBL} and constraining the existing models from observations of the temperature power spectrum of cosmic microwave background (CMB) anisotropies~[21-24].  In this review talk I will concentrate in model building and constraints from CMB anisotropies. We will describe the nature of the various primordial perturbations and give the corresponding spectra, without deriving them from quantum field theoretical arguments. The interested reader should find this in the literature. We will then compute the corresponding angular power spectra of temperature anisotropies in the CMB. Furthermore, we will give a review of the different single- and mutiple-field open inflation models and constrain their parameters from present observations of the CMB anisotropies.  Sometimes this is enough to rule out some of the models.  Finally, we will describe how future observations of the CMB temperature and polarization anisotropies might be able to decide among different inflationary models, both flat and open inflation ones. ",
        "conclusions": "Single-bubble open inflation is an ingenious way of reconciling an infinite open universe with the inflationary paradigm. In this scenario, a symmetric bubble nucleates in de Sitter space and its interior undergoes a second stage of slow-roll inflation to almost flatness. In the near future, observations of the microwave background with the new generation of satellites, MAP and Planck, will determine with better than 1\\% accuracy whether we live in an open universe or not. It is therefore crucial to know whether inflation can be made compatible with such a universe. Single-bubble open inflation models provide a natural scenario for understanding the large scale homogeneity and isotropy. Furthermore, these inflationary models generically predict a nearly scale invariant spectrum of density and gravitational wave perturbations, which could be responsible for the observed CMB temperature anisotropies. Future observations could then determine whether these models are compatible with the observed features of the CMB power spectrum. For that purpose it is necessary to know the predicted power spectrum with great accuracy. Open models have a more complicated primordial spectrum of perturbations, with extra discrete modes and possibly large tensor anisotropies. In order to constrain those models we have to compute the full spectrum for a large range of parameters.  In this review we have shown that the simplest single-field models of open inflation are not only fine tuned, but actually ruled out because they induce too large tensor anisotropies in the CMB, which is incompatible with present observations. On the other hand, two-field models generically do not lead to infinite open universes, as previously thought, but to an ensemble of very large but finite inflating `islands'. Each one of these islands will be a quasi-open universe. We may happen to live in one of those patches, where the universe {\\em appears} to be open. This new effect, semiclassical in origin, was recently discussed in Ref.~\\cite{quasi} where it was found that many of the present models are in fact quasi-open. This does not mean that they are not good cosmological models. If the co-moving size of the inflating islands is sufficiently large, then the resulting semiclassical anisotropy may be unobservable. We have shown however that such a component imposes very stringent constraints on the models. Most of them have a narrow range of parameters for which they are compatible with observations.  It is perhaps worth mentioning here some alternative proposals (not single-bubble) for the generation of an open universe in the context of inflation. First of all, the group of Roma~\\cite{Roma} proposed a model based on higher order gravity that induces bubble nucleation and later percolation, resulting in a distribution peaked at $\\Omega_0\\simeq0.2$. Perhaps the most striking recent results are those of Hawking-Turok~\\cite{HawTur} and Linde~\\cite{creation}, who claim that an open inflationary universe could have been created directly from the vacuum, without the intermediate de Sitter phase."
    },
    "9803/astro-ph9803059_arXiv.txt": {
        "abstract": "Numerical simulations of galaxy clusters including two species -- baryonic gas and dark matter particles -- are presented. Cold Dark Matter spectrum, Gaussian statistics and flat universe are assumed. The dark matter component is evolved numerically by means of a standard {\\it particle mesh} method. The evolution of the baryonic component has been studied numerically by using a multidimensional (3D) hydrodynamical code based on {\\it modern high resolution shock capturing} techniques. These techniques are specially designed for treating accurately complex flows in which shocks appear and interact. With this picture, the role of shock waves in the formation and evolution of rich galaxy clusters is analyzed. Our results display two well differenced morphologies of the shocked baryonic matter: filamentary at early epochs and quasi-spherical at low redshifts. ",
        "introduction": "Galaxy clusters are the largest systems gravitationally bounded in the Universe. Their study has been a fashion topic in Cosmology since last years. Work on topics related with galaxy clusters is worthly to: i) understand the formation, evolution, dynamics and morphology of these systems, ii) learn on the physical processes involved in them, and iii) find out some information concerning with fundamental parameters in Cosmology as density parameter ($\\Omega$) , Hubble's constant ($H$), and the spectrum of the primordial density field.  During last years technical improvements have produced huge quantity of data about galaxy clusters. Let us mention the new galaxy surveys (Guzzo 1996, and references therein) and the extensive observation in X-rays using satellites as ROSAT or ASCA. These huge volume of data strongly motivates a lot of theoretical work trying to explain the observational results.  From the theoretical point of view, numerical simulations are the best tools to understand physics involved in galaxy clusters. At the beginning, numerical simulations of galaxy clusters where performed using N-body techniques. Since then, they have been extensively used and  have produced important results ( see, e.g, Efstathiou et al. 1985, Bertschinger \\& Gelb 1991, Xu 1995). Next step in the full description of galaxy clusters was to introduce in the picture a baryonic component. The numerical methods developed in order to deal with baryonic matter were more sofisticated and expensive in computational resources. As a consequence, it was not possible to carry out numerical simulations with two species (dark matter and baryonic gas) until late eighties.   Cosmological hydrodynamic codes have been usually classified in two main categories: a) the so-called Lagrangian methods, like the {\\it Smoothed Particles Hydrodynamics} (SPH) or ulterior extensions based on them, and b) Eulerian codes.  SPH methods were first proposed by Gingold \\& Monaghan (1977), and  Lucy (1977). Among the best features of this technique, it should be pointed out its high resolution in dense regions. This property is directly derived from its Lagrangian character. The first implementations of SPH techniques had some weak points: i) The low density regions were badly described due to the Lagrangian character of the method. ii) Discontinuities and strong gradients were poorly solved and an important diffusion was introduced. iii) They were not conservative. Nevertheless, these previous problems were overcome in the modern implementations of these techniques. Improved SPH techniques have been widely developed for cosmological applications (see, e.g., Evrard 1988, Hernquist \\& Katz 1989, Navarro \\& White 1993, Gnedin 1995).   Numerical cosmological codes using an Eulerian approach to study baryonic gas inside galaxy clusters have been also developed. Some of these hydro-codes use {\\it artificial viscosity} in order to deal with shock waves (Cen 1992, Anninos et al. 1994). These techniques require a good calibration of the free parameters which are introduced by hand and state some numerical problems. Recently, a new family of finite difference methods, which use Eulerian approaches and avoid artificial viscosity, has been developed in numerical Cosmology. They are the so-called {\\it high resolution shock capturing methods} (HRSC), the modern extensions of the original Godunov's idea (1959). According to the Riemann solver and the procedure in order to achieve spatial accuracy,  we can distinguish three groups: 1) the ones following Harten's scheme (1983), like Ryu et al. (1993), 2)  those using the analytical solution of the Riemann problem for the Newtonian dynamics of ideal gases and the PPM scheme described by Collela \\& Woodward (1984) , like Bryan et al. (1994), and 3) the codes using Roe's Riemann solver (Roe 1981) plus the MUSCL or PPM cell reconstruction, like in Quilis et al. (1996). In this last reference, the code used in present paper is described and tested appropriately.  An exhaustive comparison among all these kinds of cosmological hydrodynamic codes can be found in Kang et al. (1994).   Due to the Eulerian character of our code, it does not show --in dense regions-- a resolution as good as the Lagrangian ones, and it requires more computational resources. However, HRSC schemes --by construction-- have excellent properties in order to deal with shocks, discontinuities, and strong gradients. HRSC techniques typically solve shocks in two cells. Due to their intrinsic properties, the detection of shocks is independent on the number of cells used in the simulations. It should be pointed out that this last property is really important when three-dimensional simulations are carried out. In these simulations the size of the grid is a stringent constraint due to its high cost in computational resources. Moreover, these methods are  conservative by construction, that is , quantities which should be physically conserved are numerically conserved up to the order of the method.  It should be also noticed that these methods show good results in extreme low density regions (Einfeldt et al. 1991).   As it has been pointed out by several authors, the role of shock waves can be extremely important in order to understand the heating processes in the intracluster medium (ICM). In this paper we are interested in understanding and quantifying the role of shocks. In order to do that it is crucial to use numerical codes able to manage with complex flows. Yes, indeed, one of the important features of HRSC techniques is just to treat numerically shocks and strong discontinuities giving sharp profiles (in a few numerical cells, as we have mentioned above) independently of the size of the grid. Hence, formation, evolution, and interaction of shocks in 3D flows can be analyzed accurately with HRSC schemes, and, consequently, their use is absolutely justified in order to study shocks and their consequences on the ICM's dynamics.  Hereafter, $t$ stands for the cosmological time, $t_0$ is the age of the Universe, $a(t)$ is the scale factor of a flat background. Function $\\dot{a}/a$ is denoted by $H$, where the dot stands for the derivative of $a$ with respect to the cosmological time. Hubble constant is the present value of $H$; its value in units of $100 \\ Km \\ s^{-1} \\ Mpc^{-1}$ is the reduced Hubble constant $h$. In our computations we have assumed $h=0.5\\, $. Velocities are given in units of the speed of light. Baryonic, dark matter, and background mass density are denoted by $\\rho_{_b}$, $\\rho_{_{DM}}$, and $\\rho_{_{B}}$, respectively. The density contrast is $\\delta_b=(\\rho_b-\\rho_{_{B}})/\\rho_{_{B}}$ for baryonic matter, analogously is defined $\\delta_{_{DM}}$ for dark matter.   The plan of this paper is as follows: In Section 2, our numerical cluster model is described. In Section 3, the results of the simulations are analyzed. Finally, a general discussion is presented in Section 4. ",
        "conclusions": "In this paper we have used some numerical techniques recently applied to Cosmology. These techniques , HRSC, seems to be the most suitable in order to study the role of shocks in galaxy cluster evolution. The choice is justified by their properties to handle shocks. The capability of these techniques to capture shocks with very small diffusion is independent of the resolution used in the numerical simulations. Hence, by construction, shocks are captured even using coarse grids. This property is crucial in 3D applications.   Previous sections illustrate the fact that non adiabatic processes, due to shocks, take an important role in the description of the ICM. In the model presented in this paper, that is, a baryonic fluid plus dark matter component coupled gravitationally, shocks are able to heat the ICM until values compatible with observational data.  The calculations have been carried out in two cases: with and without cooling processes. This procedure allows us to distinguish between non-adiabatic effects coming from shocks and the ones from cooling. The role of the cooling, even when it could be important in other scenarios, is  irrelevant for the simulations considered in this paper, while shocks play the most important role.  In the picture describing the dynamics of the baryonic component, there are some clues showing the presence of shocks. Examining  the quantities sensitive to shocks, all of them evidence the formation of a quasi-spherical shock. This shock seems to arise around $z\\sim 2$ at the cluster center and moves outwards. Nevertheless, some irregular shocks could form at $z \\geq 2$. This conclusion is supported by the behaviour of the entropy profiles (see Fig. 5), and the existence of shocked cells at these times (see Fig. 8). The quasi-spherical shock would form from the collapse of the quasi-spherical global structure , while other smaller shocks -- with a filamentary morphology -- would arise from some collapsing substructure and merging processes. In short, previous discussion manifest two different regimes in the shock formation.   It should be noticed that the structures simulated in this paper correspond, due to the initial conditions, to a large Abell cluster. For this kind of clusters, gravitational collapse is fast and the dynamics is violent. Shocks form earlier and are stronger than in  others smaller cluster-like objects ($ < 3\\sigma$).    Some discussion on the numerical resolution of the simulations is needed. The one used in present paper ($\\sim 0.3 Mpc$) is not enough to simulate the very center of the clusters and galaxy formation, but it suffices to study the role of the shocks in ICM. It should be kept in mind that HRSC techniques are able to resolve shocks even with coarse grids. Nevertheless, higher resolution would be desiderable to perform more complete simulations. Improvements in numerical resolution will introduce smaller scales in the problem, as a consequence, the physics of the model should be enriched in order to describe this new scenario. Chemical reactions and  radiative transfer should be considered."
    },
    "9803/astro-ph9803090_arXiv.txt": {
        "abstract": "We present a scenario for the formation and evolution of disk galaxies within the framework of an inflationary cold dark matter universe, and we compare the results with observations ranking from the present-day up to $z\\sim 1$. The main idea in this scenario is that galactic disks are built-up inside-out by gas infall with an accretion rate driven by the cosmological mass aggregation history (MAH). In Avila-Reese et al. (1997) the methods to generate the MAHs of spherical density fluctuations from a Gaussian random field, and to calculate the gravitational collapse and virialization of these fluctuations, were presented. Assuming detailed angular momentum conservation during the gas (5\\% of the total mass) contraction, a disk in centrifugal equilibrium is built-up within the forming dark matter halo. The primordial angular momentum is estimated through the Zel'dovich approximation and normalized to the spin parameter $% \\lambda $ given by analytical and numerical studies. The disk galactic evolution is followed through a physically self-consistent approach which considers (1) the gravitational interactions among the dark halo, the stellar and gas disks, and a bulge; (2) the turbulence and energy balance of the interstellar medium; (3) the star formation process due to gas disk gravitational instabilities; and (4) the secular formation of a bulge due to the gravitational instabilities of the stellar disk.  We find that the main disk galaxy properties and their correlations are basically established by the combination of three fundamental physical factors: the mass, the MAH, and the spin parameter $\\lambda $. Models calculated for a statistically significant range of values for these factors predict nearly exponential disk surface brightness profiles with realistic central surface brightnesses $\\mu _{B_0},$ and scale lengths (including low surface brightness galaxies), nearly flat rotation curves, and negative gradients in the B-V color index radial distribution. The main trends across the Hubble sequence of the global intensive properties such as B-V, $\\mu _{B_0},$ the gas fraction $f_g$, and the bulge-to-disk ratio b/d, are reproduced. For a given mass (luminosity) B-V correlates with the maximum circular velocity, and this correlation is in agreement with the scatter of the Tully-Fisher relation.  We interpret the observed color-magnitude, and ``color '' Tully-Fisher relations as a result of the empirical dependence of extinction on luminosity (mass). The model properties tend to form a biparametrical sequence, where B-V and $\\mu _{B_0}$ could be the two parameters. The star formation history depends on the MAH and on the $\\lambda$ parameter. A maximum in the star formation rate for most of the models is attained at $z\\sim 1.5-2.5$, where this rate is approximately 2.5-4.0 times larger than the present one. The scale radii and the bulge-to-total ratio decrease with $z$, while $\\mu _{B_0}$ increase. The B-band TF relation remains almost the same at different redshifts. Our scenario of disk galaxy formation and evolution reveals that the cosmological initial conditions are able to determine the main properties of disk galaxies across the Hubble sequence and predict evolutionary features for the present-day dominant galaxy population that are in agreement with very recent deep field observational studies ",
        "introduction": "The understanding of the formation and evolution of galaxies is one of the clearest challenges of contemporary astrophysics and cosmology. Since galaxies are both cosmological and astronomical objects, two general approaches can be used in order to study their formation and evolution (e.g., Renzini 1994): ({\\it i}) the deductive approach, through which, starting from some initial conditions given by a theory of cosmic structure formation, one tries to follow the evolutionary processes until the reconstruction of the observable properties of the galaxies and ({\\it ii}) the inductive approach, in which, starting from the present-day properties of galaxies, and through galactic evolutionary models, one tries to reconstruct the initial conditions of galaxy formation; the increasing observational data on galaxies at intermediate and high redshifts will enrich this approach with crucial constraints.  Most of current theories about cosmic structure formation are based on the gravitational paradigm and on the inflationary cold dark matter (CDM) cosmological models. Since these models predict more power for the small density fluctuation scales than for the larger ones, cosmic structures build up hierarchically, through a continuous aggregation of mass. From the point of view of the galaxy cosmogony, a crucial question is whether this aggregation occurs through violent mergers of collapsed substructures and/or through a gentle process of mass aggregation. This question depends on the statistical distribution of the density fluctuation field and on its power spectrum. Nevertheless, even if the dark matter structures assemble through chaotic and violent mergers of subunits, the baryon gas, because of the reheating due to the shocks implied in the collapse, virialization and star formation (SF) feedback processes, will tend to aggregate around the density peaks in a more (spatially) uniform fashion than dark matter do it. Within the framework of the hierarchical clustering theory, from the most general point of view, two could be the galaxy formation scenarios. In one case, the main properties of galaxies, including those which define their morphological types, are supposed to be basically the result of a given sequence of mergers. This picture, that we shall call the{\\it \\ merger scenario,} has been widely applied in semianalytical models of galaxy formation where galaxies are constructed from the cosmological initial conditions through preconceived recipes (e.g., Lacey et al. 1993; Kauffmann, White, \\& Guiderdoni 1993; Cole et al. 1994; Kauffmann 1995, 1996, Baugh, Cole, \\& Frenk 1996). In the other case, the formation and evolution of galaxies is related to a gentle and coherent process of mass aggregation dictated by the forms of the density profiles of the primordial fluctuations: galaxies continuously grow inside-out. We shall call this picture, firstly developed by Gunn (1981, 1987), and by Ryden \\& Gunn (1987), the e{\\it xtended collapse scenario. }% Since disk galaxies ($\\sim 80\\%$ of present-day normal galaxies) could not have suffered major mergers due to the dynamical fragility of the disks (T\\'{o}th \\& Ostriker 1992), the extended collapse scenario results more appropriate to study their evolution.  According to the merger scenario, the bulges of spiral galaxies and the elliptical galaxies arise from the mergers of galactic disks. A natural prediction of this scenario is that spirals with small bulge-to-disk ratios should have bulges older than those of spirals with large bulge-to-disk ratios (e.g., Kauffmann 1996). As Wyse, Gilmore, \\& Franx (1997) have pointed out this does not appear compatible with recent observational data (de Jong 1996a; Peletier \\& Balcells 1996; Courteau, de Jong, \\& Broeils 1997). On the other hand, if elliptical galaxies are the product of relatively recent mergers, then a big dispersion is expected in their color-magnitude relationship (but see Kauffmann 1996). Bower, Lucey, \\& Ellis (1992) showed that for the ellipticals in the Coma Cluster, this relationship is extremely tight. Ellis et al. (1997) confirmed this result for ellipticals in intermediate redshift clusters, up to $z\\sim 0.6.$ The merger scenario could also have serious difficulties from the dynamical point of view: it is not conclusive if mergers of disks are able to reproduce the high central phase-space densities of elliptical galaxies (e.g. Hernquist 1993).  The inductive approach yields the possibility to establish several constraints to the galaxy formation and evolution processes. Galactic evolutionary models{\\it } have shown that due to the rapid disk gas consumption in SF, closed models are not able to explain several properties of disk galaxies, as well as the wide range of colors, gas fractions, etc. that galaxies present across the Hubble sequence (e.g., Larson \\& Tinsley 1978; Tinsley 1980; Larson, Tinsley, \\& Caldwell 1980; Kennicutt 1983; Gallagher, Hunter, \\& Tutukov 1984; Firmani \\& Tutukov 1992, 1994). On the other hand it was shown that the SF time scale in disk galaxies is not controlled by the initial gas surface density (Kennicutt 1983; Kennicutt, Tamblyn, \\& Congdon 1994). Hence, models where gas accretion is introduced are more realistic. Gas accretion could also be necessary to maintain spiral structure. In the case of open models, galaxy formation and galactic evolution might be two related processes where the SF time scale is driven by the gas accretion rate at which the disk is being built up. Infall models of disk galaxy formation have been recently favored by studies of our own Galaxy and nearby galaxies (see for references Cay\\'{o}n, Silk, \\& Charlot 1996). Moreover these inside-out disk formation models seem also to be in agreement with constrictions provided by deep field observations (e.g., Bouwens, Cay\\'{o}n, \\& Silk 1997; Cayon et al.  1996; see also Section 4).  The gas infall rate in luminous galaxies may be controlled by the global process of galaxy formation (cosmological accretion) and/or by a self-regulated process of SF formation. This latter process proposed by White \\& Rees (1978) and White \\& Frenk (1991) is commonly applied in the merger scenario models. According to this mechanism, the gas accretion rate is driven by the cooling of the hot gas corona sustained by the supernova-injected energy. In the extreme situation of instantaneous galaxy formation, supernova gas reheating, halo self-regulated SF, and cooling flows (if the reheated gas was not completely expelled out of the system) become the dominant processes in regulating luminous galaxy evolution. However, the self-regulated halo SF model suffers from some inconsistencies. As Nulsen \\& Fabian (1996) pointed out, supernova feedback over large scales occurs on roughly the same time scales as the SF, not fast enough to tightly regulate the SF rate. In the same way, if a disk forms, then the self-regulating mechanism of SF will apply to the disk where other dynamical conditions prevail (see Firmani, Hern\\'{a}ndez, \\& Gallagher 1996), and not to the halo system. Unless the disk-halo connection is very effective, the SF in the disk will not be regulated by a balance of energy between the supernova input and the halo gas cooling. On the other hand, the X-ray gas corona predicted by the self-regulated SF mechanism lacks observational support, at least for the most massive galaxies for which the X-ray emission would have been above the minimum detection limits of the Rosat and ASCA experiments.  The galactic infall models suggested by the inductive approach, are consistent with the cosmological (deductive) extended collapse scenario of galaxy formation and evolution. In Avila-Reese, Firmani, \\& Hern\\'{a}ndez (1997, hereafter AFH), within the framework of a standard CDM model, the MAHs corresponding to fluctuations of galactic scales were generated from the statistical properties of a Gaussian random field. After calculating the virialization of the fluctuations, a range of realistic dark halo structures were obtained. Now, with the aim to explore whether these cosmological initial conditions are able to predict the evolutionary and observational disk galaxy properties and their correlations, particularly those which go across the Hubble sequence (HS), we shall construct a self-consistent and unified model of disk galaxy formation and evolution in the cosmological context. Within the framework of the extended collapse scenario and using the galactic evolutionary models of Firmani et al. (1996), we shall study the formation and evolution of disks in centrifugal equilibrium into the evolving dark halos. In section 2, the methods we use are described. The model results at $z=0,$ the main predictions of the models, and the comparisons with observations as regards the local (\\S 3.1) and global galactic properties and their main correlations (\\S 3.2) are presented in section 3. In section 4 we compare our evolutionary models with observations at intermediate redshifts ($z\\lesssim 1)$. Finally, the concluding remarks are given in section 5. ",
        "conclusions": "We have modeled the formation and evolution of disk galaxies within the framework of the extended collapse scenario, which is based on the inflationary CDM models. The gas disks in centrifugal equilibrium were built-up under the assumption of detailed angular momentum conservation into spherical virializing dark matter halos whose MAHs were calculated from the initial cosmological conditions. The disk SF is produced by global gravitational instabilities and is self-regulated by an energetic balance of the turbulent gas. The bulges are formed by secular evolution of the stellar disk based on gravitational instabilities. The main predictions of the models are:  1). The disks present exponential surface brightness profiles and negative radial B-V gradients. The scale lengths and central surface brightnesses are in agreement with the observations, including the LSB galaxies.  2). The rotation curves are nearly flat up to the Holmberg radius. Contrary to observational estimations, the rotation curve decompositions show dominion of dark matter down to the galaxy central regions. A constant density core in the dark halo solves this problem.  3). The intensive properties and their correlations (particularly those which go across the HS) of the models corresponding to the local ($z\\approx 0)$ population of disk galaxies, including the LSB galaxies, are determined by the combination of three fundamental physical factors and their statistical distributions, related to the initial cosmological conditions. These three factors are the mass, the MAH, and the primordial angular momentum expressed through the spin parameter $\\lambda .$  4). The intensive properties of the models can be described in a biparametrical sequence, where the parameters may be the color index B-V and the central surface brightness $\\mu_{B_o}$. Each one of these parameters is determined mainly by the MAH and $\\lambda $, respectively. The third fundamental physical factor, the mass, exerts no practical influence the intensive properties. We have shown that the empirical luminosity (mass)-color relation (or equivalently the color TF relation) can be explained by the effects of the metallicity and the extinction.observed dependence of extinction on luminosity (mass). These effects also contribute to decrease the slope of the B-band TF relation.  5). The SF rates of models with the average MAHs and $\\lambda =0.05$ grow by factors of 2.5-4.0 up to $z\\sim 1.5-2.5$ with respect to the SF rates at $z=0 $. After this maximum, the SF rates slowly decrease with $z.$ The SFHs of systems with early, active MAH and/or low $\\lambda ^{\\prime }s$ show high SF rates at high redfshifts ($z>3),$ while the systems with extended MAHs and/or high $\\lambda ^{\\prime }s$ present small SF rates which slowly increase until the present epoch.  6). The structural properties of the models do not change abruptly. Between $z=0$ and $z\\approx 1$ the disk scale radii in average decrease a factor $\\sim 1.3$ and the central surface brightnesses increase $\\sim 1$ mag/arcsec$^2.$ The bulge-to-total luminosity ratio also decreases with $z$ and decreases more severely for the low mass systems. The slopes of the ``structural'' and B-band TF relations do not change with $z.$ In the case of the ``structural'' TF relation, the zero-point decreases (0.75 mag for $z=0.7$ with respect to $z=0$), while for the B-band TF relation the zero-point slightly increases (0.1 mag at $z=0.7$ with respect to $z=0).$  The exploratory models presented in this work show that the main observational characteristics and correlations of disk galaxies can be well understood in the context of the extended collapse scenario, suggesting a direct connection between the conditions prevailing in the early universe and the properties of galaxies today. A serious shortcoming of galaxies emerging from Gaussian CDM cosmological models is the gravitational dominion of DM over baryon matter. The remedy to this problem is the introduction of a core in the DM halo. Fortunately, the intensive galaxy properties and their correlations are not significantly sensitive to the existence or non-existence of such a core, in such a way that all the results presented here are also true for galaxies with a core in their DM halos. The main limitations of our approach are connected to the facts that {\\it (i)} the influence of the environment on galaxy formation and evolution was not taken into account, {\\it (ii)} detailed angular momentum conservation for the baryon gas collapse was assumed, and {\\it (iii)} the mass aggregation was treated only as gas accretion neglecting the possibility of mergers of stellar systems. In future we shall address ways of overcoming these limitations with the aim to improve the model predictions and to explain galactic properties and distributions related to environment."
    },
    "9803/astro-ph9803329_arXiv.txt": {
        "abstract": "Magnetic field-aligned electric fields are characteristic features of magnetic reconnection processes operating in externally agitated magnetized plasmas. An especially interesting environment for such a process are the coronae of accretion disks in active galactic nuclei (AGN). There, Keplerian shear flows perturb the quite strong disk magnetic field leading to intense current sheets.  It was previously shown that given field strengths of 200 G in a shear flow, reconnection driven magnetic field aligned electric fields can accelerate electrons up to Lorentz factors of about 2000 in those objects thus providing us with a possible solution of the injection (pre-acceleration) problem.  However, whereas in the framework of magnetohydrodynamics the formation of the field-aligned electric fields can be described consistently, the question has to be addressed whether the charged particles can really be accelerated up to the maximum energy supplied by the field-aligned electric potentials, since the accelerated particles undergo energy losses either by synchrotron or inverse Compton mechanisms. We pre\\-sent relativistic particle simulations starting from electric and magnetic fields obtained from magnetohydrodynamic simulations of magnetic reconnection in an idealized AGN configuration including nonthermal radiative losses.  The numerical results prove that the relativistic electrons can be effectively accelerated even in the presence of an intense radiation bath.  Energies from~$50 \\Dim{MeV}$ up to~$40 \\Dim{GeV}$ can be reached easily, depending on the energy density of the photon bath. The strong acceleration of the electrons mainly along the magnetic field lines leads to a very anisotropic velocity distribution in phase space. Not even an extremely high photon energy density is able to completely smooth the anisotropic pitch angle distribution which is characteristic for quasi monoenergetic particle beams. ",
        "introduction": "\\label{sec:intro}  Active galactic nuclei (AGN) can be regarded as accreting supermassive black holes surrounded by accretion disks (Camenzind 1990; Miyoshi et al. 1995; Burke and Graham-Smith 1997 and references therein). Relativistic electrons in AGN reveal themselves by hard X-ray (probably due to pair production (cf. Svensson 1987; Done and Fabian 1989)) and $\\gamma$-ray emissions as well as radio observations of superluminous motions (e.g. Abraham et al. 1994).  The $\\gamma$-radiation observed from quasars and bla\\-zars may originate in a distance $R$ of $10^{2-3}$ gravitational radii from the central engine (Dermer and Schlickeiser 1993).  At that distance no significant pair production happens but relativistic leptons scatter via the inverse Compton process the IR-UV radiation of the accretion disk within a relativistically moving jet.  It is well known that ``standard'' mechanisms for the acceleration of high energy leptons, as diffusive shock wave acceleration and resonant acceleration by magnetohydrodynamical (MHD) turbulence can only work efficiently for Lorentz factors larger than $\\gamma_{\\rm crit}\\simeq m_{\\rm p}/{m_{\\rm e}}$ (where $m_{\\rm p}$ and $m_{\\rm e}$ denote the proton and electron masses). Consequently, charged particles accelerated via shocks or MHD~turbulence have to be pre-accelerated which confronts us with the {\\bf injection problem} (Blandford 1994; Melrose 1994) in the AGN context.  In a differentially rotating magnetized accretion disk gas, the evolution of a magnetized corona is quite hard to suppress (Galeev et al. 1979; Stella \\& Rosner 1984).  Driven by the buoyancy force magnetic flux tubes ascend into the disk corona, thereby their footpoints are sheared by the differential rotation of the disk. Either by internal shear or by encountering already pre\\-sent magnetic flux, magnetic reconnection and accompanied rapid dissipation of magnetic energy happens in the coronal plasma. Such a behavior can be studied with great detail for example in the solar corona (Parker 1994 and references therein). In a recent contribution we investigated the possible role of magnetic field-aligned electric fields ($E_\\parallel $) in the context of magnetic reconnection operating in AGN coronae for the pre-ac\\-cel\\-er\\-ation of leptons (Lesch and Birk 1997; hereafter LB). It could be shown that field-aligned electric potential structures in relatively thin current sheets form. For reasonable physical parameters such electric potentials are strong enough to accelerate electrons up to $\\gamma \\approx 2000$, in principle. However, in the framework of MHD the actual energies of the accelerated particles cannot be calculated.  It is the aim of the present contribution to corroborate the model introduced in LB with the help of relativistic particle simulations by taking macroscopic electric and magnetic field configurations obtained by the MHD simulations as an input.  In the next section we resume the MHD model in a nutshell and present the details of the resulting three-dimensional electric and magnetic fields.  In Sec.~\\ref{sec:simulation} we discuss our approach to the numerical study of high-energy particles and show the numerical results dwelling on particle spectra and energies. Eventually, we discuss our findings in Sec.~\\ref{sec:Disc}. ",
        "conclusions": "\\label{sec:Disc}   We addressed the question of charged particle acceleration during magnetic reconnection processes by means of relativistic test particle simulations.  Whereas this point is crucial for a great variety of cosmic plasmas in this contribution we dwelled on the pre-acceleration problem in the AGN context.  Possible further applications include among others as different plasma systems as the terrestrial discrete auroral arcs, the solar coronal flares, radio activity in T-Tauri magnetospheres, non-thermal emission at the edges of high-velocity clouds that hit the galactic plane and the generation of electron beams in the magnetospheres of neutron stars. The starting point for our present investigations are results obtained by an MHD simulation study (LB). In contrast to previous work carried out by different other groups, we were able to study particle acceleration in large-scale non-linearly developed reconnection electromagnetic fields rather than being restricted to the prescription of somewhat idealized analytical field solutions. Moreover, since for the considered parameter regime we have to expect an intense radiation field the test particle simulations were performed including the relevant radiative losses. A very important question is whether charged particles can really be accelerated in the reconnection region up to the pretty high energy values one might deduce from the fluid treatment.  Our findings indicate that, in fact, as expected from the fluid simulations (LB), leptons can be accelerated in reconnection zones located in AGN coronae up to high Lorentz factors; i.e. a significant portion of test particles gain about the maximum energy provided by the generalized field-aligned electric potential structures formed during the magnetic reconnection processes we have modeled within the framework of MHD. Thus, particle acceleration in reconnection zones may be considered as a way out of the injection problem we face in the AGN context.   Since sheared magnetic fields can be expected as a very common phenomena in cosmic environments like accretion disks, stellar coronae and interstellar medium, we think that our relativistic particle studies can be regarded as realistic with respect to the used configuration and the dominant forces. In this contribution the magnetic reconnection process is driven by some resistive mechanism originating from plasma instabilities. The excited electromagnetic oscillations serve as resistance. We note that in plasmas which are collisionless (both no Coulomb collisions and no turbulent wave excitation), the particle inertia presents the ultimate source of resistivity and for the magnetic dissipation. The sheared magnetic fields in collisionless systems evolve into very thin filaments, in which the lifetime of the particle determines the electrical conductivity, thereby allowing for efficient dissipation via effective particle acceleration (Lesch and Birk 1998).  Our simulations are test--particle simulations, thus, we plan for future studies to include, additionally, ponderomotive forces and the back reaction of the current carried by the high energy particles. Whether or not the latter aspect becomes important depends on the density of the run-away electrons limited by the Dreicer electric field (e.g. Benz 1993)."
    },
    "9803/cond-mat9803258_arXiv.txt": {
        "abstract": "The adsorption of large ions from solution to a charged surface is investigated theoretically.  A generalized Poisson--Boltzmann equation, which takes into account the finite size of the ions is presented.  We obtain analytical expressions for the electrostatic potential and ion concentrations at the surface, leading to a modified Grahame equation. At high surface charge densities the ionic concentration saturates to its maximum value.  Our results are in agreement with recent experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803165_arXiv.txt": {
        "abstract": "A method offering an order of magnitude sensitivity gain is described for using quasar spectra to investigate possible time or space variation in the fine structure constant $\\alpha$. Applying the technique to a sample of 30 absorption systems, spanning redshifts $0.5<z<1.6$, obtained with the Keck I telescope, we derive limits on variations in $\\alpha$ over a wide range of epochs.  For the whole sample $\\Delta \\alpha /\\alpha =-1.1\\pm 0.4\\times 10^{-5}$.  This deviation is dominated by measurements at $z>1$, where $\\Delta \\alpha /\\alpha = -1.9\\pm 0.5\\times 10^{-5}$. For $z<1$, $\\Delta \\alpha /\\alpha = -0.2\\pm 0.4\\times 10^{-5}$, consistent with other known constraints.  Whilst these results are consistent with a time-varying $\\alpha$, further work is required to explore possible systematic errors in the data, although careful searches have so far not revealed any. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803023_arXiv.txt": {
        "abstract": "We undertake a quantitative investigation, using Monte Carlo simulations, of the amount by which quasars are expected to exceed radio galaxies in optical luminosity in the context of the `receding torus' model. We compare these simulations with the known behaviour of the \\OIII~$\\lambda$5007 and \\OII~$\\lambda$3727 emission lines and conclude that \\OIII\\ is the better indicator of the strength of the underlying non-stellar continuum. ",
        "introduction": "It is widely believed that radio-loud quasars and radio galaxies differ from each other only in terms of the angle between their radio axes and the line of sight (Scheuer 1987; Barthel 1989). Quasars are observed fairly close to the line of sight ($\\simlt 45\\degree$) and therefore frequently exhibit the effects of beaming, such as superluminal motion and luminous flat-spectrum radio cores, whereas radio galaxies are observed with their axes close to the plane of the sky, and do not show these effects. However, quasars also possess a luminous non-stellar optical continuum and broad emission lines, which are absent in radio galaxies. It has therefore been proposed that there is material around the nucleus which lies preferentially in the plane perpendicular to the radio axis and obscures the central regions from view in radio galaxies. Because of the geometry of this material, it is often referred to as the ``torus'', but other geometries (e.g.\\ a warped disk) are possible. Although the broad lines are hidden from direct view by this material, the narrow line region (NLR) is much larger in size and should be less strongly affected. It is therefore possible that narrow line luminosity could be an isotropic measure of the strength of the non-stellar continuum that is otherwise invisible in radio galaxies.  Jackson \\& Browne (1990) tested this idea by comparing the \\OIII~$\\lambda$5007 luminosities of quasars and radio galaxies with similar redshifts and extended radio luminosities. They reported that quasars are 5--10 times more luminous in this line than radio galaxies, and attributed the difference to higher extinction in radio galaxies. Hes, Barthel \\& Fosbury (1993) performed a similar test using the \\OII~$\\lambda$3727 doublet and found that there was no measurable difference between the line luminosities of the two classes. Since \\OII\\ is at a shorter wavelength than \\OIII\\ and would be more greatly affected by a foreground screen of reddening, they suggested that most of the \\OIII\\ emission is produced close to the nucleus and is still obscured by the torus in radio galaxies, whereas \\OII\\ is produced farther out and is unaffected.  Unfortunately, Jackson \\& Browne's analysis was biased because, although the radio galaxies were selected from the 3CR-based complete sample of Laing, Riley \\& Longair (1983; also called 3CRR), the quasars were drawn from a number of incomplete surveys and were therefore subject to uncertain selection effects, most notably the tendency to preferentially include optically bright objects. Since line and continuum luminosities are very well correlated in quasars, their result would be biased towards finding systematically higher line luminosities in the quasars. In addition, Jackson \\& Browne's radio galaxy sample included some objects with very weak emission lines (Class~B optical spectra; Hine \\& Longair 1979) that it is now believed may not belong to the unified scheme. An unbiased analysis significantly reduces the magnitude of Jackson \\& Browne's result, but the quasars are still about twice as luminous in \\OIII\\ than the radio galaxies. This result is also seen in the higher \\OIII/\\OII\\ ratios in broad-line objects than in narrow-line objects (e.g.\\ Saunders et al.\\ 1989; Tadhunter et al.\\ 1993). At higher redshift, however, the \\OIII\\ luminosities of the two classes of object are comparable (Jackson \\& Rawlings 1997).  Rawlings \\& Saunders (1991) note that the tendency for quasars to have a higher emission line luminosity than radio galaxies could be the result of a classification bias. This could explain the discrepancy in \\OIII\\ luminosities between the two classes, but appears to run counter to the similarity in the \\OII\\ line emission. In this {\\it Letter\\/}, we apply the simple receding torus model (Lawrence 1991; Hill, Goodrich \\& Depoy 1996) to provide a quantitative explanation of the difference in \\OIII\\ luminosities between quasars and radio galaxies (and its redshift dependence), and explain why a similar effect is not seen in \\OII. The free parameters involved in our explanation are constrained by observed quantities independent of the emission line luminosities. We conclude that the luminosity of the \\OIII~$\\lambda$5007 line is an excellent indicator of the strength of the underlying non-stellar continuum. ",
        "conclusions": "We have used a simple receding torus model, whose free parameters are constrained by observation, to show that low redshift 3CRR quasars should be, on average, about twice as luminous in their ionizing continua as radio galaxies of the same radio luminosity. This difference should also be seen in their \\OIII, but not their \\OII, emission line luminosities, in agreement with observation. For samples with a higher quasar fraction, such as the high redshift 3CRR objects, the difference in ionizing luminosities between quasars and radio galaxies should be smaller, and there should therefore be less of a difference in their \\OIII\\ luminosities, again in line with observation.  This model leads to the conclusion that the \\OIII~$\\lambda$5007 emission line, and not the \\OII~$\\lambda$3727 doublet, is an unbiased indicator of the intrinsic optical--ultraviolet luminosity of both quasars and radio galaxies."
    },
    "9803/astro-ph9803215_arXiv.txt": {
        "abstract": "We report an upper limit of $9\\times 10^{12}$ cm$^{-2}$ on the column density of water in the translucent cloud along the line-of-sight toward HD 154368.  This result is based upon a search for the C-X band of water near 1240 \\AA\\ carried out using the Goddard High Resolution Spectrograph of the Hubble Space Telescope. Our observational limit on the water abundance together with detailed chemical models of translucent clouds and previous measurements of OH along the line-of-sight constrain the branching ratio in the dissociative recombination of H$_3$O$^+$ to form water. We find at the $3\\sigma$ level that no more than 30\\% of dissociative recombinations of H$_3$O$^+$ can lead to H$_2$O.  The observed spectrum also yielded high-resolution observations of the Mg II doublet at 1239.9 \\AA\\ and 1240.4 \\AA, allowing the velocity structure of the dominant ionization state of magnesium to be studied along the line-of-sight.  The Mg II spectrum is consistent with GHRS observations at lower spectral resolution that were obtained previously but allow an additional velocity component to be identified. ",
        "introduction": "Translucent clouds have total extinctions $A_{\\rm V}$ in the range of 2-5 magnitudes.  The physical and chemical conditions that characterize translucent clouds  are therefore intermediate between those in diffuse and in dense molecular clouds (Crutcher 1985; van Dishoeck \\& Black 1988).  Photodissociation rates are significantly smaller in translucent clouds than in diffuse clouds, and the column densities of molecules like CO, OH and CS are correspondingly larger. Translucent clouds can be observed through absorption lines of CN, CH, CH$^+$ and C$_2$ as well as through millimeter emission lines of CO and CS. Although they are an abundant constituent of dense molecular clouds (Jacq et al.\\ 1988; Zmuidzinas et al.\\ 1995; Gensheimer et al.\\ 1996; van~Dishoeck \\& Helmich 1996; Cernicharo et al.\\ 1997) and are expected to be the dominant coolant of warm dense regions within such clouds (Neufeld \\& Kaufman 1993; Neufeld, Lepp \\& Melnick 1995), water molecules have not been detected in diffuse or translucent molecular clouds.  In this paper, we report the results of a search for water in the translucent cloud along the line of sight to HD 154368.  HD 154368 is an O9.5 Iab star at an estimated distance of 800 pc from the Sun (Snow et al.~1996) that is situated near the Sco OB 1 association (Blades \\& Bennewith 1973). The line-of-sight toward HD 154368 lies about $\\rm 14^o$ from the core of the $\\rho$ Oph molecular cloud, which is located at a distance of 125$\\pm$25 pc (de Geus, de Zeeuw, \\& Lub 1989) and has a heliocentric radial velocity of $\\rm -6.6\\, km \\,s^{-1}$. The main H I (e.g.\\ Riegel \\& Crutcher 1972) and Na I (e.g.\\ Crawford, Barlow, \\& Blades 1989) absorption features are observed at a radial velocity that is close to that of the $\\rho$ Oph molecular cloud, a result that suggests that the main component of gas toward HD 154368 is probably not close to the star and is more likely to be the outer envelope of a dense molecular cloud only about 125 pc from the Sun. The color excess along the line of sight is E(B--V)=0.82 and most of the gas toward HD 154368 resides in two clouds centered near --3.26 (main component) and --20.95 km s$^{-1}$ (heliocentric).  This line-of-sight has been observed extensively from the ground by means of narrow H I 21 cm absorption features (e.g.~Riegel \\& Crutcher 1972); optical absorption lines of Na~I~D (e.g.~Crawford et al.~1989), CN B--X (0,0) and A--X (0,0), CH, CH$^+$; and near-infrared absorption lines of C$_2$ in the A--X Phillips system at 8750 \\AA\\ (van Dishoeck \\& de Zeeuw 1984). The CH observations can be used to infer the total column density of H$_2$ along the line-of-sight (Gredel, van Dishoeck, \\& Black 1993), a quantity that is needed to determine the atomic depletions. The red system of CN has been used by Gredel, van Dishoeck, \\& Black (1991) to derive an electron density of 0.05-0.15 cm$^{-3}$ for the molecular component. Molecular emission lines are also detectable toward HD 154368. Data on $^{12}$CO $J=1-0$, $J=2-1$, and $J=3-2$ and $^{13}$CO $J=1-0$ have been presented by van Dishoeck et al.~(1991), and have been used to constrain the column density of CO and the density of H$_2$ in the molecular component. These authors confirmed the relatively low density $n_{\\rm H}\\approx 350$ cm$^{-3}$ derived independently from the C$_2$ absorption data. The $^{12}$CO $J=1-0$ distribution over a region of $30'\\times 30'$ around the star is featureless, a result that suggests once more that the cloud and the star are not located close to one another.  The velocity structure of the gas toward HD 154368 is well known through the Na I data obtained using the Ultra High Resolution Facility at the Anglo Australian Telescope (Snow et al.~1996) and high resolution Ca II observations (Crawford 1992). The Na I data indicates seven velocity components at --27.7, --20.95, --18.2, --14.75, --10.5, --3.26, and +5.62 km s$^{-1}$ where 96\\% of the neutral sodium resides in the --3.26 km s$^{-1}$ feature. The Ca II results show five velocity components at --27.6, --20.9, --14.4, --4.3, and +6.8 km s$^{-1}$. These five coincide roughly with the Na I components, but their relative column densities are different. Approximately 50\\% of the ionized calcium resides in the --4.3 km s$^{-1}$ feature with an even distribution over the remaining velocity components.  A detailed discussion of all these observations and the implications they have for the physical conditions of the ambient medium can be found in Snow et al.~(1996). In this work, their results will be adopted and the main focus will be on the resulting OH and H$_2$O chemistry. ",
        "conclusions": "We have reported a $3\\sigma$ upper limit of $9\\times 10^{12}$ cm$^{-2}$ on the column density of water toward HD 154368. We have constructed detailed chemical models which incorporate many existing constraints on the physical conditions of the ambient medium. Together with the known column density of OH along the line of sight, we constrain the H$_2$O branching ratio for dissociative recombination of H$_3$O$^+$ to be smaller than 30\\%. Our results are in agreement with the laboratory studies of Williams et al.\\, which find an oxygen channel in the recombination of H$_3$O$^+$ and a small branching ratio for the production of water, but are mildly inconsistent with the laboratory results of Vejby-Christensen et al.  The observed spectrum of HD 154368 also yielded high-resolution observations of the Mg II doublet at 1239.9 \\AA\\ and 1240.4 \\AA, allowing the velocity structure of the dominant ionization state of magnesium to be studied along the line-of-sight.  The Mg II spectrum is consistent with GHRS observations at lower spectral resolution that were obtained previously but allow an additional velocity component to be identified. Absorption by the Ge II line at 1237.059 \\AA\\ was also detected, the first detection of interstellar germanium along the line-of-sight to HD 154368.  We acknowledge with gratitude the support of NASA grant NAGW-3147 from the Long Term Space Astrophysics Research Program and of grant GO-06739.01-95A from the HST Cycle 6 Program. In the final stages of this project MS was partially supported by NASA through grant HF-01101.01-97A, awarded by the Space Telescope Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA under contract NAS 5-26555. Most of the simulations presented in this work were performed on the Cray YMP operated by the Netherlands Council for Supercomputer Facilities in Amsterdam.  \\newpage"
    },
    "9803/astro-ph9803287_arXiv.txt": {
        "abstract": "ASCA observations of NGC 4636 and a southern region have revealed extended X-ray emission to a radius of about 300 kpc from the galaxy. The symmetric nature of the observed surface brightness around NGC 4636 indicates its association to this galaxy rather than to the Virgo cluster. Model independent estimation of the gravitational mass profile shows a flattening at a radius of $20 \\sim 35$ kpc, where the total mass reaches $\\sim 6\\times 10^{11} M_\\odot$ and a mass-to-light ratio of 23. The mass still increases to larger radii, reaching $9\\times10^{12} M_\\odot$ with a mass-to-light ratio of 300 at $\\sim$ 300 kpc from NGC~4636.  These features suggest presence of a galaxy group surrounding NGC 4636. If such optically dark groups are common among X-ray bright ellipticals, it would explain the very large scatter in their X-ray luminosities with similar optical luminosities. ",
        "introduction": "Spatial distribution of gravitational mass around elliptical galaxies has been estimated from radial distribution of X-ray emitting hot interstellar medium (ISM) (e.g. \\cite{3}; \\cite{4}). However, the outermost radius of the ISM distribution has so far remained undetermined due to a lack of sensitivity.  As a result, the total gravitating masses of elliptical galaxies depend heavily on the assumed overall extent of the ISM\\@.  High sensitivity X-ray observations of bright galaxies are therefore important in determing the mass distribution in these systems.  NGC~4636 is one of the nearest giant elliptical galaxies.  We employ a distance of 17 Mpc (\\cite{8}).  Diffuse thermal X-rays from its ISM make it one of the X-ray brightest elliptical galaxies (e.g. \\cite{3}; \\cite{4}) except those located at the center of clusters.  Although NGC~4636 is located in the Virgo Southern Extension (\\cite{nol}), it is apparently clear of the large-scale X-ray emission from hot intra-cluster gas associated with the Virgo cluster (\\cite{10}; \\cite{11}).  Trinchieri et al.\\ (1994), using {\\it ROSAT}, detected a largely extended X-ray emission up to 18$'$ from NGC~4636, and suggested it to be an extension of the Virgo cluster emission.  Here, we present new results on the extended X-ray emission around NGC~4636 based on a very long observation towards the galaxy center and on an offset observation from ASCA.  Spectral properties of NGC~4636 have been already reported by Matsushita et al.\\ (1997). ",
        "conclusions": "ASCA observations have shown a very extended X-ray emission with a radius of at least $60'$ (300 kpc) surrounding NGC~4636, which is much larger than that detected by {\\it ROSAT}\\@.  We studied radial temperature and density distributions of the gas, and determined gravitational mass profile with and without assuming a double-$\\beta$ model.  According to Fig.\\ 3, the gravitational halo of NGC~4636 appears to terminate once at $R \\sim 10-30$ kpc, where the total mass reaches several $ \\times 10^{11} ~ M_\\odot$.  This inner component probably corresponds to the mass associated with the galaxy NGC~4636. Then, the total mass starts increasing again, forming a halo-in-halo structure whose total size is comparable to that of a galaxy group.  The mass of the whole system and its mass-to-light ratio at 300 kpc reach $\\sim 1 \\times 10^{13}~ M_\\odot$ and $\\sim 300$, respectively, which are again comparable to those of galaxy groups (\\cite{28}).  At $R \\ge 200$ kpc, the X-ray emitting plasma becomes the dominant form of baryons, with its mass reaching 5--8\\% of the total gravitating mass.  This value is considerably higher than those of individual elliptical galaxies, and again close to those found in galaxy groups and poor clusters (\\cite{28}).  These features all support the presence of a gravitational potential with a size of a galaxy group around NGC~4636.  A further support is given by the observed abundance decrease in the X-ray emitting plasma (\\cite{12}; \\cite{15}; \\cite{matusita97}; \\cite{17}), from $\\sim 1$ solar within $\\sim 30$ kpc, to $\\sim 0.2$ solar beyond $\\sim 50$ kpc to a is typical level of groups and clusters of galaxies.   Nolthenius (1993) identified % NGC~4636 as a member of Virgo Southern Extension F Cloud. However the 3-dimensional location of NGC~4636 is far offset from the center of the member galaxies in the cloud, and their velocity dispersion, 463 km s$^{-1}$, is much larger than that inferred from the observed gas temperature.  The diameter of the cloud, 1.26 Mpc, is also much larger than the scale of the X-ray emission.  Therefore, the relation between the X-ray halo and the galaxy cloud remains unclear, but it is very likely that there is some galaxy concentration around NGC~4636.   Extensive studies of ASCA data have revealed (\\cite{17}) that X-ray luminous elliptical galaxies preferentially possess large-scale X-ray halos of a few 100 kpc scale, just like NGC~4636. Such a galaxy may be regarded as a dominant member of a galaxy group even though its evidence is scarce in the optical data. Thus, the X-ray emitting plasma may provide a  better tracer of the total mass distribution than the light emitting matter.  We speculate that X-ray luminous elliptical galaxies may commonly possess such an extended emission with low surface brightness.  X-ray luminosity of elliptical galaxies scatter by 2 orders of magnitude for the same optical luminosity, and its cause has long been a puzzle. The presence and absence of the extended emission can easily account for the large scatter in the total X-ray luminosity, and the low surface brightness of the extended emission, such as in NGC~4636, explains the previous undetection.  We hope that systematic deep exposures from ASCA will solve this problem and bring us a complete understanding of the mass distribution around elliptical galaxies.  \\medskip  K.M. acknowledges support by the Postdoctoral Fellowship of the Japan Society for Promotion of Science.  \\clearpage"
    },
    "9803/astro-ph9803078_arXiv.txt": {
        "abstract": "Recent hard X-ray spectroscopy of active galactic nuclei has strongly suggested that double-peaked, very broad Fe K emission arises from an accretion disk around the central engine. Model fitting of the observed Fe K emission line profile makes it possible to estimate a probable inclination angle of the accretion disk. In order to study the geometrical relationship between the accretion disk and broad emission-line regions (BLRs), we investigate the correlation between the inclination angle of the accretion disk and the velocity width of BLRs for 18 type-1 Seyfert galaxies. We found that there may be a negative correlation between them, i.e., Seyfert nuclei with a more face-on accretion disk tend to have larger BLR velocity widths, suggesting that the BLRs are not coplanar with respect to the accretion disk. The most probable interpretation may be that the BLRs arise from outer parts ({\\it r} $\\sim$ 0.01 pc) of a warped accretion disk illuminated by the central engine. ",
        "introduction": "Given the current paradigm of active galactic nuclei (AGNs), the observed huge luminosities of AGNs are powered by gravitational accretion onto a supermassive black hole (e.g., Rees 1984; Blandford 1990; Antonucci 1993; Peterson 1997). The central engines are considered to be surrounded by dusty tori whose typical inner radii are on the order of $\\sim$ 1 pc (e.g., Antonucci, Miller 1985; Pier, Krolik 1992, 1993). Therefore, in order to understand AGNs, it is very important to investigate the spatial structures of the inner $\\sim$ 1 pc regions in which a supermassive black hole, an accretion disk, warm absorbers, and broad emission-line regions (BLRs) reside.  The innermost constituent is the accretion disk; its typical radius is $\\sim$ (10 --- 100) $R_{\\rm S} \\sim (10^{-4}$ --- $10^{-3})M_{8}$ pc where $R_{\\rm S}$ is the Schwarzschild radius of a black hole and $M_{8}$ is the black-hole mass in units of $10^{8} M_{\\sun}$. The existence of accretion disks in AGNs has been demonstrated by the recent X-ray spectroscopy, collected using {\\it ASCA} (Tanaka et al. 1994). The {\\it ASCA} X-ray spectra of type-1 AGNs show the presence of very broad Fe K$\\alpha$ emission, whose line profile can be fitted well by some accretion-disk models (Tanaka et al. 1995; Fabian et al. 1995; Mushotzky et al. 1995; Iwasawa et al. 1996; Nandra et al. 1997; Reynolds 1997). An ionized accretion disk has also been detected by the recent radio continuum mapping in the archetypical, nearby Seyfert galaxy NGC 1068 (Gallimore et al. 1997).  Another inner constituent is broad-line regions (BLRs), because their typical radii are of on the order of 0.01 pc (e.g., Peterson 1993). One of the most important questions related to the BLRs is how emission-line clouds are distributed around AGNs. Although there is still no definite consensus concerning the dynamical and spatial structure of BLRs, there are three alternative models: 1) the disk model (Shields 1979; see also Osterbrock 1989), 2) the high-velocity streamer model (Zheng et al. 1991), and 3) a pair of conical regions in which photoionized clouds are orbiting randomly with Keplerian motion (Wanders et al. 1995; Wanders, Peterson 1996, 1997; Goad, Wanders 1996). In particular, a recent detailed analysis of the reverberation mapping has strongly suggested that the BLRs of the type-1 Seyfert galaxy NGC 5548 has the third type of geometry (Wanders et al. 1995). It is, however, still not known which model is the most popular one.  Since recent observations have shown that the BLRs are dominated by rotational motion, rather than the radial motion (e.g., Peterson 1993; Wanders et al. 1995), it is interesting to examine whether or not the rotational axis of the BLR is nearly the same as that of the accretion disk. If the disk model for BLRs would be correct, the BLRs may be coplanar with respect to the accretion disk. In fact, double-peaked BLRs have sometimes been considered to arise from accretion disks, themselves [e.g., P\\'erez et al. (1988); Livio, Xu (1997) and references therein; see also, however, Gaskell (1996)]. Motivated by this, we investigate the relationship between the inclination angle of the accretion disk and the width of BLR statistically using published data. ",
        "conclusions": "The most important result in this study is that there is no obvious {\\it positive} correlation between $\\sin i_{\\rm AD}$ with FWZI(H$\\alpha$)/ $L^{1/4}_{\\rm X}$. This suggests that {\\it the BLRs are not coplanar with respect to accretian disks}.  Some Seyfert nuclei in our sample show double-peaked BLRs (DBLRs). It has sometimes been considered that such DBLRs may arise from an accretion disk, itself (e.g., P\\'erez et al. 1988). The DBLR emission profiles of the four Seyfert nuclei in our sample (NGC 3227, NGC 3783, NGC 5548, and 3C 120) were studied by Rokaki et al. (1992) using a standard geometrically-thin accretion- disk model; also, the inclination angles of the DBLRs ($i_{\\rm BLR}$) were derived. We compare these inclination angles with those of the accretion disks in figure 2. This comparison also suggests a negative correlation between $i_{\\rm AD}$ and $i_{\\rm BLR}$, being consistent with our result. This strengthens our suggestion that the BLRs are not coplanar with respect to the accretion disks in the Seyfert nuclei studied here.  Let us consider what kind of geometrical configuration can explain the non-coplanar property. The negative correlation means that the normalized velocity width increases with decreasing inclination angle; i.e., {\\it Seyfert nuclei with a more face-on accretion disk tend to have larger BLR velocity widths}. There may be three alternative ideas to explain this property. One is the bipolar streamer model (e.g., Zheng et al. 1990). If we observe the accretion disk from a face-on view, the velocity width would be widest because the bipolar wind flows along our line of sight. However, this model has an intrinsic difficulty, as claimed by Livio and Xu (1996), because the emitting region on the receding flow (jet) is obscured from view by the accretion disk; the standard, optically thick accretion disk is opaque up to $\\sim$ 1 pc, and, thus, the BLR component behind the disk cannot be seen, because the typical radial distance of BLRs from the central engine is on the order of 0.01 pc (e.g., Peterson 1993). The second idea is that BLRs are located in nearly the same plane as that of an accretion disk, but are orbiting with poloar orbits. If a two-sided jet is ejected with a highly inclined angle with respect to the {\\it global} accretion disk, we can explain the negative correlation. Such a jet model is briefly described by Norman and Miley (1984). This idea is consistent with the recent reverberation mapping result for the BLRs of NGC 5548 because the most likely geometry of the BLRs of this galaxy is a pair of conical regions in which photoionized clouds are orbiting randomly with Keplerian motion (Wanders et al. 1995; Wanders, Peterson 1996, 1997; see also Goad, Wanders 1996). This model may also have the same obscuration problem as that for the above streamer model. However, if the BLR clouds are moving at randomly oriented orbits (Wanders et al. 1995), there may be no obscuration problem. The third idea is that BLRs arise from outer parts of a warped accretion disk. The disk model for BLR is the standard idea (Shields 1977; see also for a review Osterbrock 1989). It has been recently shown that accreting gas clouds probed by water-vapor maser emission at 22 GHz show evidence of significant warping (Miyoshi et al. 1995; Begelman, Bland-Hawthorn 1997). The warping of accretion disks can be driven by the effect of the radiation-pressure force (Pringle 1996, 1997). For typical AGN, the warping may occur at {\\it r} $>$ 0.01 pc (Pringle 1997), which is at a similar distance as BLRs. Therefore, the warped-disk model can explain the observed negative correlation reasonably well. This model is schematically shown in figure 3. Since the degree of warping and the viewing angle are different from AGN to AGN, the negative correlation between $i_{\\rm AD}$ and $i_{\\rm BLR}$ may be blurred as obtained in our analysis, although the poor correlation may be also due to the large errors in the estimate of $i_{\\rm AD}$."
    },
    "9803/astro-ph9803308_arXiv.txt": {
        "abstract": "We present measurements of the oxygen abundances in 64 H{\\sc ii} regions in 12 LSB galaxies. We find that oxygen abundances are low. No regions with solar abundance have been found, and most have oxygen abundances $\\sim 0.5$ to 0.1 solar. The oxygen abundance appears to be constant as a function of radius, supporting the picture of quiescently and sporadically evolving LSB galaxies. ",
        "introduction": "Low surface brightness (LSB) disk galaxies have all the characteristics of unevolved galaxies. Those discovered so far constitute a population of gas-rich, metal-poor galaxies with very low star formation rates (see the review by Bothun et al. 1997).  Their surface brightnesses are a few magnitudes lower than the values commonly found for so-called normal galaxies.  Most of them are rather late-type galaxies, with diffuse spiral arms.  A direct probe of the evolutionary state of these galaxies is the metal abundance in the interstellar medium (ISM).  A low abundance generally indicates only limited enrichment of the ISM and therefore (in a closed system) a small amount of evolution.  Because of their low surface brightness, obtaining spectra of the stellar disks of LSB galaxies is difficult and requires large amounts of telescope time.  Conclusions on metallicities must therefore be derived from spectra of H{\\sc ii} regions.  These usually are the brightest distinct objects in a LSB galaxy.  Their bright emission lines make them more easily observable than the underlying continuum.  The H{\\sc ii} regions that are observed in LSB galaxies are usually giant H{\\sc ii} regions, that are ionized by star clusters rather than by a few stars.  The first measurements of the oxygen abundances in H{\\sc ii} regions in LSB galaxies were presented in McGaugh (1994). He found, using an empirical oxygen abundance indicator, that LSB galaxies are low metallicity galaxies with typical values for the metallicity $Z<\\frac{1}{3} Z_{\\odot}$.  It shows that low metallicities can occur in galaxies that are comparable in size and mass to the bright galaxies that define the Hubble sequence. As LSB galaxies are found to be isolated (Mo et al.\\ 1994), this suggest that surface mass density and environment are as important for the evolution of a galaxy as total mass.  In this paper we present a follow-up study of oxygen abundances in LSB galaxies. We confirm the results by McGaugh (1994) that LSB galaxies are metal-poor. We present two direct measurements of the oxygen abundance from measurements of the [O {\\sc iii}]$\\lambda 4363$ line, supplemented with a large number of empirically determined oxygen abundances.  In addition, for those galaxies where sufficient data are available, we investigate the change in abundance with radius, and show that the measurements are consistent with no gradient. The steeper gradients found in HSB galaxies (Vila-Costas \\& Edmunds 1992 [VCE], Zaritsky et al. 1994, Henry \\& Howard 1995, Kennicutt \\& Garnett 1996) are not present.  It is worth noting that the exact form and magnitude of the Milky Way oxygen gradient has now been consistently reproduced in early-type stars, H{\\sc ii} regions and planetary nebulae, which supports that the extra-galactic H{\\sc ii} region gradients in HSB galaxies are real (Smartt \\& Rolleston 1997).  The lack of abundance gradients in LSB galaxies supports the picture of stochastic and sporadic evolution, where the evolutionary rate only depends on local conditions and not on the global properties of LSB galaxies as a whole.  Section 2 describes the sample selection and observations. Section 3 presents the data, while in Sect.~4 the analysis is described. Section 5 discusses the abundances found. Section 6 discusses reddening towards the H{\\sc ii} regions, while Sect.~7 concludes with presenting the gradients. In Sect.~8 the results are summarized. ",
        "conclusions": "The oxygen abundances in 64 H{\\sc ii} regions in 12 LSB galaxies have been measured.  Oxygen abundances are low. No region with solar abundance has been found, and most have oxygen abundances that are $\\sim 0.5$ to 0.1 solar.  No strong radial oxygen abundance gradients are found. The abundance seems to be constant, rather, as a function of radius, supporting the picture of quiescently and sporadically evolving LSB galaxies."
    },
    "9803/hep-ph9803471_arXiv.txt": {
        "abstract": "\\noindent From a phenomenological point of view, we study active-active and active-sterile flavour-changing (and flavour-conserving) oscillations of Dirac-Majorana neutrinos both in vacuum and in matter. The general expressions for the transition probabilities in vacuum are reported. We then investigate some interesting consequences following from particular simple forms of the neutrino mass matrices, and for the envisaged scenarios we discuss in detail neutrino propagation in matter. Special emphasis is given to the problem of occurrence of resonant enhancement of active-active and active-sterile neutrino oscillations in a medium. The peculiar novel features related to the Dirac-Majorana nature of neutrinos are particularly pointed out. ",
        "introduction": "Today we have several indications in favour of non zero neutrino masses and mixing.\\\\ The solar neutrino problem, i.e. the observed deficit of solar neutrino fluxes \\cite{SNP}, is a well established tool whose resolution requires (almost without doubt \\cite{doubt}) neutrino physics beyond the (minimal) Standard Model. The acceptable solutions to this problem, in terms of vacuum \\cite{Vsol} or matter \\cite{Msol} flavour oscillations or spin and flavour oscillations \\cite{SFsol} as well as in terms of active-sterile neutrino conversion \\cite{Stsol}, all need non vanishing neutrino masses and mixing \\cite{Vac, MSW, Akh, sterile}.\\\\ The second indication in favour of neutrino oscillations come from the observed deficit of atmospheric muon neutrinos with respect to electron neutrinos \\cite{atmo} that can be explained in terms of \\nm \\rt \\nt or \\nm \\rt \\ne or even active-sterile neutrino transitions \\cite{atsol}.\\\\ Laboratory direct searches for massive neutrinos only give, at present,  upper limits on neutrino masses \\cite{masses} and the same is valid for reactor and accelerator  neutrino oscillations experiments \\cite{reactor}, except for the LSND experiment \\cite{LSND} whose results seem to be explained in terms of $\\ov{\\nm} \\rt \\ov{\\ne}$ oscillations.\\\\ Hints for massive neutrinos also come from cosmology, looking at \\nt as the most probable candidate for the hot component of the dark matter \\cite{HDM} and from the observed peculiar velocities of pulsars \\cite{Segre}. On the other hand, in Grand Unified Theories, which attempt to give a unified view of electroweak and strong interactions, massive neutrinos are predicted \\cite{Buccella} together with other phenomena violating both lepton numbers and baryon number (such as, for example, proton decay).  However, the most intriguing fact is that a simple scenario with only three massive neutrinos cannot account for the solar neutrino problem, the atmospheric neutrino anomaly and the LSND result. This is because the three squared masses differences $\\Delta m^2$ for the three oscillation solutions to these problems are all distinct between them: the resonant MSW solution to the solar neutrino problem requires $\\Delta m^2 \\, \\sim \\, 10^{-5}$ eV$^2$, while for the atmospheric anomaly $\\Delta m^2 \\, \\sim \\, 10^{-2}$ eV$^2$ is needed and $\\Delta m^2 \\, \\sim \\, 1$ eV$^2$ for the LSND result. Many analyses \\cite{analyses} have been conducted for giving a unified view of the three problems in terms of neutrino oscillation (taking into account also the limits from laboratory experiments) and a coherent picture seems to emerge with four massive neutrinos, namely the three known neutrinos plus a sterile neutrino. Note that four neutrino mass eigenstates are needed, but not necessarily more than three neutrino flavour eigenstates. This scenario is easily realized if one considers neutrinos as Dirac-Majorana particles described by the following general mass term in the electroweak lagrangian \\cite{sterile}: \\be -{\\cal L}^{DM}_{m} \\; = \\; \\sum_{l,l^{\\prime}} \\ov{\\nu}_{l^{\\prime}R} \\; M^{D}_{l^{\\prime}l} \\; \\nu_{l^{\\prime}L} \\; + \\; \\frac{1}{2} \\; \\sum_{l,l^{\\prime}} \\ov{\\nu}^{c}_{l^{\\prime}R} \\; M^{1}_{l^{\\prime}l} \\; \\nu_{l L} \\; + \\; \\frac{1}{2} \\; \\sum_{l,l^{\\prime}} \\ov{\\nu}^{c}_{l^{\\prime}L} \\; M^{2}_{l^{\\prime}l} \\; \\nu_{l R} \\; + \\; h.c. \\label{11} \\ee Here $l,l^{\\prime} = e, \\mu , \\tau$ label the three flavour eigenstates and $M_D$, $M_1$, $M_2$ are the Dirac and the two Majorana mass matrices which, in general, are hermitian and non diagonal (however, $M_1$ and $M_2$ have to be symmetric). To construct the mass term in (\\ref{11}) we need the three known left-handed neutrinos (and their antiparticles) and other three right-handed sterile neutrinos (and their antiparticles) \\footnote{Obviously, the generalization to more than three families is possible and straightforward}. After the diagonalization of (\\ref{11}) we can obtain in general six mass eigenstates which are Majorana fields; so in this framework we can easily endow the above scenario with four massive neutrinos coming from the experiments.\\\\ Note that if neutrinos are really described by the mass term in (\\ref{11}), the total lepton number is no longer conserved and peculiar phenomena, as neutrinoless double beta decay and neutrino-antineutrino oscillations can take place.\\\\ We stress that (\\ref{11}) is predicted in many GUTs \\cite{Buccella} in which the popular ``seesaw'' mechanism \\cite{seesaw} can give rise to very small neutrino masses in a very natural way by supposing $M_1 \\approx 0$ and $M_D \\ll M_2$. However, this is not the only possibility; recently some models assuming  $M_1 \\simeq M_2$ have been proposed \\cite{equal} for accounting the three experimental indications on neutrino oscillations discussed above. Here we further explore this last scenario and study flavour transitions of Dirac-Majorana neutrinos from a completely phenomenological point of view, adopting no particular model. This work is a generalization to flavour transitions of previous papers \\cite{TE, previous} in which we studied peculiar oscillations of Dirac-Majorana neutrinos. We now assume, for simplicity, only two flavours, so $M_D$, $M_1$, $M_2$ in (\\ref{11}) are $2 \\times 2$ matrices in the flavour space. In the following section, the basic vacuum oscillations allowed by (\\ref{11}) are studied and transition probabilities are explicitly given in the general case. Some very interesting consequences due to particularly simple forms of the mass matrix are also investigated. In section 3, given the effective hamiltonian of Dirac-Majorana neutrinos interacting with a medium, resonant matter oscillations are considered along with a qualitative discussion of the phenomenon with the aid  of the level crossing diagram. Finally, in section 4, there are our conclusions and remarks. ",
        "conclusions": "In this paper we have studied the propagation both in vacuum and in matter of Dirac-Majorana neutrinos and analyzed active-active (flavour-changing) as well as active-sterile transitions, which are, in general, both possible.\\\\ For vacuum oscillations, in section 2 we have given the general expressions for the transitions probabilities for $ \\nel \\rightarrow \\nml $, \\nel \\rt \\ncml , $ \\nel \\rightarrow \\ncel $ We have then discussed some interesting limiting cases for Dirac ($M_D$) and Majorana ($M_M$) mass matrices. For pure Dirac ($M_M=0$) or pure Majorana ($M_D = 0$) neutrinos obviously we recover the usual flavour oscillation formulae \\cite{Vac}, while for both $M_D$ and $M_M$ non vanishing and diagonal the Pontecorvo formula \\cite{Pontecorvo} for neutrino-antineutrino (active-sterile) oscillations is obtained \\cite{TE}.\\\\ An interesting non trivial case is that with $M_D$ and $M_M$ given by (\\ref{226}) or (\\ref{235}) which implement the idea that neutrino mixing is essentially ruled only by the Dirac mass term while the Majorana mass term is diagonal or vice-versa, respectively. In both cases, neither pure flavour oscillations nor Pontecorvo oscillations are predicted, but only flavour-changing active-sterile transitions, such as $\\nel \\rightarrow \\ncml$, are possible. Remarkably, the transition probability for these is identical in form to that for flavour oscillations for pure Dirac or Majorana neutrinos, and this holds both in vacuum and in matter. For the latter, the resonance condition is only shifted by the neutral current contribution of \\nel to the effective potential. So, for example, the solution to the solar neutrino problem in terms of active-sterile neutrino oscillations proposed in \\cite{Stsol} applies unmodified to the present scheme. \\\\ Another interesting case, even if a bit more complicate, has been analysed for the mass matrices in (\\ref{236}) or (\\ref{237}), which implements the idea that neutrino mixing is given by the Dirac and Majorana mass terms with the same strength. In this case, $ \\nel \\rightarrow \\nml $, $ \\nel \\rightarrow \\ncml $, $ \\nel \\rightarrow \\ncel $ transitions are all possible and, in vacuum, the first two have the same transition probability, which is also identical in form to that obtained in the previous case, except for a constant suppression factor in the amplitude of oscillations of 1/4. Also in matter the pattern of neutrino transitions present in the general case is (qualitatively) reproduced in this peculiar scheme. In particular, all the flavour changing oscillations can be resonantly amplified while Pontecorvo \\nel \\rt \\ncel matter oscillations have maximum amplitude only for a given electron to neutron number ratio (see eq. (\\ref{313}) and the related footnote); the resonance conditions were discussed in section 3.1 . \\\\ Given the multiresonance structure of the oscillations pattern, it is then interesting to follow the evolution of a \\nel , for example, in a varying density medium such as the Sun; this has been done in section 3.4 with the help of a level crossing diagram reported in Fig. 1. Several scenarios are possible according to the adiabaticity properties of level crossing near the resonance points. In particular, starting from a pure \\nel beam at high density, to have a consistent conversion into \\nml at low density the resonance for \\nel \\rt \\ncml has to be crossed non adiabatically, while the passage through the one for \\nel \\rt \\nml has to be adiabatic. However, we have also shown that at very low density, and then in the vacuum, it is more appropriate to deal with the Majorana combinations $\\wt{n_\\pm}$ in (\\ref{212}) than with the pure flavour states \\nel , \\ncel , \\nml , \\ncml . This is strictly related to the Dirac-Majorana nature of neutrinos, which chooses Majorana eigenstates instead of pure flavour states as starting points. In this respect, we have to deal with ``generic'' flavour-changing or flavour-conserving transitions of Dirac-Majorana neutrinos without looking at the particular active neutrino or sterile antineutrino state. It is through the weak interactions, with which neutrinos are produced and detected, that a particular (active or sterile) component of the Majorana eigenstates is chosen.\\\\ The results obtained for the case of degenerate Dirac-Majorana mixing are qualitatively valid also in the general case in which all the entries of the Dirac and Majorana mass matrices are non zero and different between them. The main difference between the two cases is that in the general framework there are 3 mass parameters and two mixing angles ruling the evolution, while for the particular case studied in sections 2.3 and 3.4 there are only 2 mass parameters and 1 mixing angle (these parameters being not completely independent, because of relation (\\ref{239b})). The presence of more degrees of freedom in the general case allows to consider some peculiar situations which are not possible otherwise. The most remarkable one is that in the general case the proportionality of the $\\Sigma$ parameter to $\\sin^2 \\, 2 \\theta_+$ (see eq. (\\ref{239b})) is lost, so that the structure of the level crossing diagram at very low density can be altered. The eigenvalues of $H_m$ in (\\ref{39}) for zero density (vacuum) are given by \\be \\frac{1}{8k} \\, \\left( \\pm \\, 2 \\Delta m_+^2 \\; + \\; \\Sigma \\right) \\ee \\be \\frac{1}{8k} \\, \\left( \\pm \\, 2 \\Delta m_-^2 \\; - \\; \\Sigma \\right) \\ee so that one can manipulate the mass parameters to modify the low density region of the level crossing diagram without grossly altering the region where the resonance points are present. In any case, there can be present no substantial modifications of the conclusions reached above.  The oscillations of Dirac-Majorana neutrinos here studied with their peculiar features can be efficiently tested in astrophysics, in particular detecting solar or supernova neutrinos, and can have even profound implications in cosmology for the nucleosynthesis of light elements in the Universe.   \\vspace{1truecm}  \\noindent {\\Large \\bf Acknowledgements}\\\\ \\noindent We express our sincere thanks to Prof. F. Buccella for very useful talks and his unfailing encouragement, and to Prof. E. Kh. Akhmedov for enlightening discussions with one of us (S.E.)."
    },
    "9803/astro-ph9803234_arXiv.txt": {
        "abstract": "A new cataclysmic variable is identified as the optical counterpart of the faint and hard X-ray source RX\\,J0757.0+6306 discovered  during the ROSAT all-sky survey. Strong double-peaked emission lines bear evidence of an accretion disc via an S--wave which varies with a period of 81$\\pm 5$~min. We identify this period as the orbital period of the binary system. CCD  photometry reveals an additional period of 8.52$\\pm 0.15$ min. which was stable over four nights. We  suggest that \\rxj\\ is  possibly an intermediate polar, but we cannot exclude the possibility that it is a member of the  SU UMa group of dwarf novae. ",
        "introduction": "Within a project for the optical identification of a complete sample of 674 northern ROSAT All-Sky Survey (RASS) X-ray sources (which is a collaboration between the Max-Planck-Institut f\\\"ur extraterrestrische Physik, Garching, the Landessternwarte Heidelberg, Germany, and the Instituto Nacional de Astrof\\'isica, Optica y Electronica, Mexico) several new cataclysmic variables were identified. A detailed description of the project is given by Zickgraf \\etal\\ (1997). The full catalogue with all identifications is published in Appenzeller \\etal\\ (1998). Here we report the identification of the RASS X-ray source RX\\,J0757.0+6306 (= 1RXS J075700.5+630602).  Cataclysmic variables (CVs) are close binary systems with a white dwarf primary  accreting matter supplied by a late type main-sequence secondary star via an accretion disc or along magnetic field lines of the white dwarf. Magnetic CVs, where the white dwarf has a sufficiently strong magnetic field to affect the accretion trajectory, form two distinctive subclasses: the high-field polars, and the low-field intermediate polars (IPs). These subclasses are characterized by well-defined observational properties (Cropper 1990; Patterson 1994; Warner 1995). The polars are usually soft X-ray emitters and have near synchronously rotating WD, the IPs are harder X-ray sources and show a second periodicity due to the asynchronously rotating WD. In some cases a third period is observable, which is interpreted as the beat period between orbital  and spin periods.  Besides differences in the flux distribution and variability, the orbital period distribution of the various subclasses of CVs were also noticed to be different (Kolb 1995). Polars tend to cluster below the period gap (2 h $< P_{orb}<$ 3 h), while IPs are preferentially above the gap. Non-magnetic CVs are distributed almost equally. All subclasses however show deficiency of systems in the period gap and a short-period cutoff at $\\rm P_{min}=80$ min (the minimum period). The statistically significant properties of the period distribution are presumed  to have  an evolutionary origin (Verbunt \\& Zwaan 1981, Verbunt 1984, King 1988, Kolb and Ritter 1992). The rapidly increasing number of magnetic CVs discovered from the  ROSAT data has a significant impact on the above mentioned distribution and its consequences. ",
        "conclusions": "A new cataclysmic variable is discovered with  interesting features:  \\begin{enumerate} \\item The orbital period of $81\\pm 5$ min puts \\rxj\\ near the hydrogen burning period minimum where CVs experience a turning point of their evolution. Large flickering in the optical light curve and the observed spectral features of the object certainly show the presence of an accretion disc.  \\item The limited search in the Sonneberg all-sky patrol plates revealed that the system undergoes outburst activity. Another outburst was recorded (vsnet-alert No. 1379) shortly after the object's discovery was announced through the VSNET (vsnet-chat No 662). From the plate statistics we can assume that the system has rather frequent outbursts. The amplitude of the outbursts of about 4 mag are typical for dwarf novae systems, but not as large as in SU~UMa superoutbursts or the so called TOADs (tremendous outburst amplitude dwarf novae; Howell \\etal\\ 1995).  \\item There are periodic light variations with a period of  8.5 min in the light curve of the \\rxj. We observed them directly on four  out of five occasions. In the fifth night a periodic signal with a side-band frequency was detected in the power spectrum.  Very recently, the 8.5\\,min period was confirmed by R. Fried (vsnet-alert No 1387) from more prolonged observations.  \\end{enumerate}  Thus, \\rxj\\ shows mixed characteristics, making its type classification uncertain. From purely spectroscopic characteristics one may conclude that the new CV is a dwarf nova. Its short orbital period suggests that instead it may belong to the SU~UMa class or TOADs. But the repetitive detection of high-frequency pulses with a clearly fixed period indicates that it deserves a classification as an intermediate polar. This still needs to be confirmed by checking the coherency of the photometric pulses and by the detection of X-ray pulses. Intermediate polars  are CVs with the primary white dwarf rotating asynchronously due to its moderate magnetic field. Column accretion onto the magnetic poles results in the emission of high-energy radiation. This radiation is reprocessed elsewhere in the system into optical light which is modulated at periods shorter than the orbital period. The optical modulation can track the spin and/or the spin/orbit beat period of the binary (see the review by Patterson 1994).   The presence of X-ray emission in the quiescent state of  \\rxj\\ along with the moderate He~{\\sc ii} 4686 \\AA\\ emission also argue in favor of a magnetic nature. The survey of non-magnetic CVs by van Teeseling \\etal\\ (1996) shows that the majority of X-ray emitting dwarf novae are of the SU UMa type, but they all are softer sources (HR1$\\le 0.7$) than \\rxj. There are a few long-period objects classified as non-SU UMa variables that are as hard as \\rxj.  These belong to the VY~Scl, Z~Cam, and UX~UMa subclasses. We do not have any evidence which support a classification of \\rxj\\ as any of these types. Hence, since the rest of the CVs which are X-ray sources are magnetic, we conclude that \\rxj\\ is most probably magnetic.  Only two IPs have been detected with EUVE and only a few dwarf novae, all of the latter during outburst. AM Herculis stars, particularly those with high magnetic fields are  detected using EUVE. Thus, the non-detection of \\rxj\\ with EUVE does not prove that it is not an IP. It may indicate that \\rxj\\ could have a weak magnetic field (B $< 8$ MG), but given that only two IPs have EUVE detections and because of the rather large distance of \\rxj\\ the non-detection is not considered unusual. On the other hand, the intensity of the He{\\sc~ii} emission is not high enough to unambiguously classify it as a magnetic system. Silber (1992) set the following criteria for magnetic CVs: $20<$EW~(H$\\beta)<40\\AA$ and He{\\sc ii}/H$\\beta>0.4$. In our case, if the  larger equivalent width could be attributed to a shorter orbital period, the He{\\sc ii}/H$\\beta$ ratio is definitely below this criterium ($\\approx~0.15$).  The existence of outbursts and the short orbital period of the system is in some discordance with the IP classification. Most IPs cluster above the 2--3 hour period gap, while short period magnetic CVs are usually polars. However, a weak field IP will remain an IP even when it evolves towards shorter periods. In IPs, accretion outside of the Alfven radius remains in the form of a disc, while accretion inside the Alfven radius is dominated by flow along magnetic field lines. Outburst activity is uncommon since the inner part of the disc is disrupted by the magnetic field.  Nevertheless, neither outburst  activity nor short orbital period exclude the possibility of \\rxj\\ to be classified as an IP. In addition, the presence of a large disc in a short period magnetic CV suggests that the magnetic field is weak. Otherwise it would be a polar. For such a weak field  case it may not be surprising that \\rxj\\ appears to be an IP with some properties (i.e. outbursts) similar to non-magnetic CVs, yet the evidence that it is a magnetic CV is compelling. Therefore, we offer \\rxj\\ as a candidate for the shortest period intermediate polar."
    },
    "9803/astro-ph9803002_arXiv.txt": {
        "abstract": "We show that pulsar velocities may arise from anisotropic neutrino emission induced by resonant conversions of massless neutrinos in the presence of a strong magnetic field. The main ingredient is a small violation of weak universality and neither neutrino masses nor magnetic moments are required. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803016_arXiv.txt": {
        "abstract": "The $^9$Be\\,{\\sc ii} $\\lambda$ 3131 \\AA\\ doublet has been observed in the solar-type stars 16 Cyg A \\& B and in the late G-type star $\\rho^1$ Cnc, to derive their beryllium abundances. 16 Cyg A \\& B show similar (solar) beryllium abundances while 16 Cyg B, which has been proposed to have a planetary companion of $\\sim 2$ $M_{\\rm Jup}$, is known to be depleted in lithium by a factor larger than 6 with respect to 16 Cyg A. Differences in their rotational histories which could induce different rates of internal mixing of material, and the ingestion of a similar planet by 16 Cyg A are discussed as potential explanations. The existence of two other solar-type stars which are candidates to harbour planetary-mass companions and which show lithium and beryllium abundances close to those of 16 Cyg A, requires a more detailed inspection of the peculiarities of the 16 Cyg system.  For $\\rho^1$ Cnc, which is the coolest known object candidate to harbour a planetary-mass companion ($M > 0.85$ $M_{\\rm Jup}$), we establish a precise upper limit for its beryllium abundance, showing a strong Be depletion which constrains the available mixing mechanisms. Observations of similar stars without companions are required to asses the potential effects of the planetary companion on the observed depletion. It has been recently claimed that $\\rho^1$ Cnc appears to be a subgiant. If this were the case, the observed strong Li and Be depletions could be explained by a dilution process taking place during its post-main sequence evolution. ",
        "introduction": "\\label{sec1}  In very recent years, several stars have been proposed to have planetary companions on the basis of measured precise radial velocity variations. This field of research is experiencing rapid development, and updated reviews of the present situation can be found in the proceedings of the workshop on {\\it Brown Dwarfs and Extrasolar Planets} edited by Rebolo et al. (1998) and in The Extrasolar Planets Encyclopaedia\\footnote{http://wwwusr.obspm.fr/planets/} by J. Schneider.  Once a solar-type star has been suggested to harbour a planetary-mass companion, it is interesting to investigate any similarities with the Sun, as well as to find possible differences with respect to other single stars. Chemical abundances are among the most important parameters to be compared and, in particular, precise abundances of light elements such as lithium and beryllium (easy to destroy by $(p,\\alpha)$ nuclear reactions when the temperature reaches $\\sim 2.5\\times 10^6$ and $\\sim 3.5\\times 10^6$ K, respectively) combined with the abundances of other elements which are not so readily destroyed in stellar interiors, should help to understand how the presence of planets may affect the chemical composition of their parent stars. Gonzalez (1997, 1998) has derived the overall metallicities as well as abundances of different elements (including lithium) for a wide sample of proposed parent stars, finding that four of the known systems show  a metallicity significantly higher than the solar value.  A peculiar system such as 16 Cyg A \\& B, formed by twin solar-type stars of which only one has an orbiting planet (Cochran et al. 1997), is an especially suitable candidate to perform a detailed abundance study. Gonzalez (1998) found that both stars have a similar metallicity with a value slightly larger than solar, and confirmed independently a previous result of King et al. (1997a) that 16 Cyg B (the star with a suspected planet) is strongly depleted in lithium with respect to 16 Cyg A. The knowledge of their beryllium abundances is of potential value in quantifying the possible influence of a planetary companion on the mixing mechanisms operating in the stellar interior.  $\\rho ^1$ Cnc is a star with spectral type G8V, and is the coolest known object which is a candidate to have a planetary companion. Following Gonzalez (1998), this star falls into the group having roughly Jupiter-mass companions with small circular orbits and very metal-rich parent stars. Dominik et al. (1998) have shown recently that the planetary system of $\\rho ^1$ Cnc also hosts a Vega-like disk of dust, evidenced by an infrared excess at 60 $\\mu$m. The star is very depleted in lithium and its beryllium abundance could be compared with existing upper limits measured in younger stars with similar effective temperatures (Garc\\'\\i a L\\'opez et al. 1995a).  In this paper we derive the beryllium abundances of the 16 Cyg system and of $\\rho ^1$ Cnc by comparing observations with spectral syntheses of the $^9$Be\\,{\\sc ii} $\\lambda$ 3131 \\AA\\ doublet. We use those, together with their published lithium values, as well as with available abundances for other stars (with and without suggested planetary companions), and discuss briefly possible effects of planets on processes taking place in their structure and evolution. ",
        "conclusions": "\\label{sec5}  Beryllium abundances have been derived for the solar-like stars 16 Cyg A \\& B and the cooler object $\\rho ^1$ Cnc, for which there are published values of their lithium abundances. 16 Cyg B and $\\rho ^1$ Cnc are candidates to be parents of extrasolar planets, and by measuring their Be abundances we aim at studying the potential dependence on the presence of planetary companions of detailed processes operating in their structure and evolution.  16 Cyg A \\& B show very similar Be abundances, which are compatible with the solar value, while the lithium abundance of 16 Cyg B is at least a factor 6 smaller than that of 16 Cyg A. Different rates of mixing of material in their interiors associated with different angular momentum histories, as well as the hypothetical ingestion of a planetary companion by 16 Cyg A are discussed as potential explanations. The existence of two other solar-like parent stars, whose Li (and Be) does not show strong depletion, i.e. whose behaviour is like 16 Cyg A, the Sun and the majority of similar stars with Li and Be abundances available, implies that the 16 Cyg system requires special observational and theoretical attention.  A low upper limit has been derived for the beryllium abundance of $\\rho ^1$ Cnc. This is the first time a precise limit has been set and that such strong Be depletion has been observed in a late G-/early K-type MS star. This measurement clearly constrains the depletion predictions of the available mixing mechanisms, but requires observation of planet-free stars with similar age and spectral type to discard the potential effects of the planetary companion on the Li and Be depletions. Claims have also been made indicating that $\\rho ^1$ Cnc appears to be a subgiant. If this were the case, its strong Li and Be depletions could be explained by a dilution process taking place during its post-MS evolution."
    },
    "9803/astro-ph9803199_arXiv.txt": {
        "abstract": "We present a study of the polarizing power of the dust in cold dense regions (dark clouds) compared to that of dust in the general interstellar medium (ISM). Our study uses new polarimetric, optical, and spectral classification data for 36 stars to carefully study the relation between polarization percentage ($p$) and extinction ($A_V$) in the Taurus dark cloud complex.   We find two trends in our $p-A_V$ study: (1) stars background to the warm ISM show an increase in $p$ with $A_V$; and (2) the percentage of polarization of stars background to cold dark clouds does not increase with extinction. {\\it We detect a break in the $p-A_V$ relation at an extinction $1.3 \\pm 0.2$ mag, which we expect corresponds to a set of conditions where the polarizing power of the dust associated with the Taurus dark clouds drops precipitously. This breakpoint places important restrictions on the use of polarimetry in studying interstellar  magnetic fields.} ",
        "introduction": "The polarization of background starlight has been used for nearly half a century to probe the magnetic field  direction in the interstellar medium (ISM).  The observed polarization  is believed to be caused by dichroic  extinction of background starlight passing through concentrations of aligned elongated dust grains along the line-of-sight. Although there is no general consensus on which is the dominant grain alignment mechanism (Lazarian, Goodman \\& Myers 1997), it is  generally believed that the shortest axis of the ``typical'' elongated grain tends to be become  aligned to the local magnetic field. For this orientation, the observed polarization vector is parallel to the plane-of-the-sky projection of a line-of-sight-averaged magnetic field (Davis \\& Greenstein 1951).  The line-of-sight averaging inherent in background starlight polarimetric observations can make interpretation of the polarization produced by different field orientations and/or several independent dust clouds very complicated. Nonetheless, it was thought that if lines of sight with just one localized very dusty region (such as a dark cloud) between us and a background star could be found, surely the polarization would reveal the field associated with that dusty region. However, recent studies in the Taurus region (Goodman et al. 1992; Gerakines et al. 1995) and other parts of the sky (Creese et al. 1995; Goodman et al. 1995) have uncovered substantial evidence to show that dust inside cold dark clouds has lower polarizing power than dust in the general warm ISM. This means that the polarization of the light from a background star is a non-uniformly {\\it weighted} line-of-sight average of the projected plane-of-the-sky field, and that grains in cold dark clouds are systematically down-weighted. The ultimate implication of this down-weighting is that above some (column?) density threshold, the polarization of background starlight gives no information about the magnetic field in dark clouds.  It is the goal of this Letter to find and physically describe this threshold.  The evidence that grains in cold dark clouds are inefficient polarizers of background starlight is multi-faceted.  Eight years ago, Goodman et al. (1990) found that the smooth large-scale patterns apparent in polarization maps of dark cloud complexes (e.g. Vrba et al. 1976; Vrba et al. 1981; Moneti  et al. 1984; Whittet et al. 1994) do not systematically change in response to the large density enhancements represented by the dark clouds.  After this realization, it was hypothesized that perhaps optical polarimetry was incapable of seeing field changes which might occur only in the high-density, optically opaque, interiors of dark clouds.  So, near-infrared polarimetry, which can probe the optically opaque cloud interiors was undertaken.  The near-infrared observations showed that the mean direction and dispersion of the polarization vectors are virtually {\\it identical} in the cloud interiors and their peripheries (Goodman et al. 1992; 1995).  Thus, the presence of cold dark clouds still appeared to have no geometric effect on the polarization maps, implying either that: 1) the field is truly unaffected by the cloud; or 2) that background starlight polarimetry is somehow insensitive to the field in dark clouds. Polarization-extinction relations provide the best discriminant between these hypotheses.  For grains of constant polarization efficiency, $p$ should rise with $A_V$.  Using the near-infrared observations, Goodman et al. (1992, 1995) find that the {\\it percentage of polarization does not rise with extinction} in cold dark clouds.  The simplest interpretation\\footnote{Note that increased field tangling inside dark clouds cannot explain the near-infrared polarimetric observations for two reasons. 1.) The dispersion in the distribution of position angle does not increase in the cloud interior (near-IR observations) relative to the periphery (optical observations).  And, 2.) while it is true that the slope of a $p-A_V$ relation will diminish due to field tangling, it will remain positive even for highly tangled fields if all grains polarize equally well (see Jones 1989; Jones et al. 1992).}  of this result is that dust in dark clouds adds plenty to the observed extinction, but has very little ``polarizing power\" and so adds only a very small fraction to the observed net polarization. A number of factors, including poor grain alignment, grain growth, and/or changes in grain shape or composition, could be responsible for the low polarization efficiency exhibited by dust grains in cold dark clouds (see Goodman et al. 1995). Regardless of which factor(s) cause(s) the low polarization efficiency exhibited by by dust in dark clouds, the fact is that background starlight polarimetry does not reliably reveal the magnetic field {\\it in} dark clouds.  Based on the near-infrared studies, we expect that the fraction of grains with high polarization efficiency  is relatively constant in the lower-density warm ISM, but drops precipitously in dark clouds.  Therefore, we hypothesize that a breakpoint in the $p-A_V$ relationship might exist at the dark clouds' ``edges\", which previous studies in the near-IR   (Goodman et al. 1992; 1995; Gerakines et al. 1995) could not detect, due to their inability to measure low $A_V$'s accurately enough. In this Letter, we present our attempt to carefully study the $p-A_{V}$ relation near dark clouds, and thus offer a set of guidelines as to where the polarization maps can be taken as faithful representations of the magnetic field projected onto the plane of the sky, and where they cannot. ",
        "conclusions": "The breakpoint in the $p-A_{V}$ relation places important restrictions on the use of polarimetry in studying interstellar magnetic fields. Since the polarization efficiency of the dust inside dark clouds is very low, most of the polarization observed for lines of sight which pass through these extinction peaks is not due to the dark cloud; it is due to dust background and foreground to the cloud. Hence, one should not use background starlight polarimetry to map magnetic fields inside dark clouds. With the results of this study we can quantify the word ``inside.\" In regions like Taurus, {\\it it is safe to interpret the polarization of background starlight as a representation of the plane-of-the-sky projected magnetic field up to the 1.3 $\\pm$ 0.2 mag ``edge'' of the dark cloud}.  In other words the linear relation between $p$ and $A_V$ that exists in the low-density ISM breaks down for stars background to the $\\simgreat 1.3$ mag of extinction produced by a dense localized dusty region (i.e., dark cloud).  After this edge polarization no longer rises with extinction, and thus cannot reveal the field structure in the dense region. We restate that this proscription  only applies for stars background to cold dark clouds, as stars background to the warm ISM have not been shown to exhibit such behavior."
    },
    "9803/astro-ph9803150_arXiv.txt": {
        "abstract": "We investigate the effect of gravitational lensing by matter distribution in the universe on the cosmic microwave background (CMB) polarization power spectra and temperature-polarization cross-correlation spectrum. As in the case of temperature spectrum gravitational lensing leads to smoothing of narrow features and enhancement of power on the damping tail of the power spectrum. Because acoustic peaks in polarization spectra are narrower than in the temperature spectrum the smoothing effect is significantly larger and can reach up to 10\\% for $l<1000$ and even more above that. A qualitatively new feature is the generation of $B$ type polarization even when only $E$ is intrinsically present, such as in the case of pure scalar perturbations. This may be directly observed with Planck and other future small scale polarization experiments. The gravitational lensing effect is incorporated in the new version (2.4) of CMBFAST code. ",
        "introduction": "Over the next few years a number of ground based, balloon and satellite experiments will measure CMB sky with an unprecedented accuracy and detail. The promise of a one percent precision on the measured power spectrum of CMB anisotropies requires a similar accuracy in the theoretical predictions, if we are to exploit all the information present in the data. The rewards will be rich: among other things this will allow an accurate determination of a number of cosmological parameters and testing of current structure formation theories \\cite{parameters}. In principle such a program is possible, since the anisotropies were produced when the universe was still in the linear regime, which makes the calculations of model predictions very accurate. In practice there are a number of important effects that need to be included if this goal is to be realized. One of the most important among these is the gravitational lensing effect.  As photons propagate through the universe from their last scattering to our detectors they are randomly deflected by the gravitational force exerted upon them by the inhomogeneous mass distribution. Previous work has shown that gravitational lensing has an  effect on the temperature anisotropy power spectrum which is not insignificant \\cite{others,uroslens}. The random deflections smear out the sharp features in the correlation function or power spectrum, leading to a suppression of acoustic oscillations. Gravitational lensing can also enhance power on the damping tail, causing it to decay less rapidly than predicted on very small angular scales \\cite{bentonsilk}. Gravitational lensing effect on the temperature anisotropies has been discussed several times in the literature and the formalism to calculate it using the evolution of density power spectrum both in linear and nonlinear regime has been presented in \\cite{uroslens}. In this paper we extend this calculation to the two linear polarization power spectra and to the cross-correlation spectrum between temperature and polarization. Because acoustic oscillations are narrower for polarization spectra than for temperature, one expects gravitational lensing effect to be more significant in the former and indeed our results confirm this. In addition, a qualitatively new effect is the mixing between $E$ and $B$ types of polarization, which changes the pattern of polarization. The outline of the paper is the following. In \\S \\ref{formalism} we develop the formalism: this section contains all the main analytic expressions needed for a numerical implementation of the effect. These have been numerically implemented in the new version of CMBFAST package (version 2.4) and require only a marginal increase in the CPU time for their evaluations. In \\S \\ref{estimate} we compute the effect for a typical cosmological model and address the question of direct observability of the effect. We present the conclusions in \\S \\ref{conclusions}. ",
        "conclusions": ""
    },
    "9803/astro-ph9803293_arXiv.txt": {
        "abstract": "Optical multi-slit spectra have been obtained for 47 globular clusters surrounding the brightest Virgo elliptical NGC~4472 (M49). Including data from the literature, we analyze velocities for a total sample of 57 clusters and present the first tentative evidence for kinematic differences between the red and blue cluster populations which make up the bimodal colour distribution of this galaxy. The redder clusters are more centrally concentrated and have a velocity dispersion of 240 kms$^{-1}$ compared with 320 kms$^{-1}$ for the blue clusters. The origin of this difference appears to be a larger component of systematic rotation in the blue cluster system. The larger rotation in the more extended blue cluster system is indicative of efficient angular momentum transport, as provided by galaxy mergers. Masses estimated from the globular cluster velocities are consistent with the mass distribution estimated from X-ray data, and indicate that the M/L$_B$ rises to $50$~M/L$_\\odot$ at 2.5 R$_e$. ",
        "introduction": "The study of extragalactic globular cluster systems can provide important clues to the formation history of their host galaxies.  This is particularly true for elliptical galaxies for which there are two currently popular paradigms. One paradigm is the standard monolithic collapse model in which elliptical galaxies form in a single burst of star formation at high redshift (e.g. Arimoto \\& Yohii 1987). In contrast, hierarchical structure formation theories predict that spheroidal galaxies form continuously through a sequence of galaxy mergers (e.g. Cole \\etal 1994; Kauffmann 1996).  Ashman \\& Zepf (1992) explored the properties of globular clusters in models in which elliptical galaxies are the products of the mergers of spiral galaxies, and showed that the greater specific frequency of globular clusters around ellipticals relative to spirals could be explained if globular clusters form during the mergers. They also predicted that elliptical galaxies formed by mergers will have two or more populations of globular clusters - a metal-poor population associated with the progenitor spirals, and a metal-rich population formed during the merger.  In contrast, monolithic collapse models naturally produce unimodal metallicity distributions.  The discovery that the globular cluster systems of several elliptical galaxies have bimodal colour (and by implication metallicity) distributions (Zepf \\& Ashman 1993; Whitmore \\etal 1995;  Geisler \\etal 1996) provides strong support for the merger model. Geisler \\etal (1996) and Lee et al. (1998) also show that the red (metal-rich)  cluster population is more centrally concentrated than the blue (metal-poor) population, as predicted by Ashman \\& Zepf (1992).  Recently, an alternative view has been presented by Forbes et al.\\ (1997), who suggest that the bimodal color distributions may not be due to mergers, but to a multi-phase single collapse. Although the primary physical mechanism known to produce distinct formation episodes is mergers, it is important to attempt to distinguish between these competing models for the formation of globular cluster systems and their host elliptical galaxies. The kinematics of globular cluster systems may offer such a test of these models. In the multi-phase collapse picture, angular momentum conservation requires that the spatially concentrated metal-rich population rotates more rapidly than the extended metal-poor population. In contrast, simulations of merger models indicate that mergers typically provide an efficient means of angular momentum tranfer, and that the central regions have specific angular momentum that is lower than the outer regions (Hernquist 1993; Heyl, Hernquist and Spergel 1996).  Studies of the kinematics of globular cluster systems therefore provide important constraints on the formation history of elliptical galaxies. They also provide useful probes of the mass distribution of elliptical galaxies at radii larger than can be reached by studies of the integrated light. The extended nature of globular cluster systems allows the dynamical mass determined from their velocities to be compared at similar radii to masses determined through studies of the hot X-ray gas. A recent example is the study of the M87 globular cluster system by Cohen \\& Ryzhov (1997), who find that a rising mass-to-light ratio is required out to radii of $\\sim 3 R_e$, in agreement with X-ray mass determinations. However, M87 occupies a privileged position at the center of the Virgo cluster, so it is critical to test whether the rising mass-to-light ratio, and the agreement with X-ray masses, is true for more typical cluster elliptical galaxies.  In this paper we present a spectroscopic study of the globular cluster system of the elliptical galaxy NGC~4472 (M49). This is the brightest elliptical galaxy in the Virgo cluster and has been the subject of a detailed photometric study by Geisler \\etal (1996). The only previously published spectroscopic data for the NGC~4472 globular cluster population is by Mould \\etal (1990) who presented velocities and line strengths for 26 clusters.  The outline of our paper is as follows: Section 2 discusses the sample selection together with our observations and data reduction; Section 3 discusses the kinematics of the metal-rich and metal-poor populations in the context of the merger model, and analyses the implications for the overall M/L ratio in NGC~4472. Finally, we present our conclusions in Section 4. ",
        "conclusions": "We have made a detailed spectroscopic study of the globular cluster system of NGC~4472, and have more than doubled the number of confirmed clusters to 57. Whilst this remains a statistically small sample, the data show several interesting properties when combined with the accurate colour/metallicity data from Geisler \\etal (1996). When the complete sample is divided into a metal-rich (47\\%) and a metal-poor (53\\%) subset on the basis of their bimodal colour histogram, the metal-poor subset appears to have a broader distribution of velocities. We have investigated this further, and conclude that the most likely cause is a higher mean level of rotation in the metal-poor cluster system, which is consistent with that of the underlying stellar halo in amplitude and position angle (but with a much higher specific angular momentum). The metal-rich clusters on the other hand show only weak evidence for any rotation, and about an axis which is tilted $\\sim 50^o$ from that of the other components. These results are qualitatively in agreement with the predictions of a model in which the metal-rich clusters are formed during the merger of two massive gas-rich galaxies, each with its own old metal-poor cluster population. The cluster system of NGC~4472 forms a dynamically hotter population than the stellar halo, but is consistent with being in dynamical equlibrium with the halo potential defined by the hot X-ray emitting plasma, and supports the presence of a dark $\\sim 10^{12}{\\rm M}_\\odot$ halo in this giant elliptical galaxy."
    },
    "9803/hep-ph9803309_arXiv.txt": {
        "abstract": "The two possible Mikheyev-Smirnov-Wolfenstein (MSW) solutions of the solar neutrino problem (one at small and the other at large mixing angle), up to now tested mainly through absolute neutrino flux measurements, require flux-independent tests both for a decisive confirmation and for their discrimination. To this end, we perform a joint analysis of various flux-independent observables that can be measured at the SuperKamiokande and Sudbury Neutrino Observatory (SNO) experiments. In particular, we analyze the recent data collected at SuperKamiokande after 374 days of operation, work out the corresponding predictions for SNO, and study the interplay between SuperKamiokande and SNO observables. It is shown how, by using only flux-independent observables from SuperKamiokande and SNO, one can discriminate between the two MSW solutions and separate them from the no oscillation case. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803250_arXiv.txt": {
        "abstract": "This study shows one important effect of preexistent cosmic microwave background temperature fluctuations on the determination of the Hubble constant through Sunyaev-Zel'dovich effect of clusters of galaxies, especially when coupled with the gravitational lensing effect by the same clusters. The effect results in a broad distribution of the apparent Hubble constant. The combination of this effect with other systematic effects such as the Loeb-Refregier Effect seems to provide an explanation for the observationally derived values of the Hubble constant currently available based on the Sunyaev-Zel'dovich effect, if the true value of the Hubble constant is $60-80~$km/s/Mpc. It thus becomes possible that the values of the Hubble constant measured by other techniques which generally give a value around $60-80~$km/s/Mpc be reconciled with the SZ effect determined values of the Hubble constant, where are systematically lower than others and have a broad distribution. ",
        "introduction": "The Sunyaev-Zel'dovich (SZ) effect of a cluster of galaxies on the cosmic microwave background (CMB) photons can be used to determine the distance to the cluster hence the Hubble constant ($H_0$), when analysed in conjunction with X-ray observations of the cluster (Cavaliere, Danese, \\& De Zotti 1977; Gunn 1978; Silk \\& White 1978; Birkinshaw 1979). For an excellent recent review on this subject and other SZ related topics, see  Rephaeli (1995 and references therein). The accuracy of the Hubble constant determination depends upon the accuracy of several assumptions involving both sets of observations (radio and X-ray). Perhaps among the most important are the assumptions of sphericity, isothermality of clusters of galaxies (e.g., Inagaki \\etal 1995). In this {\\it Letter} we point out a completely separate effect on the determination of the Hubble constant due to preexistent, small-amplitude CMB temperature fluctuations before the photons undergo the SZ effect through a cluster. The effect is significantly amplified by the gravitational lensing of the CMB photons by the cluster, because the SZ observational beam size is typically comparable to Einstein radius of the source-lens system. This effect, when coupled with some systematic effects such as the one proposed by Loeb \\& Refregier (1997) due to the systematic over-removal of background point radio sources in the beam, may provide an explanation for the observed distribution of $H_0$ determined by SZ effect. ",
        "conclusions": "We show that the background CMB fluctuations, especially when they are coupled with the gravitational lensing effect by clusters of galaxies, have one important effect on the determination of the Hubble constant through Sunyaev-Zel'dovich effect of the clusters (for the adopted set of characteristic numbers for the cluster, which seem fairly realistic compared to those of real clusters of interest): a broad distribution of the apparent Hubble constant is produced with a FWHM about $30\\%$ of the apparent mean value. The combination of this effect with other systematic effects such as the Loeb-Refregier Effect seems to provide a reasonable explanation for the observationally derived values of the Hubble constant currently available, if the true value of the Hubble constant is $\\sim 65~$km/s/Mpc. Thus, it becomes possible that the values of $H_0$ measured by other techniques which generally give a value around $60-80$km/s/Mpc [e.g., $73\\pm 10~$km/s/Mpc ($1\\sigma$) from Freedman, Madore, \\& Kennicutt 1997 based on HST observations of Cepheids; $64\\pm 6~$km/s/Mpc ($1\\sigma$) from Riess, Press, \\& Kirshner 1996 based on type Ia supernova multicolor light-curve shapes; $64\\pm 13~$km/s/Mpc ($95\\%$ confidence level) from Kundic \\etal 1997 based on gravitational lensing time delay measurements; $70\\pm 5~$km/s/Mpc ($1\\sigma$) from Giovanelli 1997 using I-band Tully-Fisher relation; $81\\pm 6~$km/s/Mpc ($1\\sigma$) from Tonry 1997 using surface brightness fluctuations] be reconciled with the SZ effect determined values of $H_0$.  It may be possible, at least in principle, that one can use a large sample of SZ measured Hubble constant to infer the fluctuations of the CMB at the relevant scales, when the Hubble constant is independently measured to high accuracy by methods such as that using detached eclipsing binaries (Paczynski 1997)."
    },
    "9803/astro-ph9803299_arXiv.txt": {
        "abstract": "{A mechanism of ultra-high energy cosmic ray acceleration in extragalactic radio sources, at the interface between the {\\it relativistic} jet and the surrounding medium, is discussed as a supplement to the shock acceleration in `hot spots'. Due to crossing the tangential discontinuity of the velocity the particle can gain an amount of energy comparable to the energy gain at the shock crossing. However, the spectrum of particles accelerated at the jet side boundary is expected to be much flatter than the one formed at the shock. Due to this fact, particles accelerated at the boundary can dominate the overall spectrum at highest energies. In conditions characteristic to extragalactic jets' terminal shocks, the mechanism naturally provides the particles with $E \\sim 10^{20}$ eV and complies with the efficiency requirements. The spectrum formation near the cut-off energy due to action of both the shock acceleration and the tangential discontinuity acceleration is modelled with the Monte Carlo particle simulations. It confirms that the upper energy limit can surpass the shock acceleration estimate.} ",
        "introduction": "Jet-like outflows are observed in a number of astrophysical environments, starting from young stars embedded in their parent molecular clouds, up to active extragalactic objects. In the later case, a number of interesting observational phenomena are noted over orders of magnitude of linear sizes. In particular, at the smallest mili-arc-second scales, one often observes relativistic jet velocities, with flow Lorentz factors $\\gamma_u$ reaching the values above $10$ (cf. Ghisellini et al. 1996). At larger scales, velocity measurements are more difficult, but without entrainment of large amount of matter near the active galactic nuclear source the jet flow velocity must be also relativistic. A possible loading of a jet with matter is expected to be appended by a substantial amount of turbulence (Henriksen 1987) and related jet kinetic energy dissipation. However, in the FR~II radio sources, there are often observed jets efficiently transporting energy to the far-away hot spots and any jet breaking mechanism can not act too effectively near the central core. Also, the existing hydrodynamical simulations of relativistic jets show for possibility of extended stable jet structures (Marti et al. 1995, 1997; G\\'omez et al. 1995). Another argument suggesting the relativistic jet speed at all scales, may be based on the visible asymmetry of jets with respect to the nuclear source, if one believes the effect is caused by the high velocity of the essentially bi-symmetric outflow (cf. Bridle et al. 1994). Let us also note that the Meisenheimer et al. (1989) modelling of the shock acceleration process at extragalactic radio-source hot spots yields `the best-guess' jet velocities in the range ($0.1$, $0.6$)  The relativistic movement of the jet leads to shock wave formation in places where an obstacle or perturbation of the flow creates a sudden velocity jump. For jets loaded with a cold plasma the highly oblique conical shocks are formed within the jet tube. These shocks can have a non-relativistic character, involving the velocity jump perpendicular to the shock surface much smaller than the overall jet velocity $U \\sim c$. They lead to a limited kinetic energy dissipation and are usually claimed to be responsible for forming the so called `knots' along the jet. A much more powerful shock is formed at the final working surface of the jet. There, a substantial fraction of the jet energy is transferred into heating the jet's plasma, generating strong turbulence, boosting magnetic fields within the turbulent volume, and finally accelerating electrons and nuclei to cosmic ray energies. Rachen \\& Biermann (1993) considered the process of particle acceleration to ultra-high energies (UHE) at such shocks. They show that given the favourable conditions the UHE particles up to $\\sim 10^{20}$ eV can be formed. Then Rachen et al. (1993) show that assumption of UHE particle acceleration in extragalactic powerful radio sources is compatible with the current measurements of cosmic ray abundances and spectra at energies above $10^{17}$ eV. Additionally, the arrival directions of cosmic ray particles observed above $10$ EeV are correlated with the local galactic supercluster structure (Stanev et al. 1995; see, also, Medina Tanco et al. 1996, Sigl 1996, Sigl et al. 1995, 1996, Hayashida et al. 1996, Elbert \\& Sommers 1995, Geddes et al. 1996 and Medina Tanco 1998). An alternative model involving the several-Mpc-scale non-relativistic shocks in galaxy clusters is proposed by Kang et al. (1996; see also Kang et al. 1997).  As noted by us (Ostrowski 1990; henceforth Paper~I) a tangential discontinuity of the velocity field can also provide an efficient cosmic ray acceleration site if the considered velocity difference $U$ is relativistic and the sufficient amount of turbulence on both its' sides is present. The problem was extensively discussed in the early eighties by Berezhko with collaborators (see the review by Berezhko 1990) and in the diffusive limit by Earl et al. (1988) and Jokipii et al. (1989). In the present paper we consider the process of ultra high energy cosmic ray acceleration in relativistic jets including the possibility of such boundary layer acceleration. As the considerations of Rachen \\& Biermann (1993) treat the acceleration process at relativistic shock in a somewhat simplified way (see, also Sigl et al. 1995), in the first part of the next section (section 2.1) we review this process in some detail in order to understand the inter-relations between the conditions existing near the shock, the accelerated particle spectrum and the particle's upper energy limit. Then, in section (2.2), we present a short description of the basic physical model for the considered acceleration process acting at the jet boundary. We show (section 2.3) that in the conditions characteristic for relativistic jets in extragalactic radio sources, particles with energies above $10^{20}$ eV can be produced in this process without extreme parameter fitting. The required efficiency is discussed in section (2.4). We confirm the estimates presented previously for the shock acceleration, showing that the UHE particle flux observed at the Earth can be reproduced as a result of acceleration processes in jets of nearby powerful radio sources. In section 3 we discuss the problem of the particles' spectrum. With the use of Monte Carlo simulations, we consider the action of both processes acting near the terminal shock in a relativistic jet. Modification of the spectrum due to varying boundary conditions and jet velocity is discussed for the case of ($e^-$, $p$) jets expected to occur in the powerful FRII radio sources (cf. Celotti \\& Fabian 1993). The derived particle's upper energy limits are above the shock acceleration estimates and the spectrum modification at highest energies can resemble the observed above 10 EeV `ankle' structure. A short summary and final remarks are presented in section 4. A preliminary report about this work was presented in Ostrowski (1993b, 1996).  For the discussion that follows, we consider the jet propagating with the relativistic velocity, $U \\sim c$. We use $c$ = 1 units. ",
        "conclusions": ""
    },
    "9803/astro-ph9803066_arXiv.txt": {
        "abstract": "We present new Keck observations of giant arcs in the cluster Abell 2390.  High resolution two-dimensional spectra of two arcs show metal lines at $z=4.040\\pm 0.005$ with Ly$\\alpha$ emission spatially separated from the stellar continuum.  In addition, spectroscopy along the notorious `straight' arc reveals two unrelated galaxies, at z=0.913 and z=1.033, with absorption lines of MgII and FeII at z=0.913 seen against the more distant object, indicating the presence of a large gaseous halo. ",
        "introduction": "New high resolution spectra are presented for the high-redshift $z=4.04$ lensed system behind A2390. Figure 1 shows a section of the spectrum aligned along the long axes of the two main arcs.  This clearly shows that the Ly$\\alpha$ emission is spatially-separated from the continuum light and redshifted with respect to the interstellar lines, indicating an outward flow of enriched gas thought to be typical of starburst galaxies (Lequeux et al. 1995). All stellar and interstellar features are common to both spectra, confirming these arcs are images of one single highly-magnified galaxy. Note the Ly$\\alpha$ absorption is seen only in the southern portions of both arcs, coincident with the stellar continuum and with no associated Ly$\\alpha$ emission, indicating absorption of this line where the HI column is high. A lens model for this system is discussed in Frye \\& Broadhurst (1998). Keck infrared observations were also taken to study the stellar populations of this galaxy, by comparing flux ratios on opposite sides of the 4000 \\AA \\ break (Bunker, et al. these proceedings). ",
        "conclusions": ""
    },
    "9803/astro-ph9803316_arXiv.txt": {
        "abstract": "We have carried out K-band speckle observations of a sample of 114 X-ray selected weak-line T\\,Tauri stars in the nearby Scorpius-Centaurus OB association. We find that for binary T\\,Tauri stars closely associated to the early type stars in Upper Scorpius, the youngest subgroup of the OB association, the peak in the distribution of binary separations is at 90 A.U. For binary T\\,Tauri stars located in the direction of an older subgroup, but not closely associated to early type stars, the peak in the distribution is at 215 A.U. A Kolmogorov-Smirnov test indicates that the two binary populations do not result from the same distribution at a significance level of 98\\%.  Apparently, the same physical conditions which facilitate the formation of massive stars also facilitate the formation of closer binaries among low-mass stars, whereas physical conditions unfavorable for the formation of massive stars lead to the formation of wider binaries among low-mass stars. The outcome of the binary formation process might be related to the internal turbulence and the angular momentum of molecular cloud cores, magnetic field, the initial temperature within a cloud, or - most likely - a combination of all of these.  We conclude that the distribution of binary separations is not a universal quantity, and that the broad distribution of binary separations observed among main-sequence stars can be explained by a superposition of more peaked binary distributions resulting from various star forming environments. The overall binary frequency among pre-main-sequence stars in individual star forming regions is not necessarily higher than among main-sequence stars. ",
        "introduction": "Taurus-Auriga is the star forming region which has been most thoroughly surveyed for pre-main-sequence binary and multiple systems (see Mathieu 1994 for a review). For separations from 15 A.U.\\ to 2000 A.U., the binary frequency among T\\,Tauri stars in Taurus is 1.9 times as high as among nearby main-sequence stars (K\\\"ohler \\& Leinert 1998). Extrapolating over the whole range of separations yields a binary frequency of 100\\%, i.\\,e., each T\\,Tauri star in Taurus should be member of a binary or multiple systems. This apparent overabundance of binaries among pre-main-sequence stars is puzzling. One possible explanation is a decrease in the binary frequency as a function of the age of a stellar population (Patience et al.\\ 1998). However, a T\\,association like Taurus might not be the typical birthplace for low-mass stars, as up to 80\\% of all low-mass stars could originate in OB associations (Miller \\& Scalo 1978; see also Zinnecker et al.\\ 1992).  Scorpius-Centaurus is the most nearby OB association at a distance of about 145 parsec (de Zeeuw et al.\\ 1998). It consists of three subgroups (cf.\\ Figure 1) with ages ranging from 5 to 13 Myr (de Geus et al.\\ 1989). Upper Centaurus-Lupus (UCL) is the oldest subgroup of the association. Star formation started here 13 Myr ago and subsequently progressed throughout the parental giant molecular cloud (e.g.\\ Blaauw 1991 and references therein).  Based on observations with the EINSTEIN X-ray satellite, Walter et al.\\ (1994) identified 28 weak-line T\\,Tauri stars (WTTS) in Upper Scorpius (US), the youngest subgroup. 10 of these have been surveyed by Ghez et al.\\ (1993) for binary or multiple systems, and three binaries have been detected. The EINSTEIN fields covered only a small fraction of US (Fig.\\ 2) and the 28 WTTS did not allow for a statistical meaningful study of binary frequencies and separations. A search for visual binary stars among 74 ROSAT selected WTTS and post-T\\,Tauri stars in US (Kunkel et al., in prep) was carried by Brandner et al.\\ (1996). This survey was sensitive to binary separations down to 0\\farcs8, and revealed a rather high binary frequency in the region located between US and UCL (`US-B') and an apparent absence of wide visual binary stars in the center of US (`US-A').  In order to identify closer binary systems and to get a better statistics on possible spatial variations of binary star properties, we have carried out a speckle survey of 114 WTTS in US based on the lists by Walter et al.\\ (1994) and Kunkel et al.\\ (in prep).  \\begin{figure*}[htb] \\centerline{\\plotfiddle{fig1.ps}{11cm}{0}{90}{90}{-520}{-310}} \\figcaption{The spatial distribution of 532 proper motion members of the Scorpius-Centaurus  association based on HIPPARCOS measurements (adapted from de Zeeuw et al.\\ 1998). The boundaries between the subgroups Upper Scorpius (US), Upper Centaurus Lupus (UCL), and Lower Centaurus Crux (LCC) are indicated by dotted lines.  Star formation has progressed from the oldest subgroup UCL towards the younger subgroups US and LCC. The two adjacent fields of our survey for binary T\\,Tauri stars, centered on US (US-A, solid lines) and at the interface between US and UCL (US-B, dashed lines) are outlined. \\label{fig1}} \\end{figure*} ",
        "conclusions": "We have shown that the distributions of binary separations among weak-line T\\,Tauri stars in two adjacent fields (`US-A' and `US-B') in the Scorpius-Centaurus OB association are clearly distinct from each other and considerably more peaked than the (broad) distribution of binary separations observed among main-sequence field stars. In US-A, the WTTS are closely associated with B type stars, whereas in US-B only a few early type stars are present. We conclude that the same physical conditions which facilitate the formation of massive stars also facilitate the formation of closer binaries among low-mass stars, whereas physical conditions unfavorable for the formation of massive stars lead to the formation of wider binaries among low-mass stars.  The outcome of the binary formation process might be determined by the critical density at which the magnetic field support breaks down, the internal turbulence and the angular momentum of molecular cloud cores, the initial temperature within a cloud, or - most likely - a combination of all of these.  We further conclude that the distribution of binary separations is not a universal quantity. Instead, both the peak and the width of the distribution might vary from one star forming region to the next. The broad distribution of binary separations observed among main-sequence field stars can be understood as a superposition of binary populations originating in various star forming environments with very distinct peaks in the distribution of binary separations.  The apparent overabundance of binaries among T\\,Tauri stars in the Taurus-Auriga T\\,association might be explained by the fact that the distribution of binary separations there is strongly peaked towards $\\approx$ 30\\,A.U. Extrapolating from this very pronounced peak over the whole range of possible binary separations then leads to an erroneously high estimate of the overall binary frequency."
    },
    "9803/astro-ph9803120_arXiv.txt": {
        "abstract": "We investigate the non-gaussian properties of cosmic-string-seeded linear density perturbations with cold and hot dark matter backgrounds, using high-resolution numerical simulations. We compute the one-point probability density function of the resulting density field, its skewness, kurtosis, and genus curves for different smoothing scales. A semi-analytic model is then invoked to provide a physical interpretation of our results. We conclude that on scales smaller than $1.5{(\\Omega h^2)}^{-1}$Mpc, perturbations seeded by cosmic strings are very non-gaussian. These scales may still be in a linear or mildly non-linear regime in an open or $\\Lambda$-universe with $\\Gamma=\\Omega h \\lsim 0.2$. ",
        "introduction": "\\label{intro} At present,  the two main candidates for the origin of cosmic structure are inflation and topological defects (for a review, see Vilenkin \\& Shellard 1994). Although both scenarios may produce a power spectrum of density perturbations consistent with observations, they have very different predictions regarding the statistical properties of the density field. While most inflationary models produce gaussian random-phase initial conditions, defect models produce non-gaussian perturbations particularly on small scales. New results from cosmic-string-seeded structure formation using high-resolution simulations (\\markcite{ASWAs,ASWAl}Avelino {\\it et al.}~1997, 1998) were encouraging for models with $\\Gamma = \\Omega h = 0.1$--$0.2$ (see also Battye {\\it et al.}~1997); both the mass fluctuation amplitude at $8 h^{-1}$Mpc, $\\sigma_8$, and the power spectrum shape of cosmic-string-induced cold dark matter fluctuations, ${\\cal P}(k)$, were consistent within uncertainties with observational data (Peacock \\& Dodds 1994; Viana \\& Liddle 1996). However, because cosmic strings induce non-gaussian density perturbations on small scales, the power spectrum alone is insufficient to describe all the statistical properties of such a density field. This is even more important in open or $\\Lambda$-models because in those models the characteristic scales of the density field are shifted to larger scales relative to a flat model with $\\Lambda=0$.  In this Letter we investigate the non-gaussian properties of the linear density field induced by cosmic strings using higher-order statistics such as the skewness and the kurtosis of a one-point probability density function (PDF), as well as genus statistics. The non-gaussian properties we reveal provide a significant observational signature for cosmic string-seeded structure formation models on small length-scales. We note that previous analytic work has investigated the string-induced velocity field on scales above several $h^{-1}$Mpc, which was inferred to be gaussian (Vachaspati 1992; Moessner 1995), and that some of the features we study here were also observed in global topological defect models, notably for textures (Park, Spergel, \\& Turok 1991). Past work on genus statistics in the context of topological defects was made using toy models which incorporated some important features of the models in question (Brandenberger, Kaplan \\& Ramsey 1993; Albrecht \\& Robinson 1995; Avelino 1997).   Our first step in the present analysis was to perform high-resolution numerical simu\\-lations of cosmic string networks in an expanding universe (Allen \\& Shellard 1990) from which we subsequently computed the causally-sourced density field with either a cold or hot dark matter (CDM or HDM) background. The cosmic string simulations had a dynamic range extending from before the radiation-matter transition at $0.4 \\eta_{\\rm eq}$ through to deep into the matter era $8.4 \\eta_{\\rm eq}$, where $\\eta_{\\rm eq}$ is the conformal time at radiation-matter density equality. The structure formation simulation boxes contained $256^3$ grid-points and their physical volume was in the range $(4$--$100 h^{-1}$Mpc$)^3$. A much more detailed description of these methods is given by Avelino {\\it et al.}~(1997, 1998). ",
        "conclusions": "We conclude that on length scales smaller than\\ $1.5 {(\\Omega h^2)}^{-1}{\\rm Mpc}$ perturbations seeded by cos\\-mic strings are very non-gaussian, especially in the context of a CDM model. In an open or $\\Lambda$-universe with $\\Gamma=\\Omega h \\sim 0.15$, this scale will be shifted to $10 h^{-1}$Mpc, which may still be in the linear or mildly non-linear regime, thus potentially providing a strong empirical test for cosmic string models. It has been suggested that such non-gaussianity may imply that it is difficult in string models to deduce the parameter $\\beta=\\Omega_0^{0.6}/b$ from observations of density and velocity fields (van de Bruck, 1997). However, our results indicate otherwise on large scales because we find that string perturbations are very similar to gaussian-random phase fluctuations, when smoothed on sufficiently large scales."
    },
    "9803/astro-ph9803193_arXiv.txt": {
        "abstract": "Radio observations establish the B/A magnification ratio of gravitational lens 0957+561 at about 0.75.  Yet, for more than 15 years, the {\\it optical} magnfication ratio has been between 0.9 and 1.12.  The accepted explanation is microlensing of the optical source. However, this explanation is mildly discordant with (i) the relative constancy of the optical ratio, and (ii) recent data indicating possible non-achromaticity in the ratio.  To study these issues, we develop a statistical formalism for separately measuring, in a unified manner, the magnification ratio of the {\\it fluctuating} and {\\it constant} parts of the light curve.  Applying the formalism to the published data of Kundi\\'c et al. (1997), we find that the magnification ratios of fluctuating parts in both the g and r colors agrees with the magnification ratio of the constant part in g-band, and tends to disagree with the r-band value.  One explanation could be about $0.1$ mag of consistently unsubtracted r light from the lensing galaxy G1, which seems unlikely. Another could be that 0957+561 is approaching a caustic in the microlensing pattern. ",
        "introduction": "\\label{sect:intro}  The gravitational lens system 0957+561 has by now been observed at optical and radio wavelengths for nearly twenty years (Walsh, Carswell, and Weymann 1979; Porcas et al. 1979).  Radio studies have definitively established that the B/A magnification ratio of the lens, measured at the core of the radio images (which lies at the location of the optical point source images), is close to 0.75; some recent measurements are $0.75\\pm 0.02$ (Garrett et al. 1994), $0.752\\pm 0.028$ (Conner et al., 1992).  Note that while there is some controversy about the radio magnification ratio at the location of the radio jet, as opposed to the core (see, e.g., Garrett et al. 1994), only the core value interests us here.  Because the variability of the quasar in the optical is larger than in the radio, measurement of the B/A magnification ratio in the optical requires that the light curves be shifted by the correct time delay $\\tau$ before the ratio is taken.  Thus, the earliest determinations of B/A were incorrect.  For example, Young et al. (1980) obtained a ratio of 0.76, comfortingly -- yet erroneously -- close to the radio magnification.  However, at least from Vanderriest et al. (1989) on, who used a value for $\\tau$ quite close to definitive recent determinations (Kundi\\'c et al. 1997), it has been clear that the B/A magnification ratio of the optical continuum in the B and A point sources is quite different from 0.75, and moreover has remained at least fairly constant for the full history of observation.  Smoothing over observing seasons, Vanderriest et al. (1989) obtained a B/A ratio varying between about 0.9 and 1.05 over the observing seasons early-1980 through early-1986 (times referenced to A component), with a single best-fit value of 0.97. It is debatable whether the variation around the best-fit value is actual time variation of the lens magnification ratio (as distinct from time variation in the quasar luminosity, n.b.) or observational artifact. However, it does seem quite likely that the magnification ratio varied by no more than about $\\pm 8$\\% during this time.  More recently, the value recently obtained by Kundi\\'c et al. (1995, 1997) for the 1995 season (A component) is $1.12\\pm 0.01$ in g-band (with the error bar, a 95\\% confidence limit, depending somewhat on the method of reduction used).  So, it is quite plausible (and not contradicted by other measurements in the literature) that the optical B/A remained in the range 0.9 to 1.12 from 1980 through 1995, and possible that the variation has been considerably smaller than this range.  The discrepancy between the optical and radio magnification ratios has long been understood as due to microlensing (as predicted by Chang and Refsdal 1979, and Gott 1981).  The proper radius of the Einstein ring from a 0.5 $M_\\odot$ star at the lens galaxy redshift $z=0.36$, illuminated by the quasar at redshift $z=1.41$, is about $2\\times 10^{16}\\,h^{-1/2}$cm.  Since the radio emission region is much larger than this scale, it averages spatially over the microlensing pattern and is magnified by the macrolens ratio of 0.75.  If the optical magnification indeed differs by $\\sim 30$\\% from the macrolens value, then the optically emitting region must be smaller, or at most a few times larger, than the Einstein ring scale.  This accords nicely with (e.g.)  the size of an accretion disk smaller than 100 Schwarzschild radii around a $10^9$M$_\\odot$ black hole (a scale of $3\\times 10^{16}$cm).  This Einstein ring radius is only marginally, however, in accord with the apparent constancy of the microlensed magnification ratio:  Since the Earth, the microlensing star (or stars, the effect being collective), and the quasar each have (3-dimensional) peculiar velocities of at least 300 km/s, the Earth should move through $\\sim 100$\\% microlensing variations in $\\sim 10$ yr (see Kochanek, Kolatt, and Bartelmann, 1996 for related calculations). So, the observed microlensing is about a factor 10 {\\it too constant}, and one is invited to speculate on whether something other than luck is the reason.  Another invitation to speculation is the fact that Kundi\\'c et al. obtain rather different magnification ratios in their r- and g-band data, with the r ratio being $1.22\\pm 0.02$, with, again, the error bar depending on the method of analysis used.  By any interpretation of the error bars, however, the r and g results are strongly discrepant.  (Again note that there is no assumption that the fluctuations themselves have the same amplitude in the two colors, but only that the magnification {\\it ratio} should be the same.) Either a full $0.1$ mag of r-band galaxy light has escaped Kundi\\'c et al.'s careful subtraction in the B image, or something else is going on in the lens magnification ratio.  With these two hovering peculiarities (possible excess time-constancy, and possible non-achromaticity), it seems useful to try to get additional information on the magnification ratio.  This paper therefore asks the questions: Is the optical magnification ratio the same for the source region that produces the {\\it fluctuations} in quasar light as it is for the source region that produces the {\\it constant} light?  And, does the magnification ratio of the fluctuations (which we may call the ``AC'' magnification ratio or ``ACMR'') agree more closely with the r- or g-band magnification ratio previously measured (here called the ``DC'' magnification ratio or ``DCMR'')?  The answers to these questions can help diagnose the following situations: (i) If, as is true in many models, the size of the emitting region is much smaller than the microlensing scale, then all the magnification ratios should have the same value. (ii) If there is a problem with r-band galaxy subtraction -- or any other constant source of flux added to one lens component and not the other -- then the ACMR should represent the ``true'' microlens magnification ratio, and we might further expect it to be close to the g-band DCMR (where galaxy subtraction is a much smaller effect).  (iii) If the optically emitting quasar accretion disk has a scale comparable to the microlensing scale, and has (as seems almost inevitable) color gradients, then the r and g ACMRs, and r and g DCMRs, might all be distinct.  Indeed, the two ACMRs and two DCMRs then provide four distinct windows on the convolution of the accretion disk source with the microlensing pattern.  Conceptually, one measures a DCMR and an ACMR as follows: Shift one of the light curves (A,B) in time by $\\tau$ to undo the lens delay.  Fit each light curve by a constant value plus a residual time-varying part.  The ratio of the constant values is the DCMR. Now, for the two time-varying residuals, fit for a model that makes the B residual a constant times the A residual.  The best-fitting constant is the ACMR.  This conceptual formulation, while simple, is actually not quite right.  In the next section, we will give a statistical formulation of the problem that is more complete, and also more directly applicable to unevenly sampled data.  In Section 3, we discuss some implementation details, and in Section 4 we apply the formulation to the published data of Kundi\\'c et al. (1995, 1997).  Section 5 is discussion and conclusions. ",
        "conclusions": "While these data, in this analysis, do not support any very definitive conclusions, we may make the following remarks:  Occam's razor would seem to indicate that the r-band light curve of Kundi\\'c et al. has about $0.1$ mag of residual, unsubtracted, constant light, as perhaps from unmodeled small-scale variations in the lens galaxy surface brightness.  If this is the case, then all the data are compatible with a single magnification ratio for both colors and for both the fluctuating and constant pieces.  This in turn suggests an accretion disk scale much smaller than the microlensing scale, in accord with theoretical prejudice.  The utility of the ACMRs is that, taken together, they strongly favor the hypothesis that the g-band magnification ratio is the correct one, and that nothing more exotic is going wrong.  We note, however (per E. Turner, private communication), that the galaxy G1 is something like 2 magnitudes fainter than component B in r band; thus the amount of unsubtracted light would need to be comparable to the total brightness of G1, which seems quite unlikely.  It is up to the observers, not us, to decide whether Occam's razor should rule in this case.  If $0.1$ mag of residual is not possibly present, then we must conclude that the accretion disk scale is comparable to the microlensing scale, and that the constant r-band part of the disk is {\\it more} strongly magnified than (at least some of) the other three regions.  It seems likely on physical grounds that the fluctuating regions should be smaller than the constant regions, and that the g-band regions should be smaller than the r-band regions (temperature decreasing outward in the disk).  For the larger (r-band and constant) region to have a higher magnification ratio than an enclosed smaller region, the larger region must extend to a place where the magnification is a superlinear function of position in the sky.  This might indicate at least a fair chance of the B image passing through a caustic in the near ($\\sim 10$ year) future.  This possibility, as well as the reconciliation of the relative constancy of the magnification ratio over the last 15 years, will be explored by Monte Carlo simulations in another paper (Press and Kochanek, in preparation)."
    },
    "9803/astro-ph9803008_arXiv.txt": {
        "abstract": "We obtain restrictions on the universal baryon fraction, $f_{\\rm B} \\equiv \\Omega_{\\rm B}/\\Omega_0$, by assuming that the observed microlensing events towards the Large Magellanic Cloud are due to {\\em baryonic} MACHOs in the halo of the Galaxy and by extracting a bound to the total mass of the Milky Way from the motion of tracer galaxies in the Local Group.  We find a lower bound $f_{\\rm B} > 0.29^{+0.18}_{-0.15}$.  Consistency with the predictions of primordial nucleosynthesis leads to the further constraint on the total mass density, $\\Omega_0 \\alt  0.2$. ",
        "introduction": "It is a Herculean task to inventory the contents of the Universe (e.g., \\cite{fhp}).  A more modest goal might be to pin down the baryonic fraction of the total mass, $f_{\\rm B}$ (e.g., \\cite{wnef}, \\cite{xrc}).  If objects can be identified which are likely to provide a ``fair\" sample of $f_{\\rm B}$, we may avoid the daunting prospect of having to identify all the guises baryons may assume.  Large clusters of galaxies offer a very promising site (\\cite{wnef}; \\cite{xrc}; \\cite{xray1}; \\cite{xray2}).  To test the estimates of the systematic errors in $f_{\\rm B}$ derived from X-ray cluster data, it would be of value to measure $f_{\\rm B}$ in a completely different system, provided a case could be made that it will provide a ``fair\" sample. Suppose, for example, we could estimate the baryonic mass associated with the Galaxy.  If we could also measure the corresponding ``dynamical\" mass, we could obtain an independent estimate of $f_{\\rm B}$ whose systematic uncertainties (and dependence on the Hubble parameter) differ from those which accompany the X-ray cluster determinations. In this paper we focus on the Local Group of galaxies (LG), using the MACHO mass estimates (\\cite{MACHO}) for a lower bound on the baryonic mass and relying on LG dynamics to constrain the total mass estimate.  Microlensing experiments (\\cite{MACHO}) suggest that roughly half the mass in the halo of our Galaxy, out to the distance of the Large Magellanic Cloud (LMC), may be in the form of Massive Compact Halo Objects (MACHOs). One can imagine several exotic possibilities for the nature of the MACHOs. They could be very dense clusters of non-baryonic dark matter with special properties that allow them to clump inside their Einstein ring radii (\\cite{kt94}), or they could be primordial black holes.  Neither of these possibilities is especially well motivated and each has its intrinsic difficulties, but neither can be excluded a priori.  Stellar remnants such as old white dwarfs\\footnote{Neutron stars and black holes of stellar origin cannot constitute a significant halo fraction in view of the constraints arising from the observed metallicity and helium abundances (\\cite{ros90}).} appear to offer a more natural candidate (\\cite{MACHO}) which, however, is not without its problems too [e.g., white dwarfs require a rather narrow initial mass function in order to avoid overproducing low-mass stars or supernovae (\\cite{AL96})]. Dense and cold baryonic gas clouds have also been considered as a viable alternative for the observed gravitational microlenses (\\cite{cbc}; \\cite{cbc2}).  Finally, it must be kept in mind that the observed microlensing may be due to objects which are not in the halo of the Galaxy.  If the MACHOs are, indeed, stellar remnants (or cold baryonic gas clouds) in the halo of the Galaxy, then the mass of baryons within 50 kpc of the Galactic center is $M_{\\rm B} (50~{\\rm kpc}) \\geq M_{\\rm MACHO} = 2.0^{+1.2}_{-0.7} \\times 10^{11}M_\\odot$ (\\cite{MACHO}).  The purpose of the present paper is to extract information on the universal baryon fraction from this number assuming the MACHOs are revealing baryonic matter in the Galaxy halo, and from the dynamics of the Local Group of galaxies.  The constraint we obtain may be compared to the one derived from X-ray galaxy clusters (see, e.g., \\cite{xrc}; \\cite{xray1}; \\cite{xray2}), but it relies on different observations in a completely different physical system on a vastly different scale and, interestingly, has a different dependence on the Hubble parameter ($H_0 \\equiv 100h $~km~s$^{-1} $~Mpc$^{-1}$).  The value of $M_{\\rm B} (50~{\\rm kpc})$ derived from microlensing experiments is approximately 50\\% of the total mass of the Galaxy out to this distance.  The latter mass, presumably the sum of baryons and cold dark matter, is derived dynamically (see, e.g., \\cite{koch}). However, on the basis of this we cannot conclude that the primordial baryon fraction is $f_{\\rm B} \\approx 0.5$.  Baryons are ``strongly'' interacting particles, while for the (non-baryonic) cold dark matter all interactions except gravitational can be neglected.  Consequently, the density profile of the baryonic matter does not necessarily follow the density profile of the cold dark matter, and baryonic matter may be more (or less) concentrated towards the center of the gravitational well.  However, we may be able to estimate the primordial baryon fraction if we take the ratio of baryons (as revealed by the MACHOs) to the total mass on some larger scale, which should be sufficiently large so that the matter inflow or outflow across the boundary of the region is negligible.  The total mass of matter residing in such a larger region can be found dynamically; however, we cannot measure the mass of baryons separately on such larger scales.  Although the baryonic halo may be expected to extend outside of the 50 kpc scale (in the form, e.g., of MACHOs, diffuse gas, satellite Galaxies, etc.), by neglecting these extended baryons we can obtain a lower bound on $f_{\\rm B}$.  Indeed, while in the past there might have been violent processes of baryon ejection from the Galaxy accompanying, e.g., supernova explosions, analogous ejecta of cold dark matter is not expected.  Therefore, by neglecting the unknown ejected component of baryons we will be on the ``safe side'' in our inequality for $f_{\\rm B}$, which, we emphasize, does rely on our assumption that MACHOs are baryonic matter in the halo of the Galaxy.  For the larger reference scale we can choose the current turnaround radius for the  LG.  Initially, every shell of the Galaxy's building material expands with the Universe.  Gradually, this expansion slows down and eventually a gravitationally bound shell separates from the general expansion.  This shell stops expanding and then collapses (\\cite{gg72}).  The radius of this first stopping point is the turnaround radius.  With the passage of time shells that are more and more distant and less and less bound turn around sequentially, i.e., the turnaround radius propagates outward with time (for details see, e.g.,  \\cite{si}; \\cite{si2}; \\cite{stw}, 1997).  There is one shell that is turning around now, at present; the corresponding distance of this shell from the center of mass of the system is the current turnaround radius. Collisionless cold (non-baryonic) dark matter is restricted to remain within this radius, which is just what we want for the larger reference scale.  This picture of infall is valid independent of the assumption of spherical symmetry (the turnaround sphere will become a turnaround surface); for the model to be tractable analytically, we do assume spherical infall. ",
        "conclusions": "There remain several uncertainties in our LG baryon fraction estimate. One possibility which would weaken or even eliminate our constraint is if some of the observed microlensing events towards the LMC were due to an intervening satellite galaxy between us and the LMC, or due to debris in the LMC tidal tail (\\cite{z96}; \\cite{z97}).  However, the MACHO collaboration concluded (\\cite{fg}) that if the lenses were in a foreground galaxy, it must be a particularly dark galaxy; see also (\\cite{AG97}).  Moreover, the first observation of a microlensing event in the direction of the Small Magellanic Cloud (SMC) (\\cite{smc}), implies an optical depth in this direction roughly equal to that in the direction of the LMC.  This makes it unlikely that a dwarf galaxy or a stellar stream between us and the LMC is responsible simultaneously for the observed microlensing towards the LMC and the SMC (\\cite{fg}; \\cite{AG97}).  Recently, however, Gates et al. (1997) found Galactic models which explain the current microlensing data by a dark extension of the thick disk, reducing the MACHO fraction.  It is to be anticipated that as more microlensing data are accumulated, these uncertainties will be resolved.  We note that even in the absence of baryonic MACHOs there is still a limit, albeit much weaker, to $f_{\\rm B}$ from LG dynamics.  The mass of baryons in the disk of the Galaxy provides a lower bound to $M_{\\rm B}$ which is smaller by a factor of $\\sim 3$ than the microlensing estimate we have used (\\cite{fhp}).  Our lower bound to $f_{\\rm B}$ would be reduced by this factor while our upper bound to $\\Omega_0$ would be increased by the same factor.  In summary, if the observed microlensing events are the result of baryonic MACHOs in the Galaxy halo, then the dynamics of the LG may be used to infer a {\\it lower} bound to the universal baryonic mass fraction: $f_{\\rm B} > 0.29^{+0.18}_{-0.15}\\, t_{10}^2$.  If primordial nucleosynthesis is used to provide an {\\it upper} bound to the present baryonic density, we obtain an {\\it upper} bound to the present total mass density: $\\Omega_0  \\alt 0.2$ (with an extereme upper bound derived using nucleon-to-photon ratio based on the lithium abundance being $\\Omega_0  \\alt 0.47$)."
    },
    "9803/astro-ph9803187_arXiv.txt": {
        "abstract": "We write a non-relativistic Lagrangian for a hierarchical universe. The equations of motion are solved numerically and the evolution of the fractal dimension is obtained for different initial conditions. We show that our model is homogeneous at the time of the last scattering, but evolves into a self-similar universe with a remarkably constant fractal dimension. We also show that the Hubble law is implied by this model and make an estimate for the age of the universe.  \\vfill\\eject ",
        "introduction": "\\indent  This is a decisive time for cosmology since our theories for the large scale structure of the universe are being seriously challenged by the ever-growing amount of data.  The CfA1 redshift survey (de Lapparent, Geller \\& Huchra, 1986, 1988) was the first to reveal structures such as filaments and voids on scales where a random distribution of matter was expected. The most remarkable feature of these structures is the so-called ``great wall\" which is a coherent sheet of galaxies extended over an area of at least $60\\times 170 $ Mpc (Geller \\& Huchra 1989). Later on, deep pencil beam surveys (Broadhurst et al. 1990), the redshift surveys based on IRAS catalogue (Efstathiou et al. (1990a, 1990b), Saunders et al. 1991, Fisher et al 1996), the deep wide angle survey SSRS (de Costa 1988, 1994) and some others have shown inhomogeneities at scales where the galaxy-galaxy and cluster-cluster correlations were believed to be negligible. Of particular significance for the future are the two extensive redshift surveys, SLOAN and 2dF, about to commence, which aim to trace the 3-dimensional distribution of over one million galaxies across the northern and southern skies.  The first quantitative study of the cosmic inhomogeneity lead to the well-known $1.8$ power law behaviour of the galaxy-galaxy correlation function (Groth \\& Peebles 1977, Peebles 1980, Davis \\& Peebles (1983a, 1983b)). Although this law has been consistently identified in different catalogues, the break away from it at larger scales and a crossover to homogeneity has not yet been established (Davis 1996, Pietronero 1987, Coleman \\& Pietronero 1992, Pietronero 1996). Whether there is a crossover to homogeneity or not, the power law nature of the two-point and higher-order correlation functions is itself suggestive of some kind of scaling behaviour at least in some range. The simplest structure that obeys such a scaling law is a single fractal.  A theoretical model describing a non-analytic inhomogeneous scale-invariant universe is non-existent. The most-extensively-studied inhomogeneous cosmological model is Tolman spacetime (Tolman 1934, Bondi 1947). Tolman's dust solution has been used to model a hierarchical cosmology compatible with the observational analysis of the redshift surveys (Bonnor 1974, Ribeiro (1992a, 1992b), 1993). Recently, the Einstein equation for a scale-invariant spherically symmetric inhomogeneous, but isotropic, universe, which allows a non-vanishing pressure has been solved (Abdalla \\& Mohayaee 1997). However, the results obtained in this way are perturbative, assume a preferred center for the universe and violate the linearity of the Hubble law. A self-similar universe avoiding such difficulties can only be constructed for a non-analytic distribution of matter. It is rather a difficult task to construct a fractal metric and to solve Einstein equation for a self-similar universe. However, many cosmological phenomena can be accurately described by the Newtonian gravity, especially in the present matter-dominated era.  In this work, we construct a self-similar universe whose dynamics is governed by the Newtonian gravity. We divide the universe into $k$ spherical clusters each of which contains $k$ subclusters which in their turn contain $k$ sub-subclusters. This clustering cascades down all the way to the level of the galaxies which are at the lowest rung on the clustering ladder. The mass and radius of each cluster can be used to define the fractal dimension of our model.  We write the kinetic and potential energies of each cluster in terms of its center of mass energy and the internal energies of its subclusters. The thermal energy is obtained by requiring the total entropy of the canonical ensemble of the clusters and their subclusters to remain constant. The final Lagrangian is formulated in terms of two dynamical parameters: radius of the largest cluster and its ratio to the radius of its subclusters. The radius of the largest and smallest clusters, the ratio of the mass to the critical mass contained in a sphere of radius $20$ Mpc, the number of subclusters in each cluster and the ratio that characterises the relative significance of thermal and gravitational energies are left as free parameters. By fixing these to different observational values, we are able to solve the equations of motion numerically using a Pascal program. From the solutions, we can trace the evolution of the fractal dimension and verify the linearity of the velocity-distance relationship over time scales comparable to the age of the universe.  The results are remarkable. We observe that for different initial conditions a nearly homogeneous universe with a fractal dimension close to 3 evolves into a universe with a fractal dimension of the order of 2 at the present time. This fractal dimension fluctuates slightly about the value of 2 over the future times but remains on the average constant. We also show that for insignificant thermal energies, the Hubble law is closely obeyed by our model.  We also make an estimate for the age of our self-similar universe. This is one of the most challenging problems of Friedmann cosmology since the observed age of the old stars in the globular clusters is far bigger than the value estimated for the age of the universe in the nearly flat standard model. The age of the universe obtained in our model is related to the radius at which the crossover to homogeneity occurs. The farther the crossover radius the older is the universe.  This article is organized as follows. In Section 2, we formulate our clustering model. In Section 3, we obtain the kinetic, the potential and the thermal energies, write the Lagrangian and the equations of motion. In Section 4, we solve these equations numerically and discuss the validity and limitations of our Newtonian approximation. In Section 5, we show different plots of the fractal dimension for different choices of the initial condition. Hubble law is discussed in Section 6. In Section 7, we study the evolution of the scale factor and obtain a value for the age of the universe in our model. Section 8 is devoted to the conclusion. ",
        "conclusions": "\\indent  We have constructed a model for a nonrelativistic fractal universe. Our model starts off homogeneous and evolves rapidly to a self-similar universe with a remarkably constant fractal dimension of about 2. The homogeneity at the earlier times explains the isotropy of the microwave background radiation. We have also shown that the Hubble law is closely obeyed by our model for small thermal energies. We have estimated the age of the universe and have shown that it complies with the corresponding observational results. It remains an open problem to extend our model to multi-fractals and to the relativistic regime.  {\\bf Acknowledgements} We thank R. Mansouri and M. Khorrami for useful discussions. This work  has been partially supported by Conselho Nacional de Desenvolvimento Cient\\'\\i fico e Tecnol\\'ogico, CNPq, Brazil, and Funda\\c c\\~ao de Amparo \\`a Pesquisa do Estado de S\\~ao Paulo (FAPESP), S\\~ao Paulo, Brazil."
    },
    "9803/astro-ph9803134_arXiv.txt": {
        "abstract": "We have completely mapped the Galactic globular cluster NGC~1851 with large-field, ground-based $VI$ CCD photometry and pre-repair $HST$/WFPC1 data for the central region.  The photometric data set has allowed a $V $ vs. $ (V-I) $ colour--magnitude diagram for $\\sim$ 20500 stars to be constructed. From the apparent luminosity of the horizontal branch (HB) we derive a true distance modulus $(m-M)_0$ = 15.44 $\\pm$ 0.20.  An accurate inspection of the cluster's bright and blue objects confirms the presence of seven ``supra-HB'' stars, six of which are identified as evolved descendants from HB progenitors.  The HB morphology is found to be clearly bimodal, showing both a red clump and a blue tail, which are not compatible with standard evolutionary models. Synthetic Hertzsprung--Russell (HR) diagrams demonstrate that the problem could be solved by assuming a bimodal efficiency of the mass loss along the red giant branch (RGB). With the aid of Kolmogorov--Smirnov statistics we find evidence that the radial distribution of the blue HB stars is different from that of the red HB and subgiant branch (SGB) stars.  We give the first measurement of the mean  absolute $I$ magnitude for 22 known RR~Lyr variables ($<M_{I}({\\rm RR})> = 0.12 \\pm 0.20$ mag at a metallicity [Fe/H]~=~--1.28). The mean absolute $V$ magnitude is $<M_{V}(\\rm RR)> = 0.58 \\pm 0.20$ mag, and we confirm that these stars are brighter than those of  the zero-age HB (ZAHB).  Moreover, we found seven new RR~Lyr candidates (six $ab$ type and one $c$ type). With these additional variables the ratio of the two types is now $N_c$/$N_{ab} = 0.38$.   From a sample of 25 globular clusters a new calibration for $\\Delta V_{\\rm bump}^{\\rm HB}$ as a function of cluster metallicity is derived. NGC~1851 follows this general trend fairly well. From a comparison with the theoretical models, we also find some evidence for an age--metallicity relation among globular clusters.   We identify 13 blue straggler stars, which do not show any sign of variability. The blue stragglers are less concentrated than the subgiant branch stars with similar magnitudes for $r>80$ arcsec.   Finally, a radial dependence of the luminosity function, a sign of mass segregation, is found.  Transforming the luminosity function into a mass function (MF) and correcting for mass segregation by means of multi-mass King--Michie models, we find a global MF exponent $x_0=0.2\\pm 0.3$. ",
        "introduction": "Galactic globular clusters (GGC) are dynamically evolved objects. In order to understand the interplay between the internal dynamical processes and the influence of the Galactic potential, we must study a sample of GGCs comprising objects whose concentration, position in the Galaxy, luminosity and metallicity cover the whole observed range. The mass function and the radial profile must be determined for each cluster, in order to carry out a detailed dynamical analysis.   The introduction of large-size CCDs has made this kind of investigations possible. With these detectors it is also possible to obtain deep photometry for the nearest globulars, and therefore to probe their mass functions over large mass intervals, in order to reach those MS stars which are more sensitive to dynamical effects (e.g. Pryor et al. 1986).     A rich sample of stars is also essential in order to reveal and study the shortest-lived (and hence poorly known) phases of the stellar evolution (Renzini \\& Buzzoni 1986). Furthermore, the interactions between the single stars affect their evolution (e.g. Djorgovski et al. 1991).  To establish the reliability of the stellar evolutionary models, we must therefore ascertain to what extent a GC colour--magnitude diagram and luminosity function is changed by the interactions among its stars.     For the above reasons, our group started a project aimed at studying a number of globular clusters covering a wide range of the relevant parameters. NGC~1851 ($\\alpha_{2000} = 5^{\\rm h} 14^{\\rm m} 6^{\\rm s}.30$; $\\delta_{2000} = -40^\\circ 2\\arcmin 50\\farcs00$) has been selected for its position and its concentration. Its galactocentric distance, which is about twice that of the Sun, and its distance of 7.1~kpc from the Galactic plane (Djorgovski 1993) make it a typical halo object. Its concentration $c = 2.24$ is one of the highest in the list of Trager et al. (1995). A recent measurement of the cluster's proper motion has  confirmed that NGC~1851 has halo-type kinematics (Dinescu et al. 1996). According to these authors, the space velocities of the cluster are $U=256\\pm35$~km s$^{-1}$, $V=-195\\pm26$~km s$^{-1}$,  $W=-122\\pm30$~km s$^{-1}$, $\\Pi=195\\pm37$~km s$^{-1}$ and $\\Theta=167\\pm37$~km s$^{-1}$.      Past photometric studies of the cluster are given in Alcaino (1969, 1971, 1976), Stetson (1981, hereafter S81), Sagar et al. (1988), Alcaino et al. (1990) and Walker (1992a, hereafter W92). The most exhaustive analysis is that of W92. His main results are that: (1) the cluster core, although unresolved, appears to be blue; (2) the HB is bimodal, showing both a red clump and an extended blue tail; (3) there are no radial trends in the relative numbers of red horizontal branch (RHB), blue horizontal branch (BHB) and red giant branch (RGB) stars for 48\\arcsec~ $< r <$ 190\\arcsec~; (4) the RGB ``bump'' is at $V$ = 16.15 $\\pm$ 0.03 mag; (5) the population ratio $R = N({\\rm HB})/N({\\rm RGB})$ has a value 1.26 $\\pm$ 0.10, which corresponds to a helium abundance $Y$ = 0.23 $\\pm$ 0.01 (computed by means of the R-method; e.g. Renzini 1977); (6) there are six blue straggler (BS) stars and six supra-RHB stars $[$15.7 mag $<$ $V$ $<$ 16.0 mag; 0.6 $< (B-V) <$ 0.8$]$ in the region 120\\arcsec~ $< r <$ 220\\arcsec~, and there is evidence of segregation only for the BS stars, so an origin for supra-RHB stars from BS stars is not supported by W92 data; (7) no significant abundance spread is found from the colour width of the main sequence (MS); and finally (8) an age of 14 $\\pm$ 1 Gyr results from the $\\Delta$~($B$$-$$V$) method (Sarajedini \\& Demarque 1990; VandenBerg et al. 1990).    We have now obtained new {\\sl large-field} CCD $V$, $I$ photometry for NGC~1851. The new data set makes it possible to re-analyse the stellar content of the cluster with a much richer sample and, for the first time, allows a comprehensive study of its dynamical properties.  However, the central regions of the cluster cannot be studied with this ground-based material, due to the extreme crowding of the core. To overcome this limitation, pre-repair {\\it Hubble Space Telescope} ({\\it HST}) images have been retrieved from the archives and reduced in order to sample the central stellar content, in particular the radial distribution of the HB stars. For the sake of comparison, the photometric catalogue of W92 has been also used.                  \\begin{figure}[t] \\psfig{figure=n1851_neg_.ps,width=8.8cm} \\caption[]{ The observed NTT/EMMI fields, sketched over a POSS field.} \\label{field_map} \\end{figure}   \\begin{small} \\begin{table}[t] \\caption[]{Log of NTT/EMMI observations.} \\label{observs} \\begin{tabular}{cccccccc} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Nr.   &  Field & $t_{\\rm exp}$(s) &  Filter & Date         & FWHM $[\\arcsec]$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 1    &  6     & 50  &    $V$     &  1993 Feb 18  & 1.2 \\\\ 2    &  6     & 70  &    $I$     &  1993 Feb 18  & 1.2 \\\\ &&&&&\\\\ 3    &  5     & 45  &    $I$     &  1993 Feb 18  & 1.2 \\\\ 4    &  5     & 10  &    $I$     &  1993 Feb 18  & 1.4 \\\\ 5    &  5     & 10  &    $V$     &  1993 Feb 18  & 1.1 \\\\ 6    &  5     & 30  &    $V$     &  1993 Feb 18  & 1.3 \\\\ &&&&&\\\\ 7    &  2     & 30  &    $V$     &  1993 Feb 18  & 1.0 \\\\ 8    &  2     & 40  &    $I$     &  1993 Feb 18  & 1.2 \\\\ &&&&&\\\\ 9    &  3     & 60  &    $I$     &  1993 Feb 18  & 1.2 \\\\ 10    &  3     & 44  &    $V$     &  1993 Feb 18  & 1.1 \\\\ &&&&&\\\\ 11    &  1     & 45  &    $V$     &  1993 Feb 18  & 1.0 \\\\ 12    &  1     & 60  &    $I$     &  1993 Feb 18  & 1.2 \\\\ &&&&&\\\\ 13    &  4     & 60  &    $I$     &  1993 Feb 18  & 1.1 \\\\ 14    &  4     & 45  &    $V$     &  1993 Feb 18  & 1.1 \\\\ &&&&&\\\\ 15    &  7     & 45  &    $V $    &  1993 Feb 18  & 1.1 \\\\ 16    &  7     & 60  &    $I$     &  1993 Feb 18  & 1.2 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 17    &  8     & 70  &    $I$     &  1993 Feb 19  & 1.1 \\\\ 18    &  8     & 55  &    $V$     &  1993 Feb 19 & 1.1 \\\\ &&&&&\\\\ 19     &  9    & 55  &    $V$     &  1993 Feb 19 & 1.2 \\\\ 20     &  9    & 70  &    $I$     &  1993 Feb 19 & 1.2 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 21     &  back  & 120 &    $V$       & 1993 Dec 10  & 1.0 \\\\ 22     &  back  & 180 &    $I$       & 1993 Dec 10  & 1.1 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} \\end{small} ",
        "conclusions": "\\label{thediscussion}  We have presented new large-field CCD photometry for $\\sim20500$ stars in the Galactic halo globular cluster NGC~1851, from both groundbased observations and a pre-repair {\\it HST} field. The photometric catalogue has been used to build a $V$ vs. ($V$--$I$) CMD, which has been analysed in detail. An extensive comparison of our data set with the predictions of the stellar models has also been performed.  The effects of the dynamical evolution over the main sequence mass function have been investigated by means of a completeness-corrected luminosity function and the radial-count profile.   \\paragraph{The evolved stellar content}  With an accurate inspection of the cluster bright-blue objects, and a comparison with the numbers predicted from the background field and the Galactic count models, we have confirmed the presence of seven ``supra-HB'' stars in the CMD of NGC 1851. We have shown that six of the ``supra-HB'' stars could be evolved descendants from HB progenitors (post-HB or planetary nebulae).  We have shown that standard evolutionary models are not able to reproduce the observed bimodal distribution of stars along the HB.  Synthetic HR diagrams demonstrate that the problem could be solved by assuming that the efficiency of the RGB mass loss actually encompasses values going from 0.25 to 0.48.  We have found evidence that the radial distribution of the blue HB stars is different from that of red HB and SGB stars.  The BHB stars are significantly more concentrated than the SGB stars for $r>100$ arcsec. Though this distribution cannot be easily interpreted in terms of dynamical evolution, it might be related to the anomalous distribution of the BSs (see below).  All the 27 known variable stars have been identified, and 26 have been measured in both colours (the remaining one being saturated). Twenty-two of them are RR~Lyr variables. For the first time, our photometry has allowed the mean absolute $I$ magnitude of the RR~Lyr variables to be obtained at a metallicity [Fe/H]~=~--1.28. The RR~Lyr are brighter than the ZAHB in the $V$ band, in accordance with the relation given by Carney et al. (1992). The positions and the photometry for seven new RR~Lyr candidates have been given. With these additional variables the ratio of the two types is now $N_c$/$N_{ab} = 0.38$, which reduces the current estimate N$_c$/$N_{ab} = 0.47$ (Wehlau et al. 1982).  Thirteen BS stars have been identified outside the inner 80 arcsec. They do not show any sign of variability. We have investigated the radial distribution of the BSS. For $r>80$ arcsec, the BSs are less concentrated than the SGB stars with the same $V$ magnitude. We argue that the distribution of the BSs in the outer envelope of NGC 1851 might be similar to the distribution found by Ferraro et al. (1997) for the BSs in the envelope of M3.  We have considered a sample of 25 globular clusters and have derived a new calibration for the $\\Delta V_{\\rm bump}^{\\rm HB}$ parameter as a function of cluster metallicity, and we have found that NGC~1851 follows this general trend fairly well.  From a comparison with the corresponding slopes predicted by the isochrones library from Bertelli et al.  (1994), we have found that perhaps an age--metallicity relation actually exists among globular clusters, with the metal poorest possibly being older.     \\paragraph{Dynamical status of NGC~1851}  We have been able to derive a complete LF down to $V \\simeq 23.5$ mag for stars in the region 190\\arcsec~ $< r <$ 650\\arcsec~, and down to $V \\simeq 22$ mag in the region 120\\arcsec~ $< r <$ 189\\arcsec~. The external LF is steeper than the internal one, and we have interpreted this result as a sign of mass segregation. By using the most updated mass--luminosity relations we have obtained MFs which can be well fitted by power laws with distinct exponents $x$. The observed value for the external MF is $x = 1.52 \\pm 0.18$, which is steeper than the value $0.89 \\pm 0.20$ found for the internal one.   The global MF has been determined correcting the two observed mass functions for the effects of mass segregation, as predicted by the multi-mass King--Michie model which best fits the observed light profile of NGC 1851. The two values for the slope of the MF are compatible with the model if a global MF exponent $x_0=0.2\\pm 0.3$ is adopted. This value for the global MF slope is marginally smaller (MF flatter) than what would be expected from the relation between the slope of the MFs and the position in the Galaxy and the metallicity of the cluster proposed by Djorgovski et al. (1993). This might indicate that NGC 1851 has had a stronger gravitational interaction with the Galactic disc than the average of the Galactic GCs with similar position and metallicity.     \\paragraph{} The above results indicate that NGC~1851 is a cluster where the dynamical evolution has affected both its evolved and unevolved stellar content. While the single findings are not of high statistical significance (mostly due to the small size of the stellar samples), taken together they give a coherent picture. Stellar encounters have led to mass segregation, as shown by the MF, which is steeper and steeper going from external to internal regions. They have probably contributed to the creation of the observed group of blue straggler stars, and possibly have triggered the formation of a blue tail in the HB.     The internal dynamics of NGC~1851 has therefore influenced the evolution of its stars, introducing effects not reproducible by standard models.  In turn, the dynamical evolution induced by the external gravitational field of the Galaxy has also very probably contributed to the modification of the present-day stellar population of NGC 1851, as strongly suggested by the anomalously flat global mass function."
    },
    "9803/astro-ph9803302_arXiv.txt": {
        "abstract": "The shape of the radial velocity and light curves of 24 long-period ($30 \\leq P \\leq 134$ d) Cepheids in the Magellanic Clouds shows a progression with the period. The sequences of the radial velocity and light curves are based only on a small sample of stars; however, evident changes of the shape can be seen in Cepheids with period between 90 and 134 d. The Fourier parameter--period diagrams for the radial velocity curves show trends which remind in part those of Cepheids with period near 10 d. The plausible interpretation is a resonance, probably $P_0/P_1=2$ between the fundamental and the first overtone mode. The possible importance of this phenomenon for the study of stellar structure and evolution in relatively far galaxies is emphasized \\footnote {Tables 2 and 3 are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5)}. ",
        "introduction": "\\begin{table*} \\caption[]{List of the analyzed long-period Cepheids in Magellanic Clouds} \\begin{flushleft} \\begin{tabular}{lllllllllll} \\hline\\noalign{\\smallskip} HV & $P_L$ & $N_L$ & Ref$_L$  & ord$_L$ & $\\sigma_L$ (mag) & $P_{RV}$ & $N_{RV}$ & Ref$_{RV}$   & ord$_{RV}$ & $\\sigma_{RV}$ (km~s$^{-1}$)  \\\\ 904  L  & 30.400  &  58 & 10      & 6 & .034 &        &    &   &   &     \\\\ 1002 L  & 30.4694 &  60 & 1 - 6   & 7 & .032 &        &    &   &   &     \\\\ 899  L  &         &     &         &   &      & 31.027 & 49 & 7 & 7 & 1.7 \\\\ 2294 L  & 36.530  &  98 & 2 - 6   & 8 & .032 &        &    &   &   &     \\\\ 879  L  & 36.817  &  31 & 4       & 5 & .048 & 36.782 & 43 & 7 & 6 & 1.4 \\\\ 909  L  & 37.5828 &  61 & 1 - 4   & 4 & .047 & 37.510 & 64 & 7 & 5 & 1.1 \\\\ 11182 S & 39.1941 &  44 & 2,3,5   & 3 & .039 &        &    &   &   &     \\\\ 2257 L  &         &     &         &   &      & 39.294 & 48 & 7 & 6 & 1.7 \\\\ 2195 S  & 41.7988 &  39 & 1,2,5   & 5 & .049 &        &    &   &   &     \\\\ 2338 L  & 42.2153 &  55 & 2,4,5   & 7 & .048 & 42.184 & 59 & 7 & 8 & 1.4 \\\\ 837  S  & 42.739  &  55 & 2 - 5   & 6 & .037 & 42.673 & 86 & 9 & 10 & 1.6 \\\\ 877 L   & 45.107  &  78 & 2 - 5   & 3 & .041 &        &    &   &   &     \\\\ 900 L   & 47.5417 &  62 & 1 - 5   & 5 & .039 & 47.544 & 55 & 8 & 6 & 1.2 \\\\ 2369 L  & 48.3311 & 108 & 1 - 6   & 7 & .034 &        &    &   &   &     \\\\ 824 S   & 65.8306 &  62 & 1,4     & 5 & .030 & 65.755 & 40 & 8 & 5 & 1.7 \\\\ 11157 S &         &     &         &   &      & 69.06  & 62 & 9 & 5 & 1.2 \\\\ 834 S   & 73.399  & 106 & 2 - 6   & 7 & .035 & 73.648 & 53 & 8 & 7 & 1.3 \\\\ 2827 L  & 78.858  &  55 & 3,4     & 3 & .026 & 78.626 & 43 & 7 & 4 & 1.3 \\\\ 829 S   &         &     &         &   &      & 85.577 & 45 & 8 & 6 & 1.0 \\\\ 5497 L  & 99.078  &  67 & 3 - 6   & 3 & .026 & 99.040 & 41 & 8 & 3 & 1.2 \\\\ 2883 L  & 109.277 &  58 & 2 - 4,6 & 5 & .032 & 109.071 & 52 & 8 & 3 & 2.1 \\\\ 2447 L  & 117.941 &  68 & 3 - 6   & 3 & .022 & 118.841 & 39 & 8 & 2 & 1.3 \\\\ 821 S   & 127.490 & 117 & 1 - 6   & 4 & .040 & 127.083 & 50 & 8 & 4 & 1.4 \\\\ 883 L   & 133.893 &  98 & 1 - 6   & 4 & .045 & 134.75  & 92 & 9 & 6 & 2.7 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{flushleft} Ref.: 1. Gascoigne \\& Kron (1965); 2. Madore (1975); 3. Van Genderen (1977, 1983); 4. Martin \\& Warren (1979); 5. Eggen (1977); 6. Freedman et al. (1985); 7. Imbert et al. (1985); 8. Imbert et al. (1989); 9. Imbert (1994); 10. Sebo \\& Wood (1995). \\end{table*} The longer period Cepheids are particularly important in the context of the primary distance scale because they are bright enough to be visibile at a great distance. However, due to their low number, they have been poorly studied both observationally and theoretically. Simon \\& Kanbur (\\cite{sk}) considered 50 Cepheids in Galaxy and IC 4182 with period $P$ less than 70 d, compared them with hydrodynamical pulsation models and concluded that a detailed comparison between theory and observations must await a more extensive and accurate sample of observed stars. Antonello \\& Morelli (\\cite{am}) analyzed all the available photometric $V$ data of galactic Cepheids with period less than 70 d looking for possible resonance effects; they noted some small features in the Fourier parameters--period diagrams which were ascribed tentatively to expected resonances. Aikawa \\& Antonello (\\cite{aa}) tried to reproduce these observations with nonlinear models, but their conclusion was that the increasing nonadiabaticity of the pulsation with period probably reduces the effectiveness of resonance mechanisms. Finally, Simon \\& Young (\\cite{sy}) studied long period Cepheids in the period range $10 \\leq P \\leq 50$ d in Magellanic Clouds looking for galaxy-to-galaxy differences in the Cepheid distributions. Resonances between pulsation modes, which were studied essentially in Cepheids with $P$ less than about 30 d, represent a powerful comparison tool between observations and theoretical model predictions, because they affect the shape of the curves of pulsating stars in specific period ranges. The comparison of the Fourier parameters of observed and theoretical light and radial velocity curves allows to probe the stellar interior and to put constraints on the stellar physical parameters. After the work of Simon \\& Lee (\\cite{sl}) on the resonance $P_0/P_2=2$ at $P_0 \\sim 10$ d between the fundamental and second overtone mode in classical bump Cepheids, several papers by various authors were devoted to this topic, from both the observational and theoretical point of view. For example, Buchler \\& Kovacs (\\cite{bk}) and Moskalik \\& Buchler (\\cite{mb}) studied the general effects of 2:1 and 3:1 resonances in radial stellar pulsations and discussed the possible astrophysical implications, Petersen (\\cite{pe}) discussed the possible two- and three-mode resonances in Cepheids, and Antonello (\\cite{an}) looked for the expected effects in short period Cepheids. Recent reviews on galactic and Magellanic Cloud Cepheids pulsating in fundamental and first overtone mode and on the problems raised by the comparison with the pulsational models are those by Buchler (\\cite{bu}) and Beaulieu \\& Sasselov (\\cite{bs}). The resonance effects in Magellanic Cloud Cepheids cannot be reproduced by models constructed using current input physics and reasonable mass--luminosity relations; in particular the case of first overtone Cepheids characterized by $P_1/P_4=2$ (Antonello et al. \\cite{apr}) has proven to be rather difficult for theorists. We mention in passing also the recent resonance $P_2/P_6=2$, studied in the models of hypothetical second overtone mode Cepheids (Antonello \\& Kanbur \\cite{ak}).  In the present work we have considered the long-period Cepheids in Magellanic Clouds, with available photometric and radial velocity data which were suitable for Fourier decomposition. The initial purpose of the work was simply to extend the comparison between theory and observations to Cepheids with the longest known periods, but the probable discovery of a new resonance effect suggested to publish the present Letter in advance of the comparison with the hydrodynamical models (Antonello \\& Aikawa, in preparation). ",
        "conclusions": "The analysis of the radial velocity and light curves of the long-period Cepheids in the Magellanic Clouds indicates the presence of a progression which we interpret tentatively as an effect of the resonance $P_0/P_1=2$ between the fundamental and the first overtone mode. Since the longest period Cepheids are also the brightest, the present results could be of some importance for the study of stellar structure and evolution in far galaxies, because the stars with $P \\sim $ 100 d ($M_V \\sim -7$ mag) are about three magnitudes brighter than those at 10 d, in which the well known resonance $P_0/P_2=2$ is occurring. The disadvantage is the low number of such stars. Few Cepheids with $P > 80$ d have been found in relatively nearby galaxies (NGC 6822, IC 1613, NGC 300; see e.g. Madore \\cite{mm}), while in the Magellanic Clouds there are just a few of stars in comparison with a total of some thousand Cepheids. Presently the Hubble Space Telescope Key Project on the Extragalactic Distance Scale is optimized for the detection of Cepheids with period between 3 and 60 d (e.g. Ferrarese et al. \\cite{fe}), therefore it is not possible to derive reliable conclusions about the number of long-period stars. We just note that in NGC 925 Silbermann et al. (\\cite{sil}) found 4 stars with probable $P > 80$ d over a total of 80 Cepheids.  Assuming that a sufficient number of such stars is detected and our interpretation is correct, the comparison of the observed resonance effect with nonlinear model predictions will allow to test the input physics and put constraints on the physical parameters of the stars in relatively far galaxies in the same way as it is occurring for the Galaxy and Magellanic Cloud Cepheids with shorter periods."
    },
    "9803/astro-ph9803244_arXiv.txt": {
        "abstract": "The observed evolution of the galaxy cluster X-ray integral temperature distribution function between $z=0.05$ and $z=0.32$ is used in an attempt to constrain the value of the density parameter, $\\Omega_{0}$, for both open and spatially-flat universes.  We estimate the overall uncertainty in the determination of both the observed and the predicted galaxy cluster X-ray integral temperature distribution functions at $z=0.32$ by carrying out Monte Carlo simulations, where we take into careful consideration all the most important sources of possible error.  We include the effect of the formation epoch on the relation between virial mass and X-ray temperature, improving on the assumption that clusters form at the observed redshift which leads to an {\\em overestimate} of $\\Omega_0$.  We conclude that at present both the observational data and the theoretical modelling carry sufficiently large associated uncertainties to prevent an unambiguous determination of $\\Omega_{0}$.  We find that values of $\\Omega_{0}$ around 0.75 are most favoured, with $\\Omega_{0}<0.3$ excluded with at least 90 per cent confidence.  In particular, the $\\Omega_{0}=1$ hypothesis is found to be still viable as far as this dataset is concerned.  As a by-product, we also use the revised data on the abundance of galaxy clusters at $z=0.05$ to update the constraint on $\\sigma_8$ given by Viana \\& Liddle \\shortcite{VL}, finding slightly lower values than before. ",
        "introduction": "The number density of rich clusters of galaxies at the present epoch has been used to constrain the amplitude of mass density fluctuations on a scale of $8\\,h^{-1}\\,{\\rm Mpc}$ (Evrard 1989; Henry \\& Arnaud 1991; White, Efstathiou \\& Frenk 1993a; Viana \\& Liddle 1996, henceforth VL; Eke, Cole \\& Frenk 1996; Kitayama \\& Suto 1997). This is usually referred to as $\\sigma_{8}$, where $h$ is the present value of the Hubble parameter, $H_{0}$, in units of $100\\;{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$. However, the derived value of $\\sigma_{8}$ depends to a great extent on the present matter density in the Universe, parameterized by $\\Omega_{0}$, and more weakly on the presence of a cosmological constant, $\\Lambda$. The cleanest way of breaking this degeneracy is to include information on the change in the number density of rich galaxy clusters with redshift \\cite{FWED}, the use of X-ray clusters for this purpose having been proposed by Oukbir \\& Blanchard \\shortcite{OB} and subsequently further investigated \\cite{HM,OB97}. Several attempts have been made recently, with wildly differing results \\cite{Henry,FBC,Grossetal,BB,Ekeetal,Retal}.  The best method to find clusters of galaxies is through their X-ray emission, which is much less prone to projection effects than optical identification. Further, the X-ray temperature of a galaxy cluster is at present the most reliable estimator of its virial mass.  This can then be used to relate the cluster mass function at different redshifts, calculated for example within the Press--Schechter framework \\cite{PS,BCEK}, to the observed cluster X-ray temperature function. We can therefore compare the evolution in the number density of galaxy clusters seen in the data with the theoretical expectation for large-scale structure formation models, which depends significantly only on the assumed values of $\\Omega_{0}$ and $\\lambda_{0}\\equiv\\Lambda/3H^{2}_{0}$, the latter being the contribution of $\\Lambda$ to the total present energy density in the Universe.  However, the X-ray temperature of a cluster of galaxies is not an easily measurable quantity, as compared to its X-ray luminosity.  A minimum flux is required, so that there is a high enough number of photons for the statistical errors in the temperature determination to be reasonably small. Because of this, although estimates of the present-day cluster X-ray temperature function have been available since the early 90's \\cite{ESFA,HA}, the change in the cluster X-ray temperature function as we look further into the past has been much more difficult to determine.  Estimates for the X-ray temperatures of individual clusters with redshifts as high as 0.3 have been available for some years (e.g.~see David et al.~1993), but only with the advent of the {\\em ASCA} satellite has it been possible to measure X-ray temperatures for clusters of galaxies around that redshift in a systematic way, and to go to even higher redshifts.  The evolution of the cluster X-ray luminosity function with redshift, though easier to determine, provides much weaker constraints on $\\Omega_{0}$ and $\\lambda_{0}$, due to the fact that the X-ray luminosity of a galaxy cluster is not expected to be a reliable estimator of its virial mass (e.g.~Hanami 1993). Though it could in principle provide some indication of the change of the cluster X-ray temperature function with redshift, the problem is that not only is there considerable scatter in the present-day cluster X-ray temperature verses luminosity relation \\cite{Davetal,Fetal}, but it is also not known how the relation may change with redshift, though recently it has been argued that at least up to $z=0.4$ it does not seem to evolve \\cite{MScharf,AF,Retal,SBO}.  The deepest complete X-ray sample of galaxy clusters presently available is the one obtained from the {\\em Einstein Medium Sensitivity Survey} ({\\em EMSS}) \\cite{Getal,Hetal}.  This sample is restricted to objects with declination larger than $-40^{\\rm o}$ and is flux-limited, with $F_{{\\rm obs}}\\geq1.33\\times10^{-13}\\;{\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$, where $F_{{\\rm obs}}$ is the cluster flux in the 0.3 to 3.5 keV band which falls in a $2'.4\\times2'.4$ {\\em EMSS} detect cell.  It presently contains 90 clusters of galaxies, after a few misidentifications were recently removed \\cite{GioiaL,Netal}.  This is the only complete galaxy cluster catalogue beyond a redshift of 0.3, and as such unique in providing the means to distinguish between different possible values for $\\Omega_{0}$ and $\\lambda_{0}$. However, until the recent effort by Henry \\shortcite{Henry}, very few X-ray temperatures were known for those galaxy clusters in the {\\em EMSS} sample with redshifts exceeding 0.15 (see Sadat et al.~1998 for a recent compilation).  Henry \\shortcite{Henry} used {\\em ASCA} to observe all galaxy clusters in the {\\em EMSS} cluster sample with $0.3\\leq z \\leq 0.4$ and $F_{{\\rm obs}}\\geq2.5\\times10^{-13}\\;{\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$. The resulting sub-sample of 10 clusters has a median redshift of 0.32, and the data obtained for each cluster, the X-ray flux, luminosity and temperature, can be found in Table 1 of Henry (1997).  We will use this data together with the present-day (median redshift 0.05) cluster X-ray temperature function. We work within the extended Press--Schechter formalism proposed by Lacey \\& Cole (1993, 1994), which allows an estimation of the formation times of dark matter halos.  We will assume the dark matter to be cold, and consider the cases of an open universe, where the cosmological constant is zero, and a spatially-flat universe, such that $\\lambda_{0}=1-\\Omega_{0}$. ",
        "conclusions": "{}From the above analysis, we conclude that {\\em at present} it is not possible to reliably exclude any interesting value for $\\Omega_{0}$ on the basis of X-ray cluster number density evolution alone, due to the limited statistical significance of the available observational data and to uncertainties in the theoretical modelling of cluster formation and evolution.  However, we do find that values of $\\Omega_{0}$ below 0.3 are excluded at least at the 90 per cent confidence level.  Values of $\\Omega_{0}$ between 0.7 to 0.8 are those most favoured, though not strongly.  These results are basically independent of the presence or not of a cosmological constant.  Our conclusions support those of Colafrancesco, Mazzotta \\& Vittorio \\shortcite{CMV}, who tried to estimate the uncertainty involved in the estimation of the cluster X-ray temperature distribution function at different redshifts based on its present-day value. They found this uncertainty, given the still relatively poor quality of the data, to be sufficiently large to preclude the imposition of reliable limits on the value of $\\Omega_{0}$.  Our results disagree with those of Henry \\shortcite{Henry} and Eke et al.~\\shortcite{Ekeetal}, as they found the preferred $\\Omega_{0}$ to lie between 0.4 to 0.5, with the $\\Omega_{0}=1$ hypothesis strongly excluded. This disagreement is mainly the consequence of our focus on the threshold X-ray temperature of 6.2 keV, while they draw their conclusions based on the analysis of the results obtained for several threshold X-ray temperatures. Further below we will repeat our calculations assuming a threshold X-ray temperature of 4.8 keV, and we will find that when we calculate the joint probability of some value for $\\Omega_{0}$ being excluded on the basis of the results concerning either one or both threshold X-ray temperatures of 6.2 keV and 4.8 keV, the favoured value for $\\Omega_{0}$ decreases to around 0.55. Some of the reasons for our choice of deriving the conclusions solely based on the results obtained for the 6.2 keV threshold were mentioned at the end of subsection 3.1 and others will be detailed below.  Other less important contributions to the difference between our results and those presented by Henry \\shortcite{Henry} and Eke et al.~\\shortcite{Ekeetal} are the different assumed normalization for the virial mass to X-ray temperature relation, and the corrections in the expected values in the Universe for both $N(>6.2\\,{\\rm keV},\\,0.05)$ and $N(>6.2\\,{\\rm keV},\\,0.32)$ due to the uncertainties in the X-ray cluster temperature measurements. Note that changing the mean of the bootstrap distribution obtained for $N(>6.2\\,{\\rm keV},\\,0.32)$ to its theoretically-expected overall value in some $\\Omega_{0}$ universe and then calculating the exclusion level on the estimated value for $N(>6.2\\,{\\rm keV},\\,0.32)$ in the Universe given the dataset in Henry \\shortcite{Henry}, rather than just using the original bootstrap distribution to impose an exclusion level on the theoretically-expected overall value for $N(>6.2\\,{\\rm keV},\\,0.32)$ in that $\\Omega_{0}$ universe, does not seem to make much difference. This is a reflection of the fact that the bootstrap distributions recovered do not have a strongly asymmetric shape.  Our disagreement with Eke et al.~\\shortcite{Ekeetal} on the level of exclusion of the $\\Omega_{0}=1$ hypothesis is also due to our much larger assumed uncertainty in the theoretically-expected overall value for $N(>6.2\\,{\\rm keV},\\,0.32)$.  For the $\\Omega_{0}=1$ hypothesis to be favoured, one requires the lowest possible observed value for $N(>6.2\\,{\\rm keV},0.32)$. This is best achieved if, for the sample of 10 galaxy clusters used in its calculation, the X-ray temperatures turn out to be on average lower than the assumed mean, and the X-ray fluxes higher. A higher ratio between the extended and detect cell fluxes for the {\\em EMSS} at $z=0.32$ would also help.  On the theoretical side, the higher one decides the expected value for $N(>6.2\\,{\\rm keV},0.32)$ is, the more compatible with the data the $\\Omega_{0}=1$ hypothesis becomes.  This can be best achieved if, in decreasing order of importance, the cluster virial mass at fixed X-ray temperature is being underestimated, $\\delta_{{\\rm c}}$ is lower than the canonical value 1.7 and $f$, the assembled fraction of a cluster virial mass after which the X-ray temperature does not change significantly, is assumed greater than 0.75.  However, the single most important factor in determining the theoretically-expected overall value for $N(>6.2\\,{\\rm keV},0.32)$ is the present-day normalization for the dispersion of the density field, $\\sigma_{8}$, which in turn results from the observational value for the present density $N(>6.2\\,{\\rm keV},\\,0.05)$.  Although we worked with all X-ray clusters that make up the dataset in Henry \\shortcite{Henry}, and even estimated the effect of also considering the 5 clusters with lower X-ray fluxes present in the {\\em EMSS} in the redshift bin from 0.3 to 0.4, in fact we only used the abundance of clusters with X-ray temperatures in excess of 6.2 keV to constrain $\\Omega_{0}$. We mentioned some of the reasons for this choice in Section~3. Nevertheless, we decided to repeat the same calculations for a threshold X-ray temperature of 4.8 keV. This value also well represents the mean curve going through the observed cumulative X-ray temperature distribution function at both $z=0.05$ and $z=0.32$.  \\begin{figure} \\centering \\leavevmode\\epsfysize=5.4cm \\epsfbox{zclus_fig4.eps}\\\\ \\caption[Figure4]{The absolute exclusion levels for different values of $\\Omega_{0}$ in both the open and spatially-flat cases, when the threshold X-ray temperature of 4.8 keV is used.} \\end{figure}  The results regarding the best-fit value for $\\Omega_{0}$, presented in Figure~4, are somewhat different from those we obtained when the threshold X-ray temperature was assumed to be 6.2 keV.  This is particularly true if the correction for the possibility of any of the 5 clusters with the lowest X-ray fluxes in the $0.3<z<0.4$ EMSS sub-sample having X-ray temperatures in excess of 4.8 keV is included, as can be seen in Figure 5 for the open case. While the standard analysis without these 5 X-ray clusters prefers a value for $\\Omega_{0}$ between 0.4 to 0.5, when the correction for the scatter in the relation between the cluster X-ray temperature and luminosity is included, in the way described in subsection 3.3, the preferred value for $\\Omega_{0}$ decreases to about 0.3.  Now the $\\Omega_{0}=1$ hypothesis is excluded at more than the 95 per cent confidence level, with or without the correction.  At the 90 per cent confidence level, one finds that $\\Omega_{0}>0.8$ is excluded without the correction, being this limit lowered to 0.7 when the correction is included.  One can also estimate the joint probability of some $\\Omega_{0}$ value being excluded on the basis of the results relative to either one or both X-ray temperature thresholds. Assuming the data used in the calculations for the two thresholds is independent, the results then imply that the favoured value for $\\Omega_{0}$ is close to 0.55 (0.50 if the incompleteness correction is included) and the $\\Omega_{0}=1$ hypothesis is excluded at the 99 per cent level. This agrees very well with the results of Henry \\shortcite{Henry} and Eke et al.~\\shortcite{Ekeetal}, leading us to believe that the main difference between our analysis and theirs is our decision to draw our conclusions solely based on the exclusion levels obtained for the X-ray temperature threshold of 6.2 keV.  A further potential problem one must consider when working with clusters whose observed X-ray temperature is as low as 4.8 keV is the possibility that the energy in the intracluster gas has increased as a result of (pre-)heating by supernovae and starbursts in the cluster galaxies.  In fact this is the leading hypothesis (e.g.  Navarro, Frenk \\& White 1995; Markevitch 1998) put forward to explain the discrepancy between the observed slope of the X-ray temperature--luminosity relation, close to 0.3, and the expected value of 0.5 if clusters evolve in a self-similar way \\cite{Kaiser}.  Following Eke et al.~\\shortcite{Ekeetal}, we assume that in a cluster whose observed X-ray temperature is 4.8 keV, 17 per cent of its energy, that is 0.8 keV per intracluster gas particle, was due to (pre-)heating produced by processes occurring inside the cluster galaxies. This is approximately the amount of energy that gets injected into the intracluster gas particles in the simulation of Metzler \\& Evrard \\shortcite{ME}, where a galaxy cluster's X-ray temperature, which would otherwise be 5.6 keV, increased to 6.4 keV. Note however that in the scheme proposed by Eke et al.~\\shortcite{Ekeetal} a cluster this large would not be (pre-)heated to the extent simulated by Metzler \\& Evrard \\shortcite{ME}, as in their proposal Eke et al.~\\shortcite{Ekeetal} assume that the energy gained by each intracluster gas particle due to (pre-)heating decreases as a galaxy cluster becomes larger, being close to zero for galaxy clusters with X-ray temperatures exceeding 6.2 keV.  \\begin{figure} \\centering \\leavevmode\\epsfysize=5.4cm \\epsfbox{zclus_fig5.eps}\\\\ \\caption[Figure5]{The absolute exclusion levels for different values of $\\Omega_{0}$ only for the open case, when the threshold X-ray temperature of 4.8 keV is used. The full curve includes a correction (FC) for the possibility of the 5 clusters with lowest fluxes in the {\\em EMSS} located between $z=0.3$ and $z=0.4$ having X-ray temperatures in excess of 4.8 keV. The dashed curve includes a correction (HC) for the possibility of (pre-)heating of the intracluster medium due to processes within the cluster galaxies. The dotted curve includes both corrections.} \\end{figure}  The above assumption means that the observed values for $N(>4.8\\,{\\rm keV},\\,z)$, when $z=0.05$ and $z=0.32$, should now be compared with the theoretically-expected values for $N(>4.0\\,{\\rm keV},\\,z)$ at those redshifts.  The resulting exclusion levels on the value of $\\Omega_{0}$ can be seen in Figure 5 for the open case.  There is little difference compared to the results in Figure 4 that follow from the standard no-heating calculation.  The lower value for $\\sigma_{8}$, required to match theory and observations at $z=0.05$, more than compensates for the expected increase in the number of galaxy clusters with X-ray temperatures in excess of 4.8 keV at $z=0.05$, in effect bringing this number down.  In fact, the standard no-heating calculation for a threshold X-ray temperature of 4.8 keV requires a value for $\\sigma_{8}(\\Omega_{0})$, so that the observed value for $N(>4.8\\,{\\rm keV},\\,0.05)$ is reproduced, that is less than 3 per cent below that required by the $>6.2$ keV data, quoted in equation (\\ref{final1}).  On the other hand, including the (pre-)heating correction, the required $\\sigma_{8}(\\Omega_{0})$ value drops to 19 per cent below that preferred by the $>6.2$ keV data. Though the coincidence between the $\\sigma_{8}$ values obtained for the two X-ray temperature thresholds 4.8 keV and 6.2 keV under the no-heating assumption may be accidental, it could indicate that (pre-)heating was relatively unimportant at least for the galaxy clusters observed at $z=0.05$ with X-ray temperatures exceeding 4 keV. If (pre-)heating was more important in the past than today, then the required $\\sigma_{8}(\\Omega_{0})$ value would be that obtained through the standard no-heating hypothesis, but the comparison at $z=0.32$ would include the (pre-)heating correction.  This would push the theoretically-expected value for $N(>4.8\\,{\\rm keV},\\,0.32)$ up, favouring higher values for $\\Omega_{0}$.  This is not as far-fetched as it may seem, given that it is well known that the star-formation rate peaks before $z=1$ (e.g.  Madau, Ferguson \\& Dickinson 1998; Baugh et al.  1998), and consequently so does the rate of supernovae Type II (the rate of supernovae Type Ia peaks a few Gyr later) and the probability of starbursts.  The results for the 4.8 keV threshold X-ray temperature are close to those found by Eke et al.~\\shortcite{Ekeetal}, leading us to believe that their exclusion levels for $\\Omega_{0}$ are dominated by the information associated with the threshold X-ray temperatures 4.0 keV and 5.0 keV. In our view, the analysis for these X-ray temperature thresholds carries with it a sufficient number of uncertainties, due to the problems mentioned above, so as to render the constraints imposed on $\\Omega_{0}$ not very trustworthy.  Only the data regarding clusters with X-ray temperatures in excess of about 6 keV seems sufficiently free of modelling problems so as to be potentially useful in constraining $\\Omega_{0}$.  Another possible complication has arisen from recent work by Blanchard, Bartlett and Sadat \\shortcite{BBS} who use a sample of 50 galaxy clusters with mean redshift of 0.05, which were identified through the {\\em ROSAT} satellite, to estimate the cumulative X-ray temperature distribution function at $z=0.05$. They claim the number density of galaxy clusters at $z=0.05$ with X-ray temperatures exceeding 4 keV is being {\\em underestimated} when the Henry \\& Arnaud cluster sample is used. Through the X-ray cumulative temperature distribution function at $z=0.05$ they obtain, they then estimate $\\Omega_{0}$ using the {\\em EMSS} cluster abundance in the redshift bin $0.3<z<0.4$ and the X-ray temperature data gathered in Henry \\shortcite{Henry}. They find the favoured value for $\\Omega_{0}$ to be 0.75, while $\\Omega_{0}<0.3$ is excluded at more than the 95 per cent level. These results coincide very well with ours when only the 6.2 keV threshold X-ray temperature is considered, thus perhaps implying that the discrepancy between the favoured value for $\\Omega_{0}$ found when different X-ray temperature thresholds are considered may arise from a underestimation of the cumulative distribution function at $z=0.05$ for X-ray temperatures below about 6 keV.  Unfortunately, due to uncertainties associated both with the observational measurements and the theoretical modelling of cluster evolution, the presently-available data on galaxy clusters with X-ray temperatures exceeding about 6 keV is not able to strongly discriminate between cosmologies with different values for $\\Omega_{0}$. And in any case, the data available is probably not yet statistically significant.  More is needed to support or disclaim the preliminary conclusions that can be obtained from it.  In particular there are some oddities with the sub-sample of {\\em EMSS} galaxy clusters observed by Henry, such as the strange redshift distribution, strongly clustered around $0.32$, and the unexpectedly low X-ray temperature of MS2137.3, that makes one have some doubts about how representative this dataset is of the Universe.  Within the next few years, with the launch of the {\\em XMM} satellite, possibly in late 1999, a significant increase in the quantity and quality of the available data is expected to occur \\cite{Romer}. It should then be possible to place stronger constraints on $\\Omega_{0}$ on the basis of the evolution of the galaxy cluster X-ray temperature function. This would be helped by improvements in the theoretical modelling of cluster evolution, perhaps based on the high-resolution hydrodynamical $N$-body simulations on cosmological scales expected in the near future."
    },
    "9803/astro-ph9803323_arXiv.txt": {
        "abstract": "BATSE, Ulysses, and TGRS and KONUS on WIND detected four gamma-ray events within 1.8 days during 1996 October 27-29, consistent with coming from the same location on the sky.  We assess the evidence that these events may be due to a series of bursts from a single source by calculating the probability that such a clustering in position and in time of occurrence might happen by chance.  The calculation of this probability is afflicted by the usual problem of a posteriori statistics.  We introduce a clustering statistic, which is formed from the \"minimum circle radius\" (i.e. the radius of the smallest circle that just encloses the positions of all the events) and the minimum time lapse (i.e. the time elapsed between the first and last event). We also introduce a second clustering statistic, which is formed from the \"cluster likelihood function\" and the minimum time lapse. We show that the use of these statistics largely eliminates the \"a posteriori\" nature of the problem.  The two statistics yield significances of the clustering of $3.3\\times 10^{-4}$ and $3.1\\times 10^{-5}$, respectively, if we interpret the four events as four bursts, whereas the clustering is not significant if we interpret the four events as only three bursts. However, in the latter case one of the bursts is the longest ever observed by BATSE. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803115_arXiv.txt": {
        "abstract": "{   We have used the results of recent smoothed particle hydrodynamic simulations of colliding stars to create models appropriate for input into a stellar evolution code.  In evolving these models, we find that little or no surface convection occurs, precluding angular momentum loss via  a magnetically-driven stellar wind as a viable mechanism for slowing rapidly rotating blue stragglers which have been formed by collisions. Angular momentum transfer to either a circumstellar disk (possibly collisional ejecta) or a nearby companion are plausible mechanisms for explaining the observed low rotation velocities of blue stragglers  Under the assumption that the blue stragglers seen in NGC 6397 and 47 Tuc have been created solely by collisions, we find that the majority of blue stragglers cannot have been highly mixed by convection or meridional circulation currents at anytime during their evolution.  Also, on the basis of the agreement between the predictions of our non-rotating models and the observed blue straggler distribution, the evolution of blue stragglers is apparently not dominated by the effects of rotation.  }      \\newpage ",
        "introduction": "Blue stragglers lie roughly along an extension of a star cluster's main sequence (MS) and are generally bluer and brighter than the turn-off (TO) stars.  While several possible mechanisms for their formation have been proposed (e.g. delayed star formation, binary mass transfer, binary coalescence, stellar collisions; see Livio, 1993, and Stryker, 1993, for reviews), \\nocite{livio93}\\nocite{stryker93} the collision scenario, in which initially unbound stars come into contact and merge during a dynamical encounter, has received the most attention recently (e.g. Sandquist \\hbox{\\em et al.} 1997; Ouellette \\& Pritchet 1996; Sills, Bailyn \\& Demarque 1995). \\nocite{sbh97}\\nocite{op96}\\nocite{sbd95} Blue stragglers formed by collisions are of particular interest because their formation rate tells us about the stellar interaction rate in their cluster environment and the dynamical history of the cluster itself (Bailyn, 1995).\\nocite{bailyn95}   There are at least two ways to investigate the formation rate of blue stragglers by collisions.  One is to perform N-body simulations, modelling the stellar environment and determining the collision rate directly. This requires knowledge of the stellar interaction cross-section, local stellar density, mass function, and binary fraction.  The second is to model the structure of the merger remnants and to evolve them using a stellar evolution code.  Just as  the age of a cluster is found by comparing theoretical isochrones, based on standard stellar models,  with an observed colour-magnitude diagram (CMD), comparison of the evolution of the merger models to the observed numbers and distribution of blue stragglers in the CMD can tell us about the distribution of lifetimes, and the production rates, of these stars.   In this paper, we develop and evolve models of collisional merger remnants and compare the predictions of these models with the observed blue straggler distributions in NGC 6397 and 47 Tuc. ",
        "conclusions": "We have evolved stellar models which are appropriate for stars created during collisions between equal and unequal mass stars.  In doing so we have used the results of SPH simulations of colliding polytropes, most importantly the predictions of entropy stratification.  To produce these models, we have made several assumptions to facilitate the incorporation of the results of the SPH simulations of collisions between {\\em polytropic} stars into models of real stars; these assumptions include neglecting mass loss, rotation, and shock heating during the collision.  Because we have restricted our attention to head-on parabolic collisions, we believe these approximations to be reasonably valid;  mass loss and rotational velocity are shown to be small in this case by the simulations of LRS; shock heating is shown to be small and, most importantly, should not affect the distribution of helium-rich gas in the merger remnant (Sec~\\ref{sec:shock}), although it is likely that shock heating will affect the distribution of material near the surface of the star.  The form of the energy injection term $\\epsilon_x$ which we use to expand our merger remnant models is constrained by the results of SPH, but its form is nonetheless not crucial as long as significant mixing is not artificially introduced (via the Vogt-Russell Theorem; Vogt 1926, Russell 1927).  In comparing the predictions of these models with the blue stragglers seen in two dense clusters, NGC 6397 and 47 Tuc, we have also assumed that all of the blue stragglers in these clusters have been formed through collisions.  The apparent agreement between the predictions of our unequal mass merger models and the observed blue straggler distributions  NGC 6397 and 47 Tuc suggests that \\begin{itemize} \\item little or no mixing occurred during the formation of these blue stragglers, either of fresh hydrogen to the cores, or of  helium to the surfaces.  As explained earlier, such mixing would produce a significant blueward scatter in the distribution of blue stragglers in the CMD --- as it is, our models with the best agreement with the observations, the unequal mass mergers, produce rather red stars, relative to the clusters' ZAMS.  Regardless of the formation mechanism for the blue stragglers in these clusters, it is apparent from their location redward of the ZAMS that they are formed as {\\em evolved} stars (Ouellette \\& Pritchet, 1996). \\item there may be some form of efficient AML mechanism acting to slow the rotation of these stars.  Blue stragglers which are formed by collisions are expected to be rapidly rotating: the affect of rapid rotation on the observed distribution of blue stragglers in the CMD should result in a poor agreement with our non-rotating models, and yet the agreement appears to be quite good.   The thin, short-lived convective envelopes seen in our models precludes a magnetically driven stellar wind AML mechanism to explain the slow rotation rates and apparent lack of rotational effects; however, angular momentum transfer to a circumstellar disk or nearby companion are still plausible.   \\end{itemize}    The lack of surface convection in our most massive models and the thin convective envelopes seen in our less massive models means that surface abundances of collisionally merged blue stragglers might not be altered from those of the parent stars, unless some amount of hydrodynamical mixing occurs during the collision.  Additionally, meridional circulation currents, which should be confined to the stellar envelope except in the most extreme cases, may act to mix the surface layers, regardless of the amount of convection.   The extreme blue stragglers observed in both NGC 6397 and 47 Tuc occupy the region of the CMD which should be populated by highly mixed stars.  It is possible that these blue stragglers are the few that were born rotating rapidly enough that  the meridional circulation currents are able to penetrate more deeply in to the star; if so, these stars should still be rapidly rotating.  It is, however, possible that some fraction of these stars are merely photometric errors."
    },
    "9803/astro-ph9803265_arXiv.txt": {
        "abstract": "s{Fourteen classical double radio galaxies with redshifts between zero and two were used to determine the cosmological parameters $\\Omega_m$, $\\Omega_{\\Lambda}$, and $\\Omega_k$, where these are the normalized values of the mean mass density, cosmological constant, and space curvature at the present epoch.  A low value of $\\Omega_m$ is obtained, and $\\Omega_m = 1$ is ruled out with 97.5 \\% confidence.  The low value of $\\Omega_m$ determined using the radio source method described here is also indicated by several independent tests.  Thus, it appears that either a cosmological constant, or space curvature, is significant at the present epoch. This means that the universe is undergoing, or has recently undergone, a transition away from a state of matter domination and into a state where either a cosmological constant or space curvature is determining the expansion rate of the universe.  The low value of $\\Omega_m$ presented here and by Guerra \\& Daly (1998) means that we can state with 97.5 \\% confidence that the universe will continue to expand forever.  } ",
        "introduction": "The structure and fate of the universe can be described by the cosmological parameters $\\Omega_m$, $\\Lambda$, and $k$, where $\\Omega_m$ is the average density of matter in the universe at the present time divided by the critical density, $\\Lambda$ is the current value of the cosmological constant (generally, though not always, taken to be time-independent), and $k$ describes the global geometry of the universe.  The normalized values of $\\Lambda$ and $k$ are denoted $\\Omega_{\\Lambda}$, and $\\Omega_k$ (e.g. Peebles 1993).  Assuming that any relativistic component, such as the microwave background radiation, has a negligible contribution to the total mass-energy density at the present time, the following equation must be satisfied: $\\Omega_m$ + $\\Omega_{\\Lambda}$ + $\\Omega_k$ = 1.  As originally discussed by Dicke (1970), and later by Peebles (1993), each term evolves differently with redshift, so it is unlikely that two terms will be comparable at any given time. However, if $\\Omega_m < 1$, then, the universe is undergoing a transition away from a state of matter domination into a state where either space curvature or the cosmological constant is dominant.  Following Dicke's argument, it is rather unlikely that all three terms will be significant at the present time.  The cosmological parameters can only be believably determined when several independent methods of estimating the parameters all yield similar values, both within each category of test, and comparing results from different categories. At present, there are three main categories of estimates of $\\Omega_m$: (1) local, low-redshift dynamical tests; (2) tests that depend on the coordinate distance to high-redshift sources through the angular size distance or the luminosity distance; and (3) tests using the fluctuations of the microwave background radiation on different angular scales.  At the present time, values of $\\Omega_m$ have been determined using methods (1) and (2); results obtained using method (1) are mentioned below.  Method (2) has been applied to radio galaxies with redshifts between zero and two by Daly (1994, 1995), Guerra \\& Daly (1996, 1998), and is applied here to radio sources with redshifts between zero and two.  Method (2) has also been applied to supernovae by Garnavich {\\it et al.} (1998) and Perlmutter {\\it et al.} (1998), who study sources to a redshift of about one. At some point in the future, there may be independent constraints from measurements of fluctuations of the microwave background radiation on different angular scales.  Several local, low-redshift estimates of $\\Omega_m$ indicate that it is significantly less than unity (Hudson {\\it et al.} 1995; Shaya, Peebles, \\& Tully 1995; Carlberg {\\it et al.}\\ 1997; Bahcall \\& Fan 1998). These tests are interesting and important, though they are all subject to caveats. Many depend on whether the mass is distributed like the light, known as ``biasing,\" and many local measurements only indicate the amount of mass that is clustered on the scales of galaxies and clusters of galaxies.  The concordance of many local, low-redshift tests suggests that the amount of mass that clusters with galaxies is significantly less than the critical value.  Estimates of cosmological parameters that utilize the coordinate distance to sources at high redshift (category [2] defined above), such as tests involving angular size distance or luminosity distance, are fundamentally different from local, low-redshift tests.  The value of $\\Omega_m$ estimated through the coordinate distance, the angular size distance, or the luminosity distance, is truly the global value of $\\Omega_m$. The test is independent of any biasing of matter relative to light, of how or whether the mass is clustered, of the nature of the dark matter (e.g. baryonic, cold dark matter, hot dark matter, etc.), and of the origin of fluctuations that lead to galaxy formation (e.g. cosmic strings, textures, adiabatic or isocurvature fluctuations, etc.).  Two tests currently being used to determine global cosmological parameters through the coordinate distance to high-redshift sources (category [2] defined above) are the radio source method, and the supernova method. ",
        "conclusions": "The constraints on cosmological parameters obtained here arise primarily from the relative position of the data points as a function of redshift, and are not sensitive to any particular point, or to the normalization of the curves (which is left as a free parameter in the fits). For example, if the one low-redshift point near z = 0 is excluded, the results do not change (see Guerra \\& Daly 1998).  It is clear from figure 1a that the radio source model is working extremely well; the data fall right along the cosmological curves all the way from zero redshift to a redshift of about two.  It is clear from figure 1b that, given this data, it is very unlikely that $\\Omega_m=1$.  The results presented here are very similar to those obtained by the supernovae groups.  Any potential problems, such as dust extinction or evolution, are completely different for the two methods.  The fact that they yield nearly identical results suggests that both are correct.  In addition, the low value of $\\Omega_m$ obtained using this method agrees with the low values indicated by local dynamical tests, which further suggests that $\\Omega_m$ is low, and the universe will continue to expand forever.  We are in the process of obtaining radio information on 67 additional radio galaxies with redshifts between zero and two.  With these additional sources we will be able to constrain cosmological parameters to very high accuracy."
    },
    "9803/astro-ph9803053_arXiv.txt": {
        "abstract": "We show that the kinematics of the shells seen around some elliptical galaxies provide a new, independent means for measuring the gravitational potentials of elliptical galaxies out to large radii.  A numerical simulation of a set of shells formed in the merger between an elliptical and a smaller galaxy reveals that the shells have a characteristic observable kinematic structure, with the maximum line-of-sight velocity increasing linearly as one moves inward from a shell edge.  A simple analytic calculation shows that this structure provides a direct measure of the gradient of the gravitational potential at the shell radius.  In order to extract this information from attainable data, we have also derived the complete distribution of line-of-sight velocities for material within a shell; comparing the observed spectra of a shell to a stellar template convolved with this distribution will enable us to measure the gradient of the potential at this radius.  Repeating the analysis for a whole series of nested shells in a galaxy allows the complete form of the gravitational potential as a function of radius to be mapped out.  The requisite observations lie within reach of the up-coming generation of large telescopes. ",
        "introduction": "The gravitational potentials of elliptical galaxies have proved difficult to derive, especially at large radii.  In disk galaxies, both stars and gas have relatively simple orbit structures dominated by nearly circular orbits, and it is therefore straightforward to deduce the underlying gravitational potential (at least in the disk plane) from the observed kinematics in these systems.  However, the corresponding orbital structure in elliptical galaxies is much more complicated, making the derivation of the gravitational potential from observed kinematics in such a system far from simple.  In fact, basic kinematic observations of projected density and line-of-sight velocity dispersion are not sufficient to solve unambiguously for both the functional form of the gravitational potential and the distribution of stellar orbits in elliptical galaxies (Binney \\& Mamon 1982).  This ambiguity has recently been partially resolved by using the extra kinematic information that can be derived from the shapes of line profiles in high quality spectral data (e.g. Gerhard 1993, Carollo et al.\\ 1995). However, these results are restricted to the inner few effective radii of the stellar light; it is hard to extend these methods to much larger radii because of the low surface brightnesses of galaxies further out.  Independent mass measurements using gravitational lensing (e.g. Kochanek 1995) are restricted to still smaller radii. The only technique that reaches to large radii comes from mapping the distribution of hot gas around some ellipticals (Buote \\& Canizares 1997). This method relies on the assumption of hydrostatic equilibrium in the analysis of X-ray intensity and temperature profiles, from which the gravitational potential may be deduced. The requisite data are hard to obtain with current X-ray telescopes, and the analysis is subject to possible systematic uncertainties if the hot gas is not in a single phase. Furthermore, many of the X-ray halos may be associated with the group hosting the elliptical rather than with the galaxy itself.  It is therefore of interest to develop alternative probes of the mass distribution in the outer parts of elliptical galaxies.  If we wish to measure the gravitational potential of an elliptical galaxy unambiguously using kinematic tracer particles, then we need a set of test bodies whose orbital structure is simple and observationally well-constrained.  In this paper, we show that the large, regular systems of faint shells seen in some elliptical galaxies (Schweizer 1980, Malin \\& Carter 1980) offer just such a tracer.  These shells are believed to be the remnants of a small galaxy after a head-on collision with a larger system -- any other collision geometry does not result in extensive, nested shell systems. From the geometry of the merger, we know that all the stars in each shell must be on radial orbits with very similar energies.  As we show below, this particularly-simple orbital structure means that the kinematics of the shells provide a useful new probe of the gravitational potential out to large radii in elliptical galaxies. ",
        "conclusions": "In this paper, we have set out to show how the kinematics of the shells produced when a small galaxy merges with a large system can be used to measure the underlying gravitational potential.  As mentioned earlier in the paper, although the minor merger model is the most widely accepted interpretation of extensive shell systems, several other possible explanations have been advanced.  Thus, not only would the detection of the characteristic kinematic structure illustrated in Fig.~\\ref{fig-shellxw} allow us to measure the gravitational potential, but it would also confirm the origin of the shells themselves.  An alternative method for using shells to constrain the form of the gravitational potential was originally proposed by Quinn (1984).  He suggested that use be made of the distribution of shells with radius. Recall that, at any given time, the photometrically-observed shells are delineated by those stars that have completed integer number of radial oscillations.  The ratio of the radii of successive shells therefore gives the oscillation periods at different energies, which is directly related to the form of the potential.  The constraints on the potential obtained from such analyses are non-dimensional; they could, for example, provide information about the power-law index of the potential, but they cannot yield a mass normalization.  By contrast, the kinematic method described here relies on measured velocities, which can be translated directly into the fully-dimensional gravitational potential.  The shell counts have also been shown to be susceptible to uncertainties associated with dynamical friction acting on the shredding satellite, since different shells will have followed substantially different orbits, with different numbers of passages through the center of the host galaxy. Quinn's method also requires that we know how many radial oscillations the outermost shell has executed, which can only be inferred indirectly.  The method discussed here, on the other hand, only uses information from each shell singly, and only uses the fact that phase wrapping is an efficient mechanism for concentrating stars of very nearly equal energies.  Our analysis has assumed that the merger occurred on an exactly radial orbit. To test the importance of this assumption, we have repeated our simulations in the isochrone potential for satellite galaxies on slightly non-radial orbits. These calculations revealed that even in systems where the encounter is sufficently off-axis to produce shells that do not appear regular and aligned, the line profiles of individual shells are still accurately reproduced by the radial model calculations.  Thus, the kinematics of an aligned shell system will almost certainly be well-modelled by the simple radial orbit approximation.  Furthermore, altering the viewing angle of the shell system tends to diminish the shell visibility before altering the velocity structure, again suggesting that by selecting aligned, regular shell systems, we will preferentially find systems to which the above analysis is applicable.  It should be noted, however, that the tests we have carried out have all been based on the isochrone potential, since the orbit calculations are computationally simple in a potential of this form.  It would be interesting to carry out a more extensive numerical study of shells in other potentials in order to check that they, too, are insensitive to the assumption of radial orbits.  Any flattening of the host galaxy's potential is also a possible problem.  It may introduce orientation-dependent effects, as well as modify the orbit structure of the stars making up each shell. Simulations of the type described earlier in flattened potentials show that regular shell systems only occur for nearly radial encounters in spherical potentials: a degree of flattening as small as 5\\% in the potential will cause significant beating between the motions in the various directions, and leads to shells with alternately large and small opening angles. The velocity structure of such shells is also rather disturbed, and it is not clear that the velocity profiles will provide much information about the overall potential. However, the detection of a regular uniform shell system provides us with strong {\\it a priori} evidence that the potential must be very close to spherically-symmetric.  Shells are photometrically faint structures, and so obtaining the kinematic measurements that we advocate here will not be simple.  It is, however, noteworthy from a comparison of Figs.~\\ref{fig-shellxy} and \\ref{fig-shellxw} that the kinematic observation of these systems significantly enhances their contrast against any background.  It is also worth noting that the brightnesses of observed shells do not decrease very rapidly with radius, and so, unlike other techniques for measuring gravitational potentials using stellar kinematics, the difficulty of applying this method does not increase prohibitively with radius.  We have carried out signal-to-noise ratio calculations for some of the brighter shell galaxies such as NGC~3923, and have ascertained that data of the requisite quality could be obtained with a couple of nights' integration using a 4-metre telescope.  Clearly, projects of this nature will prove quite feasible with the up-coming generation of large telescopes."
    },
    "9803/astro-ph9803047_arXiv.txt": {
        "abstract": "For faint photometric surveys our ability to quantify the clustering of galaxies has depended on interpreting the angular correlation function as a function of the limiting magnitude of the data. Due to the broad redshift distribution of galaxies at faint magnitude limits the correlation signal has been extremely difficult to detect and interpret. We introduce a new technique for measuring the evolution of clustering. We utilize photometric redshifts, derived from multicolor surveys, to isolate redshift intervals and calculate the evolution of the amplitude of the angular 2-pt correlation function. Applying these techniques to the the Hubble Deep Field we find that the shape of the correlation function, at $z=1$, is consistent with a power law with a slope of $-0.8$. For $z>0.4$ the best fit to the data is given by a model of clustering evolution with a comoving \\rn = 2.37 \\Mpc\\ and $\\epsilon = -0.4^{+0.37}_{-0.65}$, consistent with published measures of the clustering evolution. To match the canonical value of \\rn = 5.4 \\Mpc, found for the clustering of local galaxies, requires a value of $\\epsilon = 2.10^{+0.43}_{-0.64}$ (significantly more than linear evolution). The log likelihood of this latter fit is 4.15 less than that for the \\rn = 2.37 \\Mpc\\ model. We, therefore, conclude that the parameterization of the clustering evolution of $(1+z)^{-(3+\\epsilon)}$ is not a particularly good fit to the data. ",
        "introduction": "The evolution of the clustering of galaxies as a function of redshift provides a sensitive probe of the underlying cosmology and theories of structure formation. In an ideal world we would measure the spatial correlation function of galaxies as a function of redshift and type and use this to compare with the predictions of different galaxy formation theories.  Observationally, however, our ability to efficiently measure galaxy spectra falls rapidly as a function of limiting magnitude and consequently we are limited to deriving spatial statistics from small galaxy samples and at relatively bright magnitude limits (e.g.\\ $I_{AB} < 22.5$, Le Fevre et al.\\ 1996, Carlberg et al.\\ 1997).  To increase the size of the galaxy samples and thereby reduce the shot noise the standard approach has been to measure the angular correlation function, i.e.\\ the projected spatial correlation function (Brainerd et al.\\ 1996, Woods and Fahlman 1997). While this allows us to extend the measure of the clustering of galaxies to fainter magnitude limits ($R<29$, Villumsen et al.\\ 1997) it has an associated limitation. For a given magnitude limit the amplitude of the angular correlation function is sensitive to the width of the galaxy redshift distribution, N(z). At faint magnitude limits N(z) is very broad and consequently the clustering signal is diluted due to the large number of randomly projected pairs.  In this letter we introduce a new approach for quantifying the evolution of the angular correlation function; we apply photometric redshifts (Connolly et al.\\ 1995, Lanzetta et al.\\ 1996, Gwyn and Hartwick 1996, Sawicki et al.\\ 1997) to isolate particular redshift intervals. In so doing we can remove much of the foreground and background contamination of galaxies and measure an amplified angular clustering. We discuss here the particular application of this technique to the Hubble Deep Field (HDF; Williams et al.\\ 1996). ",
        "conclusions": "Photometric redshifts provide a simple statistical means of directly measuring the evolution of the clustering of galaxies. By isolating narrow intervals in redshift space we can reduce the number of randomly projected pairs and detect the clustering signal to high redshift and faint magnitude limits. Applying these techniques to the HDF we can characterize the evolution of the angular 2 pt correlation function out to $z=1.6$. For redshifts $0.4<z<1.6$ we find that the amplitude of the angular correlation function is best parameterized by a comoving \\rn=2.37 \\Mpc\\ and $\\epsilon = -0.4^{+0.37}_{-0.65}$.  To match, however, the canonical local value for the clustering length, \\rn=5.4 \\Mpc, requires $\\epsilon = 2.1^{+0.4}_{-0.6}$, significantly more than simple linear growth.  It must be noted that while these results are in good agreement with those from published photometric and spectroscopic surveys (Le F\\`{e}vre et al 1996, Hudon and Lilly 1996) there are two caveats that should be considered before applying them to constrain models of structure formation. The small angular extent of the HDF (at a redshift, $z=1$, the field-of-view of the HDF is approximately 0.9 \\Mpc) means that fluctuations on scales larger than we probe will contribute to the variance of the measured clustering (Szapudi and Colombi 1996). Secondly, the requirement that the HDF be positioned such that it avoids bright galaxies ($I_{814}<20$) biases our clustering statistics by artificially suppressing the number of low redshift galaxies (a bias that will be present in most deep photometric surveys). Therefore, the clustering evolution in the HDF may not necessarily be representative of the general field population. Given this, there is enormous potential for the application of this technique to systematic wide angle multicolor surveys, such as the Sloan Digital Sky Survey,"
    },
    "9803/astro-ph9803271_arXiv.txt": {
        "abstract": "We argue that due to various restrictions cosmic strings and monopole-string networks are not likely to produce the observed flux of ultra-high energy cosmic rays (UHECR).  Among the topological defects studied so far, the most promising UHECR sources are necklaces and monopolonia. Other viable sources which are similar to topological defects are relic superheavy particles. All these sources have an excess of pions (and thus photons) over nucleons at production. We demonstrate that in the case of necklaces the diffuse proton flux can be larger than photon flux, due to absorption of the latter on radiobackground, while monopolonia and relic particles are concentrated in the Galactic halo, and the photon flux dominates. Another signature of the latter sources is anisotropy imposed by asymmetric position of the sun in the Galactic halo. In all cases considered so far, including necklaces, photons must be present in ultra-high energy radiation observed from topological defects, and experimental discrimination between photon-induced and proton-induced extensive air showers can give a clue to the origin of ultra-high energy cosmic rays. ",
        "introduction": "The observation of cosmic ray particles with energies higher than $10^{11}~GeV$ \\cite{EHE}  gives a serious challenge to the known mechanisms of acceleration. The shock acceleration in various astrophysical objects typically gives maximal energy of accelerated protons less than $(1-3)\\cdot 10^{10}~GeV$ \\cite{NMA} (see however \\cite{Bier97}). The unipolar induction can provide the maximum energy $1\\cdot 10^{11}~GeV$ only for extreme values of the parameters \\cite{BBDGP}. Much attention has recently been given to acceleration by ultrarelativistic shocks \\cite{Vie},\\cite{Wax}. The particles here can gain a tremendous increase in energy, equal to $\\Gamma^2$, at a single reflection, where $\\Gamma$ is the Lorentz factor of the shock. However, it is known (see e.g.  the simulation for pulsar relativistic wind in \\cite{Hosh}) that particles entering the shock region are captured there or at least have a small probability to escape.  Topological defects (TD) (for a review see \\cite{Book}) can naturally produce particles of ultrahigh energies (UHE) well in excess of those observed in cosmic rays (CR). In most cases the problem with topological defects is not the maximum energy but the fluxes. However, in some cases the predicted fluxes are comparable with observations.  Usually, UHE particles appear at the decays of superheavy (SH) particles produced by TD. (We shall refer to these SH particles as $X$-particles).  Examples discussed in the literature include ejection of X-particles from superconducting strings, emission of X-particles from cusps or intersections of ``ordinary'' strings, and production of such particles in monopole-antimonopole annihilations. Metastable SH particles can also be relics of an earlier epoch, produced by a thermal or some other mechanism in the early Universe.  A rather exceptional mechanism of UHE particle production is given by radiation of accelerated monopoles connected by strings. In  this case a monopole can radiate gluons with very large Lorentz factors and with virtualities of the order of the monopole acceleration.  A common signature of all extragalactic UHECR is the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{GZK}. It reveals itself as a steepening of the spectrum of UHE protons and nuclei due to their interaction with microwave radiation. The steepening starts at $E \\approx 3\\cdot 10^{10}~GeV$. Apart from this steepening there is another signature of interaction of extragalactic CR with microwave radiation: a  bump in the spectrum preceeding the cutoff. The bump is a consequence of the proton number conservation in the spectrum: protons loose energy and are accumulated before the cutoff.  In this paper we will discuss the signatures of UHECR from TD distinguishing them from particles produced by astrophysical accelerators.  We will confine ourselves here to the case of the conventional primary particles, protons and photons, and will not consider the other UHE signal carriers discussed in the literature such as neutrinos \\cite{bz,bordes}, Lightest Supersymmetric Particles \\cite{cfk,BeKa}, relativistic monopoles \\cite{wk} and vortons \\cite{bona}.  Throughout the paper we shall use the following numerical values and abbreviations: the dimensionless Hubble constant $h=0.65$, the Cold Dark Matter density in terms of critical density $\\Omega_{CDM}=0.2 h^2$, the size of Dark Matter halo $R_h=100~kpc$, UHECR - for Ultra High Energy Cosmic Rays and UHE - for Ultra High Energy, TD - for Topological Defect, SH - for Superheavy, CDM - for Cold Dark Matter, SUSY - for Supersymmetry, LLA - for Leading Logarithmic Approximation, AGN - for Active Galactic Nucleus, GC and AC - for Galactic Center and Anticenter, respectively. ",
        "conclusions": "We studied observational constraints on various TD models, as well as possible signatures of TD as sources of the observed UHE radiation ($E\\geq 10^{10}~GeV$).  The most stringent constraint is due to electromagnetic cascades. It depends on the energy spectrum of particles from decays and on astrophysical quantities which determine the development of the cascade (most notably on the flux of intergalactic infra-red/optical radiation and on intergalactic magnetic field ). The SUSY-QCD spectrum makes the cascade constraint weaker, because this spectrum predicts less higher energy particles and more low energy particles as compared with ordinary QCD spectrum. In case of very large $m_X$ it means that for a given UHECR flux, less energy is transferred to the e-m cascade radiation. There are considerable uncertainties in the extragalactic flux of infra-red radiation and extragalactic magnetic field. The conservative limit on the energy density of the cascade radiation imposed the latest EGRET data is $\\omega_{cas} \\approx 2\\cdot 10^{-6}~eV/cm^3$. It could be $3 - 5\\cdot 10^{-6}~eV/cm^3$ with astrophysical uncertainties mentioned above. The further progress in the study of origin of EGRET extragalactic flux and in calculation of SUSY-QCD spectrum, in the pessimistic case, can exclude such TD as e.g. necklaces as the sources of observed UHECR.     Another important constraint arises from the fact that at ultra-high energies, the proton attenuation length $R_p(E)$ and the photon absorption length $R_{\\gamma}(E)$ are both small compared to the Hubble radius.  Models in which the typical distance between defects is $D>>R_p$ are disfavored.  In such models, the observed spectrum would have an exponential cutoff, unless a source is accidentally close to the observer.  In the latter case the flux would be strongly anisotropic.  Finally, in many cases TD give UHECR fluxes lower than the observed ones. We showed here that this is the case for monopole-string networks. Superconducting and ordinary cosmic strings probably belong to this category as well, although some loopholes still remain to be closed.     With all these constraints taken into account, it appears that only necklaces, monopolonium and relic SH particles survive as potential UHE sources.  The most important observational signature of TD as sources of UHE CR is the presence of photon-induced EAS. For all known mechanisms of UHE particle production the pions (and thus photons) dominate over nucleons. At energies lower than $1\\cdot 10^{12}~GeV$, protons have considerably larger attenuation length than photons and the observed proton flux can be dominant. Nevertheless, even in this case photons reach an observer from sources located inside  the sphere of radius $R_\\gamma (E)$ (assuming that $R_\\gamma >D$). Unlike protons, photons propagate rectilinearly, indicating the direction to the sources.  {\\em Necklaces} with a large value of $r=m/\\mu d > 10^7$ have a small separation $D <R_\\gamma$. They are characterized by a small fraction, $R_{\\gamma}/R_p$, of photon-induced EAS at energies $10^{10} - 10^{11}~GeV$. This fraction increases with energy and becomes considerable at the highest energies. For smaller values of $r \\sim 10^4 - 10^6$, when the separation is larger than $R_{\\gamma}$ but still smaller than $R_p$, most of UHE particles are expected to be protons (with a chance of incidental proximity of a source, seen as a direct gamma-ray source). Thus, in all cases necklaces are characterized by an excess of proton-induced showers.  However, some fraction of photon-induced showers is always present, and it can be large at the highest energies.  {\\em Monopolonium}, decaying vortons {\\em and SH relic particles} are characterized  by an enhanced density in the Galactic halo. They give a photon-dominated flux without a GZK cutoff. Due to asymmetric position of the Earth in the Galaxy, this flux is anisotropic. The largest flux is expected from the direction of the Galactic Center, where the density of sources is the largest. Unfortunately, the Galactic Center is not seen by gigantic arrays , such as Akeno, Fly's Eye, Haverah Park, and Yakutsk array. However, these detectors can observe a minimum in the direction of the Galactic anticenter in particles with energies $E>1\\cdot 10^{10}~GeV$, as compared with the direction perpendicular to the Galactic Plane.  A flux from the Virgo cluster might be another signature of this model.  The search for photon induced showers is not an easy experimental task. It is known (see e.g. Ref.\\cite{AK}) that in the UHE photon-induced showers the muon content is very similar to that in proton-induced showers. However, some  difference in the muon content between these two cases is expected and may be used to distinguish between them observationally.  A detailed analysis would be needed to determine this difference.  The Landau-Pomeranchuk-Migdal (LPM) effect \\cite{LPM} and the absorption of photons in the geomagnetic field are two other important phenomena which affect the detection of UHE photons \\cite{AK,Kasa}; (see \\cite{ps} for a recent discussion). The LPM effect reduces the cross-sections of electromagnetic interactions at very high energies. However, if the primary photon approaches the Earth in a direction characterized by a large perpendicular component of the geomagnetic field, the photon is likely to decay into electron and positron \\cite{AK,Kasa}. Each of them emits a synchrotron photon, and as a result a bunch of photons strikes the Earth atmosphere. The LPM effect, which strongly depends on energy, is thus suppressed. If on the other hand a photon moves along the magnetic field, it does not decay, and LPM effect makes shower development in the atmosphere very slow. At extremely high energies the maximum of the showers can be so close to the Earth surface that it becomes \"unobservable\" \\cite{ps}.  We suggest that for all energies above the GZK cutoff the showers be analyzed as candidates for being induced by UHE photons, with the probability of photon splitting in the geomagnetic field determined form the observed direction of propagation, and with the LPM effect taken into account. The search for photon-induced showers can be especially effective in the case of Fly's Eye detector which can measure the longitudinal development of EAS. The future Auger detector will have, probably, the highest potentiality to resolve this problem."
    },
    "9803/astro-ph9803101_arXiv.txt": {
        "abstract": "We searched for cluster X-ray luminosity and radius evolution using our sample of 201 galaxy clusters detected in the 160~deg$^2$ survey with the \\ROSAT\\/ PSPC (Vikhlinin et al.\\ 1998). With such a large area survey, it is possible, for the first time with \\ROSAT\\/, to test the evolution of luminous clusters, $L_x>3\\times10^{44}\\,$\\ergpersec\\ in the 0.5--2~keV band. We detect a factor of 3--4 deficit of such luminous clusters at $z>0.3$ compared to the present.  The evolution is much weaker or absent at modestly lower luminosities, 1--$3\\times10^{44}\\,$\\ergpersec. At still lower luminosities, we find no evolution from the analysis of the $\\log N - \\log S$ relation. The results in the two upper $L_x$ bins are in agreement with the {\\em Einstein\\/} EMSS evolution result (Gioia et al.\\ 1990a, Henry et al.\\ 1992) while being obtained using a completely independent cluster sample.  The low-$L_x$ results are in agreement with other \\ROSAT\\/ surveys (e.g.\\ Rosati et al.\\ 1998, Jones et al.\\ 1998).  We also compare the distribution of core radii of nearby and distant ($z>0.4$) luminous (with equivalent temperatures 4--7~keV) clusters, and detect no evolution. The ratio of average core radius for $z\\sim0.5$ and $z<0.1$ clusters is $0.9\\pm0.1$, and the core radius distributions are remarkably similar. A decrease of cluster sizes incompatible with our data is predicted by self-similar evolution models for high-$\\Omega$ universe. ",
        "introduction": "The cluster evolution rate is a strong test of cosmological parameters (e.g., White \\& Rees 1978, Kaiser 1986, Eke, Cole \\& Frenk 1996). It is best to study evolution using X-ray selected samples of distant clusters which are much less affected by projection than the optically selected samples (van Haarlem et al.\\ 1997).  Of all the interesting cluster parameters such as mass, velocity dispersion, and temperature, the X-ray luminosity is the most accessible to measurements with present-day instruments, and most of the earlier studies were focused on evolution of the cluster X-ray luminosity function.  A strong evolution of cluster luminosities at $z\\sim0.1$ was reported from the EXOSAT survey (Edge et al.\\ 1990), but was later disproved by the \\ROSAT\\/ All-Sky Survey (Ebeling et al.\\ 1997). At higher redshifts, negative evolution of the cluster X-ray luminosity function was first reported by Gioia et al.\\ (1990a) using the {\\em Einstein}\\/ Extended Medium Sensitivity Survey (EMSS; Gioia et al.\\ 1990b, Stocke et al.\\ 1991). Gioia et al.\\ and later Henry et al.\\ (1992) compared the cluster luminosity functions below and above $z=0.3$.  They found that while the number of the low luminosity clusters does not evolve, there is a significant deficit of luminous, $L_x(0.3-3.5\\mbox{~keV})>5\\times10^{44}\\,$\\ergpersec, clusters at high redshift.  This EMSS result was questioned recently. Nichol et al.\\ (1997) reanalyzed the EMSS cluster sample using \\ROSAT\\/ X-ray and new optical observations and argued that the evolution reported in the original EMSS papers was not significant.  Several groups pursued independent searches for distant clusters in archival \\ROSAT\\/ PSPC observations. Collins et al.\\ (1997) found that the redshift distribution of 35 clusters detected in their 17~deg$^2$ survey is consistent with no evolution.  This contradicted the earlier claim by Castander et al.\\ (1995) of a strong evolution in a similar sample; however, the latter authors used an X-ray source detection algorithm not optimized for the cluster search.  Jones et al.\\ (1998) presented the $\\log N - \\log S$ relation for 46 clusters from their 16~deg$^2$ survey and found that this relation is consistent with no evolution of the $L_x<2\\times10^{44}\\,$\\ergpersec\\ (0.5--2~keV band) clusters.  Rosati et al.\\ (1998) derived cluster luminosity functions up to $z\\sim 0.8$ from their sample of 70 clusters detected in a 33~deg$^2$ survey, and found no evolution at low luminosities, $L_x<3\\times10^{44}\\,$\\ergpersec. However, none of these \\ROSAT\\/ surveys covers an area large enough to probe the evolution of the luminous clusters, and their no-evolution claims do not contradict the EMSS results.   Our 160~deg$^2$ survey (Vikhlinin et al.\\ 1998, hereafter Paper~I) is the first \\ROSAT\\/ survey comparable with the EMSS in sky coverage for distant clusters. We are able to test, and confirm, the EMSS evolution results even with the incomplete redshift data currently at hand. Eventually, when the spectroscopic work is complete, we will be able to characterize the luminosity evolution more accurately. In this Letter, we also show that the cluster X-ray core radii do not evolve between $z\\sim0.5$ and now. Throughout the paper, we use definitions $f_{-14}$ and $L_{44}$ for flux and luminosity in the 0.5--2~keV energy band in units of $10^{-14}\\,$\\ergs\\ and $10^{44}\\,$\\ergpersec, respectively. We also use $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0=0.5$. ",
        "conclusions": "We present a first \\ROSAT\\/ analysis of the evolution of luminous, $L_x>3\\times10^{44}\\,$\\ergpersec\\ distant clusters. We find a significant, factor of 3--4, decrease in the number of such clusters at $z>0.3$, confirming the detection of evolution in the EMSS (Gioia et al.\\ 1990a, Henry et al.\\ 1992). At lower luminosities, 1--$3\\times10^{44}\\,$\\ergpersec, the evolution is undetectable, with a decrease in number by a factor of only $1.3\\pm0.2$.  This is also consistent with the EMSS and other \\ROSAT\\/ surveys (e.g.\\ Rosati et al.\\ 1998). The absence of evolution of low luminosity clusters is also supported by the analysis of the $\\log N - \\log S$ distribution from which we find that the cluster volume emissivity, dominated by low-luminosity objects, does not evolve. The observed evolution can be reproduced by a model in which the characteristic luminosity decreases with redshift, but the comoving number density of clusters increases. Such models arise naturally in the hierarchical cluster formation scenario (e.g.\\ Kaiser 1986).  We compare the distribution of core-radii of distant, $z>0.4$ and nearby clusters.  We find that the distribution of core radii at $z>0.4$ is very similar to that in nearby clusters; the average radius has changed at $z>0.4$ by a factor of only $0.9\\pm0.1$. A stronger change is expected for hierarchical cluster formation in a flat universe (Kaiser 1986, Cen \\& Ostriker 1994). We also note that the assumption of no evolution of cluster sizes has been essentially used in flux measurements and area calculations in several X-ray surveys (e.g.\\ EMSS, Nichol et al. 1997), but is only verified here for the first time.  \\bigskip \\centerline{\\includegraphics[width=3.25in]{axdistr.ps}} \\vskip -15pt \\figcaption{Core-radius distribution for distant, $z>0.4$, clusters, derived from our survey (solid and dashed histogram for $q_0=0.5$ and $q_0=0$, respectively). The shaded histogram shows the core-radius distribution for nearby luminous clusters from Jones \\& Forman (1998). The angular resolution limit of our survey (15\\arcsec) corresponds to 120~kpc at the median redshift of distant clusters, well below the peak of the distribution. \\label{fig:axdistr}} \\medskip"
    },
    "9803/astro-ph9803337_arXiv.txt": {
        "abstract": "If the theoretical relationship between white dwarf mass and orbital period for wide-orbit binary radio pulsars is assumed to be correct, then the neutron star mass of PSR J2019+2425 is shown to be $\\sim 1.20 M_{\\odot}$. Hence the mass of the neutron star in this system prior to the mass transfer phase is expected to have been $< 1.1 M_{\\odot}$. Alternatively this system descends from the accretion induced collapse (AIC) of a massive white dwarf.\\\\ We estimate the magnetic inclination angles of all the observed wide-orbit low-mass binary pulsars in the Galactic disk using the core-mass period relation and assuming that the spin axis of an accreting neutron star aligns with the orbital angular momentum vector in the recycling process of the pulsar. The large estimated magnetic inclination angle of PSR J2019+2425, in combination with its old age, gives for this system evidence against alignment of the magnetic field axis with the rotational spin axis. However, in the majority of the similar systems the distribution of magnetic inclination angles is concentrated toward low values (if the core-mass period relation is correct) and suggests that alignment has taken place. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803294_arXiv.txt": {
        "abstract": "s{ Samples of high-redshift galaxies are easy to select in the millimetre/submillimetre (mm/submm) waveband using sensitive telescopes, because their flux density--redshift relations are expected to be flat, and so the selection function is almost redshift-independent at redshifts greater than 0.5. Source counts are expected to be very steep in the mm/submm waveband, and so the magnification bias due to gravitational lensing is expected to be very large, both for lensing by field galaxies and for lensing by clusters. Recent submm-wave observations of lensed images in clusters have constrained the submm-wave counts directly for the first time. In the next ten years our knowledge of galaxy evolution in this waveband will be greatly enhanced by the commissioning of sensitive new instruments and telescopes, including the CMBR imaging space mission {\\em Planck Surveyor}. This paper highlights the important features of gravitational lensing in the submm waveband and discusses the excellent prospects for lens searches using these forthcoming facilities. } ",
        "introduction": " ",
        "conclusions": "\\begin{enumerate} \\item The surface density of distant galaxies in the submm waveband, and therefore the expected surface density of gravitational lenses and the effects of source confusion in future observations, is now known with reasonable accuracy. \\item The {\\em Planck Surveyor} survey and surveys using other forthcoming mm/submm-wave telescopes will produce catalogues of distant sources that will be of great interest for studies of galaxy evolution. Lensed galaxies and quasars will be detected with an efficiency of up to 10\\% in these surveys, and a sample of order 1000 lenses could be compiled. \\end{enumerate}"
    },
    "9803/hep-ph9803418_arXiv.txt": {
        "abstract": "The neutrino capture rate measured by the Russian-American Gallium Experiment is well below that predicted by solar models. To check the response of this experiment to low-energy neutrinos, a 517 kCi source of \\nuc{51}{Cr} was produced by irradiating 512.7~g of 92.4\\%-enriched \\nuc{50}{Cr} in a high-flux fast neutron reactor.  This source, which mainly emits monoenergetic 747-keV neutrinos, was placed at the center of a 13.1 tonne target of liquid gallium and the cross section for the production of \\nuc{71}{Ge} by the inverse beta decay reaction $\\mnuc{71}{Ga}(\\nu_e,e^-)\\mnuc{71}{Ge}$ was measured to be {[5.55 \\+/-~0.60~(stat) \\+/-~0.32~(syst)]} \\E{-45} cm$^2$.  The ratio of this cross section to the theoretical cross section of Bahcall for this reaction is 0.95 \\+/- 0.12 (expt) $^{+0.035}_{-0.027}$ (theor) and to the cross section of Haxton is 0.87 \\+/- 0.11 (expt) \\+/- 0.09 (theor).  This good agreement between prediction and observation implies that the overall experimental efficiency is correctly determined and provides considerable evidence for the reliability of the solar neutrino measurement. ",
        "introduction": "Gallium experiments are uniquely able to measure the principal component of the solar neutrino spectrum.  This is because the low threshold of 233 keV \\cite{Audi and Wapstra 95} for inverse beta decay on the 40\\% abundant isotope \\nuc{71}{Ga} is well below the end point energy of the neutrinos from proton-proton fusion, which are predicted by standard solar models to be about 90\\% of the total flux.  The Russian-American Gallium Experiment (SAGE) has been measuring the capture rate of solar neutrinos with a target of gallium metal in the liquid state since January 1990.  The measured capture rate \\cite{Abdurashitov et al. 94,Gavrin 98} is 67~\\+/-7~(stat)~$^{+5}_{-6}$ (syst) SNU\\footnote{1 SNU corresponds to one neutrino capture per second in a target that contains 10$^{36}$ atoms of the neutrino absorbing isotope.}, a value that is well below solar model predictions of 137 $^{+8}_{-7}$ SNU \\cite{Bahcall and Pinsonneault and Wasserburg 95} and 125~\\+/- 5 SNU \\cite{Turck-Chieze and Lopes 93}.  In addition, the GALLEX Collaboration, which has been measuring the solar neutrino capture rate with an aqueous GaCl$_3$ target since 1991, observes a rate of 70~\\+/- 7~$^{+4}_{-5}$ SNU \\cite{Hampel et al. 96}.  The other two operating solar neutrino experiments, the chlorine experiment \\cite{Cleveland et al. 98} and the Kamiokande experiment \\cite{Suzuki et al. 95}, have significantly higher-energy thresholds, and thus are not able to see the neutrinos from $pp$ fusion.  When the results of these four solar neutrino experiments are considered together, a contradiction arises which cannot be accommodated by current solar models, but which can be explained if one assumes that neutrinos can transform from one species to another \\cite{Bahcall 94,Berezinsky 94,Parke 95,Hata et al. 94,Castellani et al. 94,Bahcall et al. 95,Heeger and Robertson 96}.  The gallium experiment, in common with other radiochemical solar neutrino experiments, relies on the ability to extract, purify, and count, all with well known efficiencies, a few atoms of a radioactive element that were produced by neutrino interactions inside many tonnes of the target material.  In the case of 60 tonnes of Ga, this represents the removal of a few tens of atoms of \\nuc{71}{Ge} from $5 \\times 10^{29}$ atoms of Ga.  To measure the efficiency of extraction, about 700 $\\mu$g of stable Ge carrier is added to the Ga at the beginning of each exposure, but even after this addition, the separation factor of Ge from Ga is still 1 atom in 10$^{11}$.  This impressively stringent requirement raises legitimate questions about how well the many efficiencies that are factored into the final result are known.  It has been understood since the outset that a rigorous check of the entire operation of the detector (i.e., the chemical extraction efficiency, the counting efficiency, and the analysis technique) would be made if it is exposed to a known flux of low-energy neutrinos.  In addition to verifying the operation of the detector, such a test also eliminates any significant concerns regarding the possibility that atoms of \\nuc{71}{Ge} produced by inverse beta decay may be chemically bound to the gallium (so-called ``hot atom chemistry'') in a manner that yields a different extraction efficiency than that of the natural Ge carrier.  In other words, it tests a fundamental assumption in radiochemical experiments that the extraction efficiency of atoms produced by neutrino interactions is the same as that of carrier atoms.  This article describes such a test, in which a portion of the SAGE gallium target was exposed to a known flux of \\nuc{51}{Cr} neutrinos and the production rate of \\nuc{71}{Ge} was measured.  Similar tests have also been made by GALLEX \\cite{Hampel et al. 98}.  Although a direct test with a well-characterized neutrino source lends significant credibility to the radiochemical technique, we note that numerous investigations have been undertaken during the SAGE experiment to ensure that the various efficiencies are as quoted \\cite{Abdurashitov et al. 94}.  The extraction efficiency has been determined by a variety of chemical and volumetric measurements that rely on the introduction and subsequent extraction of a known amount of the stable Ge carrier.  A test was also carried out in which Ge carrier doped with a known number of \\nuc{71}{Ge} atoms was added to 7 tonnes of Ga.  Three standard extractions were performed, and it was demonstrated that the extraction efficiencies of the carrier and \\nuc{71}{Ge} follow each other very closely.  Another experiment was performed to specifically test the possibility that atomic excitations might tie up \\nuc{71}{Ge} in a chemical form from which it would not be efficiently extracted. There is a concern that this might occur in liquid gallium because the metastable Ga$_2$ molecule exists with a binding energy of \\about1.6 eV.  In this experiment the radioactive isotopes \\nuc{70}{Ga} and \\nuc{72}{Ga} were produced in liquid gallium by neutron irradiation.  These isotopes quickly beta decay to \\nuc{70}{Ge} and \\nuc{72}{Ge}.  The Ge isotopes were extracted from the Ga using our standard procedure and their number was measured by mass spectrometry.  The results were consistent with the number expected to be produced based on the known neutron flux and capture cross section, thus suggesting that chemical traps are not present.  This experiment is not conclusive, however, because the maximal energy imparted to the \\nuc{70}{Ge} nucleus following beta decay of \\nuc{70}{Ga} is 32 eV, somewhat higher than the maximal energy of 20 eV received by the \\nuc{71}{Ge} nucleus following capture of a 747-keV neutrino from \\nuc{51}{Cr} decay (and considerably higher than the maximum nuclear recoil energy of 6.1 eV after capture of a 420-keV neutrino from proton-proton fusion).  Further evidence that the extraction efficiency was well understood came from monitoring the initial removal from the Ga of cosmogenically produced \\nuc{68}{Ge}.  This nuclide was generated in the Ga as it resided on the surface exposed to cosmic rays.  When the Ga was brought underground, the reduction in the \\nuc{68}{Ge} content in the initial extractions was the same as for the Ge carrier.  These numerous checks and auxiliary measurements have been a source of confidence in our methodology, yet it is clear that a test with an artificial neutrino source of known activity provides the most compelling validation of radiochemical procedures.  This article is an elaboration of work that previously appeared in Ref.~\\cite{Abdurashitov et al. 96}.  The experimental changes since the previous Letter are some minor refinements in the selection of candidate \\nuc{71}{Ge} events and in the treatment of systematic errors; recent cross section calculations are also included.  The central experimental result given here is almost identical to what was reported earlier. ",
        "conclusions": "The primary motivation for the \\nuc{51}{Cr} source experiment was to determine if there is any unexpected problem in either the chemistry of extraction or the counting of \\nuc{71}{Ge}, i.e., to see if there is some unknown systematic error in one or both of the efficiency factors in $\\epsilon$, the product of extraction and counting efficiencies.  If some such systematic error were to exist, then the value of $\\epsilon$ that we have used in the preceding will be in error by the factor $E$, defined as $E \\equiv \\epsilon_{\\text{true}}/ \\epsilon_{\\text{measured}}$.  Since the cross section is inversely proportional to $\\epsilon$, this hypothetical error is equivalent to the cross section ratio, $E = \\sigma_{\\text{measured}}/\\sigma_{\\text{true}}$.  An experimental value for $E$ can be set from our measured cross section, Eq.~(\\ref{cross section result}), if one assumes that the true cross section is equivalent to the theoretically calculated cross section.  Then $E \\approx R \\equiv \\sigma_{\\text{measured}}/\\sigma_{\\text{theoretical}}$.  Neutrino capture cross sections averaged over the four neutrino lines of \\nuc{51}{Cr} have been calculated by Bahcall \\cite{Bahcall 97} and by Haxton \\cite{Haxton 98}.  Bahcall, assuming that the strength of the two excited states in \\nuc{71}{Ge} that can be reached by \\nuc{51}{Cr} neutrinos is accurately determined by forward-angle $(p,n)$ scattering, gives a result of 5.81 (1.0 $^{+0.036}_{-0.028})$ \\E{-45} cm$^2$.  The upper limit for the uncertainty was set by assuming that the excited state strength could be in error by as much as a factor of 2; minor contributions to the uncertainty arise from forbidden corrections, the \\nuc{71}{Ge} lifetime, and the threshold energy.  An independent consideration of the contribution of excited states has been made by Hata and Haxton \\cite{Hata and Haxton 95} and very recently by Haxton \\cite{Haxton 98}.  They argue that, because of destructive interference between weak spin and strong spin-tensor amplitudes in \\nuc{71}{Ge}, the strengths determined from $(p,n)$ reactions are, for some nuclear levels, poor guides to the true weak interaction strength.  In particular, Haxton finds the weak interaction strength of the $(5/2)^-$ level in \\nuc{71}{Ge} at an excitation energy of 175 keV to be much greater than the value that is measured by the $(p,n)$ scattering reaction, and calculates a total \\nuc{51}{Cr} cross section of (6.39~\\+/- 0.68) \\E{-45} cm$^2$.  This cross section was deduced from the measured $(p,n)$ cross sections for the two excited states, and uses a large-basis shell model calculation to correct for the presence of spin-tensor contributions.  Since not all known theoretical uncertainties were included, the stated error here is a lower bound.  Combining our statistical and systematic uncertainties for the cross section in quadrature into an experimental uncertainty, we can thus give estimates for $E$: \\begin{eqnarray} E & \\approx & R \\equiv \\frac{\\sigma_{\\text{measured}}} {\\sigma_{\\text{theoretical}}} \\\\ & = & \\left\\{ \\begin{array}{l l} 0.95 \\pm 0.12 \\text{ (expt)}\\ ^{+0.035}_{-0.027} \\text{ (theor)} & \\text{ (Bahcall)}, \\\\ 0.87 \\pm 0.11 \\text{ (expt)}\\ \\pm 0.09 \\text{ (theor)} & \\text{ (Haxton)}. \\end{array} \\right. \\nonumber \\end{eqnarray} \\noindent With either of these theoretical cross sections, $R$ is consistent with unity, which implies that the total efficiency of the SAGE experiment to the neutrinos from \\nuc{51}{Cr} is close to 100\\%.  The measurement reported here should not be interpreted as a direct calibration of the SAGE detector for solar neutrinos. This is because the \\nuc{51}{Cr} neutrino spectrum differs from the solar spectrum, there is a 10\\%--15\\% uncertainty in the theoretical value for the \\nuc{51}{Cr} cross section, and the total experimental efficiency for each solar neutrino measurement is known to a higher precision than the 12\\% experimental uncertainty obtained with the \\nuc{51}{Cr} source.  As a result, the solar neutrino measurements reported by SAGE should not be scaled by the factor $E$.  Rather, we consider the Cr experiment as a test of the experimental procedures, and conclude that it has demonstrated {\\it with neutrinos} that there is no unknown systematic uncertainty at the 10\\%--15\\% level.  The neutrino spectrum from \\nuc{51}{Cr} is very similar to that of \\nuc{7}{Be}, but at slightly lower energy.  Since the response of \\nuc{71}{Ga} to \\nuc{7}{Be} neutrinos is governed by the same transitions that are involved in the \\nuc{51}{Cr} source experiment, we can definitely claim that, if the interaction strength derived from the \\nuc{51}{Cr} experiment is used in the analysis of the solar neutrino results, then the capture rate measured by SAGE includes the full contribution of neutrinos from \\nuc{7}{Be}.  This observation holds independent of the value of $E$ or of cross section uncertainties.  This demonstration is of considerable importance because a large suppression of the \\nuc{7}{Be} neutrino flux from the sun is one consequence of the combined analysis of the four operating solar neutrino experiments \\cite{Bludman et al. 93,Akhmedov et al. 95}.  GALLEX has completed two \\nuc{51}{Cr} measurements whose combined result, using the cross section of Bahcall \\cite{Bahcall 97}, can be expressed as $R$ = 0.93~\\+/- 0.08 \\cite{Hampel et al. 98}, where the uncertainty in the theoretical cross section has been neglected.  Both SAGE and GALLEX, which employ very different chemistries, give similar results for the solar neutrino capture rate and have tested their efficiencies with neutrino source experiments.  The solar neutrino capture rate measured in Ga is in striking disagreement with standard solar model predictions and there is considerable evidence that this disagreement is not an experimental artifact."
    },
    "9803/astro-ph9803011_arXiv.txt": {
        "abstract": "We describe narrowband and spectroscopic searches for emission-line star forming galaxies in the redshift range 3 -- 6 with the 10 m Keck\\,II Telescope. These searches yield a substantial population of objects with only a single strong (equivalent width $\\gg 100$ \\AA) emission line, lying in the $4000 - 8500$ \\AA\\ range.  Spectra of the objects found in narrowband--selected samples at $\\lambda \\sim5390$ \\AA\\ and $\\sim6741$ \\AA\\ show that these very high equivalent width emission lines are generally redshifted Ly$\\alpha\\ \\lambda\\,1216$ \\AA\\ at $z\\sim3.4$ and 4.5.  The density of these emitters above the 5$\\sigma$ detection limit of $1.5\\times 10^{-17}$ ergs cm$^{-2}$ s$^{-1}$ is roughly 15,000/$\\sq^{\\circ}$/unit $z$ at both $z\\sim3.4$ and 4.5.  A complementary deeper ($1\\ \\sigma \\sim 10^{-18}$ ergs cm$^{-2}$ s$^{-1}$) slit spectroscopic search covering a wide redshift range but a more limited spatial area ($200\\ \\sq''$) shows such objects can be found over the redshift range $z=3 - 6$, with the currently highest redshift detected being at $z=5.64$.  The Ly$\\alpha$ flux distribution can be used to estimate a minimum star formation rate in the absence of reddening of roughly $0.01\\ M_{\\odot}$ Mpc$^{-3}$ yr$^{-1}$ ($H_0 = 65\\ {\\rm km}\\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1}$, $q_0 = 0.5$).  Corrections for reddening are likely to be no larger than a factor of two, since observed equivalent widths are close to the maximum values obtainable from ionization by a massive star population. Within the still significant uncertainties, the star formation rate from the Ly$\\alpha$--selected sample is comparable to that of the color-break--selected samples at $z\\sim 3$, but may represent an increasing fraction of the total rates at higher redshifts.  This higher-$z$ population can be readily studied with large ground-based telescopes. ",
        "introduction": "The search for high-redshift galaxies and the effort to map the star formation history of galaxies have progressed rapidly in the last several years as magnitude--limited spectroscopic surveys pushed into the $z=1-3$ range (Cowie et al.\\markcite{large_sample} 1996; Cohen et al.\\markcite{cohenhdf} 1996), while color--based selection techniques produced many objects in the $z=2.5-5$ range (Steidel et al.\\markcite{stei96a} 1996a, \\markcite{stei96b}1996b; Franx et al.\\markcite{franx} 1997; Steidel et al.\\markcite{stei98} 1998), particularly in the exquisite Hubble Deep Field (HDF) sample (Lowenthal et al.\\markcite{low97} 1997).  However, the galaxies chosen by these techniques correspond to objects with ongoing massive star formation and small amounts of extinction, and may represent only part of the populations at these early epochs.  More evolved objects may be heavily dust--reddened and more easily picked out at submillimeter wavelengths (Smail, Ivison, \\& Blaine\\markcite{smail97} 1997; Hauser et al.\\markcite{dirbe} 1998), while earlier stages in evolution may have relatively little continuum light and be too faint to be seen in the magnitude--limited samples or selected with the color-break techniques at current sensitivity limits.  This latter class of objects may represent the earliest stages of the galaxy formation process, in which substantial amounts of metals have yet to form.  These galaxies may have much stronger Ly$\\alpha$ emission relative to the stellar continuum, since they have massive star formation that can excite the Ly$\\alpha$ emission line, but without so much dust that the line is suppressed, and this can result in very high observed equivalent widths in the range of 100--200\\,(1+$z$) \\AA\\ (e.g., Charlot \\& Fall\\markcite{charl93} 1993).  Such objects may be hard to pick out with color-break techniques but be detectable in Ly$\\alpha$ searches of sufficient depth (Cowie\\markcite{cowie88} 1988; Thommes\\markcite{thommes} 1996).  An increased incidence of strong Ly$\\alpha$ emission does, indeed, appear in the color-break samples at fainter continuum magnitudes (Steidel et al.\\markcite{stei98} 1998).  Earlier blank-field Ly$\\alpha$ surveys (e.g., Thompson, Djorgovski, \\& Trauger\\markcite{tdt95} 1995; Thompson \\& Djorgovski\\markcite{td95} 1995) failed to find such objects, as Pritchet\\markcite{pri94} (1994) has summarized, but their sensitivity lay at the margin of where such objects would be expected in significant numbers (Cowie\\markcite{cowie88} 1988; Thommes \\& Meisenheimer\\markcite{thommes95} 1995).  However, Hu \\& McMahon\\markcite{br2237} (1996), using very deep targeted narrowband searches ($1\\ \\sigma = 1.5 \\times 10^{-17}$ ergs cm$^{-2}$ s$^{-1}$), found $z\\sim4.55$ Ly$\\alpha$-emitting galaxies with the very strong emission and weak or undetected continuua predicted for early star-forming objects. The advent of 10 m telescopes, along with this successful detection of Ly$\\alpha$ emitters, prompted us to undertake a new survey that has been successful in detecting blank-field high-$z$ Ly$\\alpha$ emitters. The present Letter describes the early results of this search, which uses extremely deep narrowband filter exposures taken with LRIS (Oke et al.\\markcite{lris} 1995) on the Keck II telescope to search for emission-line populations at extremely faint flux levels ($1\\ \\sigma=3 \\times 10^{-18}$ ergs cm$^{-2}$ s$^{-1}$).  This survey picks out sources of extremely high equivalent width ($W_{\\lambda}>100$ \\AA) emission lines as candidates, and then uses followup LRIS spectroscopy (\\S2) to determine if these are Ly$\\alpha$ emitters.  The first results from the Hawaii survey with a 5390/77 \\AA\\ filter, corresponding to Ly$\\alpha$ emission at $z\\sim3.4$, were described in Cowie \\& Hu\\markcite{smitty1} 1998 (hereafter, Paper I), and yielded a number of candidate high-$z$ galaxies similar to the redshift $z\\sim4.55$ Ly$\\alpha$-emitting galaxies found by Hu \\& McMahon\\markcite{br2237} (1996) and emitters at $z\\sim 2.4$ (Pascarelle et al.\\markcite{pasc96} 1996; Francis\\markcite{fra97} et al.\\ 1997) in targeted searches.  In Paper I, we showed that the use of color selection on emission-line selected objects of high equivalent width ($> 100$ \\AA) picked out objects with continuum colors similar to those of color--selected Lyman break galaxies with measured redshifts of $z\\sim3.4$.  They also recovered the one field object whose previously measured redshift placed Ly$\\alpha$ within the filter bandpass.  However, the emission-line galaxies selected by their high equivalent width also comprised objects with very faint continuua, that would have fallen below the magnitude threshold of current Lyman break surveys, and also included two objects that could not be detected in Keck imaging of the optical continuum ($1\\ \\sigma$ $B=27.8$, $V=27.5$, and $I=25.8$; $W_{\\lambda}>400$ \\AA).  In the present Letter we first present spectroscopic followups for the narrowband candidates of Paper I.  In a small fraction of the cases, the spectrum shows both Ly$\\alpha$ and \\civ\\ $\\lambda$ 1550 \\AA, confirming the redshift identification but suggesting AGN-like properties.  However, the majority of the spectra show only a single strong emission line.  The absence of other detectable features, in combination with the high equivalent width of the selected candidates, identifes the single line as redshifted Ly$\\alpha$ emission, and argues that the equivalent width criterion ($W_{\\lambda}\\gg 100$ \\AA) can, in fact, be used as a good diagnostic of high-$z$ Ly$\\alpha$-emitting galaxies.  We then (\\S3) present results of a second deep narrowband search with a 78 \\AA\\ bandpass $\\lambda \\sim6741$ \\AA\\ filter (Ly$\\alpha$ at $z\\sim 4.54$) and followup spectroscopy, that confirmed two Ly$\\alpha$-emitting galaxies at this higher redshift. In \\S4, we describe a very deep blank-field spectroscopic search (6 hr LRIS integration covering $\\lambda\\lambda\\,\\sim5000-10000$ \\AA) that yielded four emitters at redshifts 3.04 -- 5.64.  Finally (\\S5), the data on emission-line objects from the imaging surveys in the two redshift intervals are combined with various spectroscopic surveys at lower redshift to show the evolution of the emission-line fluxes with redshift. We argue that the Ly$\\alpha$-emitting objects are significant contributors to the integrated star formation of the galaxy population throughout the $z=3 - 6$ redshift range, and that the integrated star formation rate of the Ly$\\alpha$ selected objects is flat, or possibly rising through this redshift range, with a value greater than 0.01 $M_{\\odot}$ Mpc$^{-3}$ for $q_0 = 0.5$ and $H_0 = 65\\ {\\rm km}\\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1}$.  At the highest redshifts, most of the star formation may be occurring in objects of this class. ",
        "conclusions": "Since resonant scattering enhances the effects of extinction, it is harder to convert the Ly$\\alpha$ emission into a massive star formation rate than it is for line luminosity diagnostics such as H$\\alpha$ and \\oii\\ 3727.  For the present calculation, we assume that extinction may be neglected in computing the required massive star formation rates, which then constitute a minimum estimate. However, upward corrections to this value are unlikely to be larger than a factor of two, since the observed rest-frame equivalent widths lie in the $100-200$ \\AA\\ range --- close to the maximum values that are obtainable from ionization by a massive star population (Charlot \\& Fall\\markcite{charl93} 1993).  Then, assuming case B recombination, we have $L($Ly$\\alpha$) = 8.7\\,$L($H$\\alpha$) (Brocklehurst\\markcite{brock} 1971), which using Kennicutt's\\markcite{kenn83} (1983) translation of $\\dot{M}$ from H$\\alpha$ luminosity, gives $\\dot{M}=(L($Ly$\\alpha$)/$10^{42}$ ergs s$^{-1}$) $M_{\\odot}$ yr$^{-1}$.  In order to cross-calibrate to \\oii\\ fluxes at lower redshift we assume $f$(H$\\alpha$+\\nii) = $1.25\\,f$(\\oii) based on the mean values of the ratio in both the Gallego et al.\\markcite{gal95} (1995) and Hawaii Deep Survey (Cowie et al.\\markcite{large_sample} 1996) samples.  We also assume $f$(H$\\alpha$+\\nii) = $1.33\\,f$(H$\\alpha$) (Kennicutt\\markcite{kenn83} 1983).  In Figure~\\ref{fig:5} we compare the range of Ly$\\alpha$ luminosities to the range of line luminosities in lower redshift objects.  The plot shows the quantity $z^2\\,f$ vs redshift, where H$\\alpha$+\\nii, \\oii, and Ly$\\alpha$ fluxes have been converted to H$\\alpha$ fluxes using the relationships above, and we have restricted ourselves to the imaging data in which the fluxes of the Ly$\\alpha$ are well determined.  The comparison objects are drawn from the surveys of Gallego et al.\\markcite{gal_95} 1995 ({\\it filled boxes}), Songaila et al.\\markcite{ksurvey_3} 1994 ({\\it pluses}), and the Hawaii $B=25$ sample [{\\it open boxes\\/} (H$\\alpha$+\\nii), {\\it triangles\\/} (\\oii), and {\\it circles\\/} (Ly$\\alpha$)]. The solid ($q_0=0.5$) and dashed ($q_0=0.02$) lines on Fig.~\\ref{fig:5} show the fluxes corresponding to stellar mass production rates of 10 $M_{\\odot}$ yr$^{-1}$, 1 $M_{\\odot}$ yr$^{-1}$, and 0.1 $M_{\\odot}$ yr$^{-1}$ for $H_0 = 65\\ {\\rm km}\\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1}$.  Maximum formation rates locally are around a few $M_{\\odot}$ yr$^{-1}$, rising to values of just over 10 $M_{\\odot}$ yr$^{-1}$ above $z=0.6$.  The Ly$\\alpha$ fluxes at the higher redshifts are then consistent with this value or just slightly smaller depending on the extinction correction.  The CADIS results (Thommes et al.\\markcite{cadis} 1998) would lie a factor of several times higher in flux than the two Hawaii filter samples, but both Keck spectroscopy and repeat Fabry-P\\'erot observations have disproved the original $z\\sim5.7$ candidate selection (Meisenheimer\\markcite{meisen} 1998).  The minimum integrated star formation rates at $z=3.4$ and $z=4.5$ are 0.006 $M_{\\odot}$ Mpc$^{-3}$ yr$^{-1}$ and 0.01 $M_{\\odot}$ Mpc$^{-3}$ yr$^{-1}$ respectively for $q_0=0.5$ where the first value is slightly smaller than that given in Paper I since AGN-like objects are excluded.  Both values are lower limits calculated in the absence of extinction and the $z=4.5$ value is based on a single field. Within the substantial uncertainties of the as yet small samples, the results suggest that the star formation rates in the strong emission line population are constant or may possibly be increasing with redshift from $z=3 - 6$, in contrast to color-based samples where the rate is declining at higher redshifts (Madau et al.\\markcite{madau96} 1996, \\markcite{madau98} 1998).  This is consistent with the broad expectation that as we move to higher redshifts and earlier stages of galaxy formation, an increasingly larger fraction of the star formation should be in strong Ly$\\alpha$ emitters that correspond to the youngest galaxies."
    },
    "9803/astro-ph9803227_arXiv.txt": {
        "abstract": "Black holes are by definition {\\it black}, and therefore cannot be directly observed by using electromagnetic radiations. Convincing identification of black holes must necessarily depend on the identification of a very specially behaving matter and radiation which surround them. A major problem in this subject of black hole astrophysics is to quantify the behaviour of matter and radiation close to the horizon. In this review, the subject of black hole accretion and outflow is systematically developed. It is shown that both the stationary as well as the non-stationary properties of the observed spectra could be generally understood by these solutions. It is suggested that the solutions of radiative hydrodynamic equations may produce clear spectral signatures of black holes. Other circumstantial evidences of black holes, both in the galactic centers as well as in binary systems, are also presented. ",
        "introduction": "Stellar mass black holes are the end products of stars. After the fuel is exhausted inside a normal star, the core collapses and the supernova explosion occurs. If the mass of the core is lower than, say, $\\sim 3M_\\odot$, the object formed at the center may be a neutron star. Otherwise, it is a black hole. Therefore, some of the compact binary systems should contain black holes. Similarly, core collapse in the proto-galactic phase could also produce supermassive black holes ($M \\sim 10^6$ to $10^9M_\\odot$). In spiral galaxies, the central black holes are less massive (say, $10^{6-7}\\ M_\\odot$), while in elliptical galaxies the central black holes are more massive (say, a few times $10^{8-9}\\ M_\\odot$).  Astrophysical community generally believes that the black holes should exist because of the solid foundation of the theory of general relativity which predicts them. The problem remains that of identification. Black holes do not emit anything except Hawking radiation, which, for any typical mass of the astrophysical black holes is so cold (typically $60$ nano Kelvin for a solar mass black hole, and goes down inversely with increase in mass) that it would be entirely masked by the much hotter microwave background radiation. Classically, black holes are point-like with infinite density and are surrounded by an imaginary one-way membrane called `event horizon' of radius $R_g=2GM_{BH}/c^2$. Here, $G$ and $c$ are gravitational constant and velocity of light respectively, $M_{BH}$ is the mass of the black hole. $R_g$ is known as the Schwarzschild radius and is roughly equal to $30$ kilometers for a $10\\ M_\\odot$ black hole. For a maximally rotating (Kerr) black hole, the radius is half as small. Surrounding matter and radiation are pulled by the black hole only to disappear inside never to be observed again. Not even light, what to talk about matter, can escape to distant observers from regions within the horizon, making it {\\it impossible} to detect a black hole through direct observations. A positive identification must therefore rely on indirect and circumstantial evidences. In fact, the problem of identification of black holes boils down to the identification of surrounding matter which may behave in a `funny' way. We shall quantify the degree of `funniness' as we go along.  In this {\\it review}, we discuss how a black hole could be identified. We first present elementary properties of the spacetime around a black hole and compare them with those of a Newtonian star. We discuss in great length the properties of the global solutions of equations which govern the behaviour of matter. We then show that the observations in the last couple of decades do agree with these properties. Towards the end we make a comparative study of methodologies of black hole detection and present our judgment on the best way to detect black holes. ",
        "conclusions": "That black holes, which represent the end product of massive stars and star clusters, must exist somewhere in this universe is beyond any doubt. The issues discussed in this review were: whether they are in principle detectable, how to detect them and whether they have been detected. It seems that a few cases at least they {\\it have been detected}. If the observations of Genzel et al. [96, 125] is correct, then the mass of the central $0.1pc$ region of our galaxy would be $2.5-3.2 \\times 10^6 M_\\odot$ and the corresponding mass density would be $6.5 \\times 10^9 M_\\odot /pc^3$, the highest measured concentration so far. The water mega-maser measurement of the nucleus of NGC4258 within $0.1pc$ has the central mass of $4 \\times 10^7 M_\\odot$ and corresponding mass density is $6.5 \\times 10^9 M_\\odot /pc^3$. The central mass of M87 from the estimation of Keplerian and non-Keplerian components is $ \\sim 4 \\times 10^9 M_\\odot$ and the corresponding mass density is $2.0 \\times 10^7 M_\\odot /pc^3$. Although, Cyg X-1 is the most studied black hole candidate so far, its mass function is very low. Its confirmation as a black hole comes from its spectral features, especially the weak power-law slope of the bulk motion Comptonization in its soft state. The only candidates with mass function higher than, say, $3 M_\\odot$, are $GRS1124-683$, $GRO J1655-40$, $H 1705-250$, $GS 2000+25$ and $GS2023+338$ and are possible stellar mass black holes. With the improvements of the future observational techniques, one needs to focus on more detailed predictions of the advective disks, such as variation of the solution topology with specific energy, or equivalently, accretion rate. With the emergence of gravitational wave astronomy, the wave signals from galactic centers should be detectable. The proposal presented in Section (5.3) would for the first time correlate the distortions of the gravitational wave signals with those from the spectral signatures. Together they would not only verify black holes, they may also become the strongest test of general relativity to date.  \\newpage  \\centerline {Reference}"
    },
    "9803/astro-ph9803157_arXiv.txt": {
        "abstract": "The effects which star cluster concentration and binarity have on observable parameters, that characterise the dynamical state of a population of stars after their birth aggregate dissolves, are investigated. To this end, the correlations between ejection velocity, binary proportion, mean system mass, binary orbital period and mass ratio are quantified for simulated aggregates. These consist of a few hundred~low-mass binary and single stars, and have half-mass radii in the range~2.5 to 0.08~pc.  The primordial binary-star population has a period distribution similar to that observed in Taurus-Auriga for pre-main sequence binaries.  The findings presented here are useful for interpreting correlations between relative locations and proper motions, binary properties and masses of young stellar systems within and surrounding star forming regions, and of stellar systems escaping from Galactic clusters.  For the low-concentration binary-rich aggregates, the proportion of binaries decreases monotonically as a function of increasing ejection velocity after aggregate dissolution, as expected. However, this is not the case for initially highly concentrated binary-rich aggregates. The reason for this difference is the interplay between the disruption of binary systems and the initial depth of the potential well from which the stellar systems escape.  After aggregate dissolution, a slowly expanding remnant population remains. It can have a high binary proportion (80~per cent) with a high mean system mass, or a low binary proportion (less than about 20~per cent) with a low mean system mass, if it was born in a low- or a high-concentration aggregate, respectively.  It follows that adjacent regions on the sky near some star-forming clouds can have young populations with different binary proportions and different mass functions, even if the binary proportion at birth and the initial mass function (IMF) were the same.  Binary systems that are ejected from the aggregate tend to be massive, and their mass ratio tends to be biased towards higher values.  The mean system mass is approximately independent of ejection velocity between~2 and~30~km/s. Dynamical ejection from binary-rich aggregates adds, within~10~Myr, relatively massive systems to regions as far as~300~pc from active star-forming centres.  Long-period systems cannot survive accelerations to high velocities. The present experiments show that a long-period ($>10^4$~d) binary system with a large velocity ($>30$~km/s) cannot be ejected from an aggregate.  If such young systems exist, then they will have been born in high-velocity clouds. ",
        "introduction": "\\label{sec:intro} \\noindent Stellar systems (i.e. single or multiple stars) form in groups. The dynamical processes within these alter the properties of the young systems when they leave the site where they formed.  The dynamical properties of a stellar system are its mass (i.e. luminosity if age is known), the multiplicity and the orbital parameters if it is a multiple system.  The distribution of velocities of young stars emanating from star-forming centres (i.e. the kinematical signature of star formation) will also be affected by the dynamical interactions within the young groups. Both, the distribution of {\\it dynamical properties} and the {\\it kinematical signature of star formation} bear an imprint of the dynamical configuration at the time when the stellar group was born.  Star formation in Taurus-Auriga gave birth to aggregates with sizes of roughly 0.5--1~pc consisting of about 20--50 stars.  It is now well established that most stars form in binary systems in Taurus-Auriga (e.g. K\\\"ohler \\& Leinert 1998).  The same appears to hold true in other star-forming regions (Ghez et al. 1997).  Embedded clusters may also have a binary proportion that is higher than in the Galactic field (Padgett, Strom \\& Ghez 1997).  In the Trapezium cluster, which is a very dense embedded cluster and probably less than 1~Myr old, Prosser et al. (1994) find a binary proportion that is at least as large as in the Galactic field. In the cluster core, Petr et al. (1998) observe, for low-mass stars, a binary proportion similar to the Galactic field, and smaller by about a factor of three than the binary proportion in Taurus-Auriga.  These findings are particularly interesting, because binary destruction is expected to be efficient in such an environment.  A review of pre-main sequence binary stars, and their relation to Galactic field systems, is provided by Mathieu (1994, see also Kroupa 1995a, Simon et al. 1995).  Young Galactic clusters also contain binary systems. The particularly well studied Pleiades and Praesepe clusters have binary proportions of 40--50~per cent, for systems of spectral type earlier than K0 (Raboud \\& Mermilliod 1998a, 1998b).  There exists thus evidence that the formation of binary systems may be by far the dominant star-formation mode in both loose groups and highly concentrated embedded clusters, some of which may evolve to bound Galactic clusters. The term {\\it aggregates} is used henceforth to mean loose groups or embedded clusters of more than 10~stars.  If stars form predominantly in aggregates of binary systems, then the kinetic energy distribution after aggregate dissolution should be enhanced at high energies, when compared to dissolved aggregates of single stars, because binary star binding energy can transform into kinetic energy (Heggie 1975, Hills 1975, Hut 1983).  Large accelerations are destructive to binary systems, so that the proportion of binaries should be a decreasing function of increasing final kinetic energy.  Additionally, different initial aggregate concentrations lead to different final binary proportions and kinematical signatures, as will be shown here.  This is also true under the extreme assumption that {\\it all} stars always form in binary systems with the same initial dynamical properties. This dynamical mechanism of producing variations in binary proportion and associated dynamical properties stands in contrast to a possible variation of these parameters determined by the star-formation process. Durisen \\& Sterzik (1994) make the interesting point that the binary proportion may be smaller in molecular clouds with a higher temperature than in lower-temperature clouds.  It is important to study the signatures that arise from purely dynamical interactions in stellar groups, for a comparison with outcomes from usually less well-understood alternative scenarios.  In a study of the large-scale distribution of young stars around active star-forming regions, Sterzik \\& Durisen (1995) find that the dynamical decay of small stellar groups can lead to sufficiently large velocities to populate large areas on the sky with young stars, so that these need not have formed near their observed location. They find that special initial dynamical configurations of the stars (e.g. cold thin strings) lead to enhanced production of ejected stars. Initial decay of cold sub-groups within larger complexes also has this effect (Aarseth \\& Hills 1972), and scattering of proto-stars on cloud clumps during an even earlier dynamical phase may likewise eject very young low-mass stars (Gorti \\& Bhatt 1996).  However, the number of ejected stars cannot account for the observed number of widely distributed young stars (Feigelson 1996).  The evolution of circum-stellar discs around stars ejected from small stellar groups is studied by Armitage \\& Clarke (1997), and McDonald \\& Clarke (1995) show that the presence of circum-stellar material in small proto-stellar groups increases the number of binaries formed and randomises the mass-ratio distribution.  That the number of dynamically ejected stars is increased significantly in binary-rich stellar aggregates, when compared to clusters consisting initially only of single stars, is shown by Kroupa (1995c).  These simulations show that a mass-ratio distribution produced by randomly associating masses from the IMF, decays to the observed distribution for G-dwarf binaries, if most stars form in aggregates similar to observed embedded clusters.  Also, initially more concentrated aggregates produce more stars with a high ejection velocity, the maximum of which increases with decreasing cluster radius.  De la Fuente Marcos (1997) investigates the dependence on cluster richness, and finds that the mean ejection velocity increases for more initially populous clusters. Ejection velocities larger than a few hundred~km/s can be achieved in young star clusters containing massive primordial binaries (Leonard \\& Duncan 1990). This may explain the location of OB stars far from active star-forming sites. Leonard (1991) finds, on the basis of many scattering experiments, that the maximum ejection velocity is of the order of the escape velocity from the stellar surface of the most massive star. If its mass is $60\\,M_\\odot$, then a similar star can attain an ejection velocity of up to 700~km/s. A low-mass star may find itself fleeing with a velocity of up to 1400~km/s, after a surface-grazing encounter with such a star. A critical discussion of the possible origin of runaway OB stars is provided by Leonard (1995). He stresses that collisions of two stars during binary-binary interactions can produce runaway OB stars with very similar properties as in the alternative scenario, in which such stars result from a supernova explosion in close binary systems. An interesting and insightful discussion of the implications of the binary properties of runaway OB stars for the dynamical configuration of massive stars at birth is to be found in Clarke \\& Pringle (1992).  In this paper, the correlations between stellar velocity, system mass and binary proportion that arise from aggregates with different initial concentration and consisting initially either of 400~single stars or of 200~binary systems, is studied.  The resulting correlations are useful for interpreting the properties and distribution of young stars near and in star forming regions (see for example Brandner et al. 1996, Feigelson 1996, Frink et al. 1997).  In Section~\\ref{sec:method} the assumptions, simulations and definitions are described.  The results are presented in Section~\\ref{sec:results}, and Section~\\ref{sec:conclusions} contains the conclusions. ",
        "conclusions": "\\label{sec:conclusions} \\noindent The correlations between ejection velocity and the proportion of binaries, as well as their orbital parameters, have been quantified for a range of initial dynamical configurations.  The correlations are useful in the study of stellar systems that are apparently ejected from Galactic clusters (see e.g. Frink et al. 1997), some of which are known to be rich in binaries (e.g. Raboud \\& Mermilliod 1998a, 1998b). Observed ejected binaries should show correlations as presented in Figs.~\\ref{fig:orbit1}, \\ref{fig:orbit2}, ~\\ref{fig:mass1} and~\\ref{fig:mass2}.  The results for the binary-rich aggregates modelled here are also relevant for an understanding of the large-scale distribution of young stars, because most stars appear to form in aggregates with a high binary proportion.  Additionally, the correlations contain information about the dynamical configuration at birth.  For binary-rich aggregates containing a few hundred stars the following correlations result: (i) more tightly clustered aggregates lead to more stellar systems having larger ejection velocities and a smaller overall binary proportion, (ii) the large population of primordial binaries leads to a significantly enhanced number of systems with high-ejection velocities compared to single-star aggregates, (iii) systems with high ejection velocities have a significantly reduced binary proportion, (iv) binary stars with high ejection velocities have short-period orbits, and tend to be more massive with a mass-ratio biased towards unity, (v) the average system mass as a function of ejection velocity is defined above about 2~km/s by stochastic close encounters, so that systems more massive than $0.5\\,M_\\odot$ with high ejection velocities occur, and (vi) aggregates with $R_{0.5}\\le0.25$~pc lead to a complex dependence of the resulting binary proportion on velocity, whereas a stellar population emerging from less concentrated aggregates shows a monotonic decrease of the binary proportion with increasing velocity.  For aggregates of a few hundred single stars one obtains: (i) more tightly clustered aggregates lead to an increased number of stellar systems with larger ejection velocities (but significantly less so than in the binary-rich aggregates), and an enhanced overall binary proportion that remains significantly below the observed binary proportion in the Galactic field, (ii) the binary proportion increases with ejection velocity, (iii) is as (iv) above, and (iv) is as (v) above.  Remnant unbound young populations take long to disperse because they have a small velocity dispersion. The binary proportion and mean system mass (and thus the inferred IMF) of such a remnant population, sensitively depends on the initial dynamical configuration of the binary-rich birth aggregate. After emerging from the birth aggregate, the distribution of velocities of a young stellar population changes with time in the gravitational potential of the nearby molecular clouds.  A substantial proportion of emerging stars is likely to remain bound to the parent molecular cloud until it ceases to exist.  These findings are important for interpreting the spatial distribution, kinematics and binarity of young stars within and surrounding star-forming regions.  Molecular clouds, in which stars form preferentially in dense embedded binary-rich clusters, should have an enhanced halo population of ejected and relatively binary poor ($f\\approx0.25$) young stellar systems.  Also, young but binary-depleted groups of stars can be misinterpreted to be evidence for an environmental dependency of the binary-formation mechanism. For example, in fig.~6 of Brandner et al. (1996), the region US-B has more binaries than the region US-A, which also contains many more B~stars than US-B. The presence of B~stars suggests that the stars in US-A may have formed in dense embedded clusters. The stars seen in US-A would then constitute the $v\\simless1$~km/s remnant population for $R_{0.5}\\simless0.25$~pc (Fig.~\\ref{fig:binprop1}).  Given the results of the present study, it is suggested that such a difference in binary proportion between two regions may be due to different initial dynamical configurations, and need not imply a dependence of the binary proportion on the star-forming environment.  Important for the interpretation of the large-scale distribution of young stars surrounding star forming sites is the realisation that relatively massive systems are ejected with relatively large velocity (2--30~km/s, Fig.~\\ref{fig:binprop1}), which is a point also stressed by Sterzik \\& Durisen (1995).  The X-ray surveys are flux limited and detect the massive stars (Wichmann et al. 1996), the presence of which around star-forming regions may be a natural consequence of the processes studied here.  However, if some young binary systems are found to have orbital periods that place them above the dashed lines in the right panels of Figs.~\\ref{fig:orbit1}--\\ref{fig:orbit3}, then this would support the suggestion by Feigelson (1996), that some star-formation occurs in small high-velocity clouds."
    },
    "9803/astro-ph9803196.txt": {
        "abstract": "} \\nc{\\eab}{ \\noindent The explanation of the observed galactic magnetic fields may require the existence of a primordial magnetic field. Such a field may arise during the early cosmological phase transitions, or because of other particle physics related phenomena in the very early universe reviewed here. The turbulent evolution of the initial, randomly fluctuating microscopic field to a large-scale macroscopic  field can be described in terms of a shell model, which provides an approximation to the complete magnetohydrodynamics. The results indicate that there is an inverse cascade of magnetic energy whereby the coherence of the magnetic field is increased by many orders of magnitude. Cosmological seed fields roughly of the order of $10^{-20}$ G at the scale of protogalaxy, as required by the dynamo explanation of galactic magnetic fields, thus seem plausible. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Apart from the baryon number and the spectrum of energy density fluctuations, the physical processes that took place in the very early universe do not have many consequences that could  still be directly detectable today. Most observables have been washed away by the thermal bath of the pre-recombination era. One possibility, which has recently received increased attention, is offered by the large-scale magnetic fields observed in a number of galaxies, in galactic halos, and in clusters of galaxies  \\cite{observe,becketal}. The astrophysical mechanism responsible for the origin of the galactic magnetic fields is not understood. Usually one postulates a small seed field, which can then be either enhanced by the compression of the protogalaxy, and/or exponentially amplified by the turbulent fluid motion as in the dynamo theory  \\cite{dynamo}. The exciting possibility is that the seed field could be truly primordial \\cite{kulsrud}, in which case cosmic magnetic fields could provide direct information about the very early universe.  Early magnetic fields could then play an important role in particle cosmology by modifying the dispersion or clustering properties of various particles.  One particular example is the fate of the neutrino: because of their magnetic moments, Dirac neutrinos propagating in the background of a magnetic field would be subject to a  spin flip \\cite{eers}, so that a left-handed neutrino can be turned into a right-handed neutrino, giving rise to an extra effective neutrino degree of freedom and thereby affecting primordial nucleosynthesis. Dark matter particles could also be sensitive to the presence of a magnetic field. For instance, axions couple to magnetic fields, but perhaps surprisingly, it can be shown that despite the coupling, cold axion oscillations are not much affected by the presence of a primordial magnetic field \\cite{jarkkoax}.  The issue at hand is then: is it possible that primordial magnetic fields of significant strength exist? To answer this,  first one has to find a mechanism in the early  universe which is able to produce a large enough magnetic field.  There are various proposals, a number of which are based on the early cosmological phase transitions, which are discussed in Sect 3. The second problem is to explain how the initial field, which is expected to be random as it is created  by microphysics and having correlation lengths typical to microphysics, can grow up to be coherent enough at large length scales. This is a problem in magnetohydrodynamics which is discussed in Sect 4. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Explaining the galactic magnetic fields in terms of microphysical processes that took place when the universe was only ten billionth of a second old is a daunting task, which is not made easier by the complicated evolution of the magnetic field as it is twisted and tangled by the flow of plasma. It is nevertheless encouraging that mechanisms for generating primordial magnetic  fields of suitable size exist, and in particular those based on the early cosmological phase transitions discussed in Sect. 3 look promising. At the same time the fact that there are so many possibilities tends to underline our ignorance of the details of the subsequent evolution of the magnetic field. The step from microphysics to macroscopic fields is a difficult one because of the very large magnetic Reynolds number of the early universe. However, different considerations, both analytic approximations, 2d simulations, as well as the full-fledged shell model computations which can account for turbulence, seem to point to the existence of an inverse cascade of magnetic energy. Moreover, as discussed in Sect. 4.3, the inverse cascade is obtained also in the presence of a large plasma viscosity. Therefore the primordial origin of the galactig magnetic fields is quite possible.  Much theoretical work remains to be done, though. At the same time it is very important that progress is made on the observational front. In particular, measuring or setting a stringent limit on the intergalactic field, which could be possible in the near future as indicated in Sect. 2.3, would provide the testing ground for all theoretical scenarios. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "9803/astro-ph9803143_arXiv.txt": {
        "abstract": "s{We describe two new -- {\\it stochastic-geometrical} -- methods to obtain reliable velocity field statistics from N-body simulations and from any general density and velocity fluctuation field sampled at a discrete set of locations. These methods, the {\\it Voronoi tessellation method} and {\\it Delaunay tessellation method}, are based on the use of the Voronoi and Delaunay tessellations of the point distribution defined by the locations at which the velocity field is sampled. Adjusting themselves automatically to the density of sampling points, they represent the optimal estimator for volume-averaged quantities. They are therefore particularly suited for checking the validity of the predictions of quasi-linear analytical density and velocity field perturbation theory through the results of N-body simulations of structure formation. We illustrate the subsequent succesfull application of the two methods to estimate the bias-independent value of $\\Omega$ in the N-body simulations on the basis of the predictions of perturbation theory for the $\\Omega$-dependence of the moments and PDF of the velocity divergence in gravitational instability structure formation scenarios with Gaussian initial conditions. We will also shortly discuss practical and complicating issues involved in the obvious extension of the Voronoi and Delaunay method to the analysis of observational samples of galaxy peculiar velocities.} \\vskip 1.0cm ",
        "introduction": "The study of the large-scale cosmic velocity field is a very promising and crucial area for the understanding of structure formation. The cosmic velocity field is in particularly interesting because of its close relation to the underlying field of mass fluctuations. Indeed, on these large and (quasi)-linear scales the acceleration, and therefore the velocity, of any object is expected to have an exclusively gravitational origin so that it should be independent of its nature, whether it concerns a dark matter particle or a bright galaxy. Moreover, in the linear regime the generic gravitational instability scenario of structure formation predicts that at every location in the Universe the local velocity is related to the local acceleration, and hence the local mass density fluctuation field, through the same universal function of the cosmic density parameter $\\Omega$ (Peebles 1980), $f(\\Omega) \\ \\propto \\ \\Omega^{0.6}$. Because linear theory provides a good description on scales exceeding a few Megaparsec, the use of this straightforward relation implies the possibility of a simple inversion of the measured velocity field into a field that is directly proportional to the field of local mass density fluctuations. Such a procedure can then be invoked to infer the value of $\\Omega$, through a comparison of the resulting field with the field of mass density fluctuations in the same region.  However, such a determination of $\\Omega$ may be contrived as the estimate of the mass density fluctuation field on the basis of the observed galaxy distribution may offer a biased view of the underlying mass distribution. By lack of a complete and self-consistent physical theory of galaxy formation, the commonly adopted approach is to make the simplifying assumption that the galaxy density $\\delta_g$ and the mass density $\\delta$ are related via a linear bias factor $b$, $\\delta_g\\,=\\,b\\,\\delta$. The comparison between the observed galaxy density fluctuation field and the local cosmic velocity field will therefore yield an estimate of the ratio $\\beta\\,=\\,{f(\\Omega)/b} \\, \\approx\\, {\\Omega^{0.6}/b}\\,.$. However, while numerous studies have yielded estimates of $\\beta$ in the range $\\beta \\approx 0.5-1.2$ (see Dekel 1994, Strauss \\& Willick 1995, for compilations of results), it has proven very cumbersome to subsequently disentangle the contribution of $\\Omega$ and $b$ to the quantity $\\beta$. In fact, it turns out to be impossible within procedures based on the linearity of the analysed velocity field. ",
        "conclusions": "The main incentive for developing the Delaunay and Voronoi method is provided by the wish to be able to infer a bias-independent value of $\\Omega$ through comparison of the velocity statistics obtained from the discrete point sample with those of analytical distributions. In figure 3 we show the PDFs of the velocity divergence $\\theta$ that were numerically determined by the Delaunay method for a range of N-body simulations, each with a different value of $\\Omega$. The solid curve shows the corresponding analytical distribution function $p(\\theta)$, for the cosmic epoch with the same dispersion $\\sigma_{\\theta}$. For contrast, each of the four frames also contains the dashed curve for the PDF in an Einstein-de Sitter universe with the same value of $\\sigma_{\\theta}$. Evidently, the Delaunay method is highly succesfull in reproducing the correct statistical distribution, and via the relations between the moments of the PDF we indeed obtain very good estimates of $\\Omega$.  While figure 3 illustrates the potential power of the tessellation methods, we are obviously motivated to apply them to more practical situations and hence more cumbersome cases where selection and sampling effects and sampling errors are of crucial influence. In particular we hope to be able to develop a formalism capable of dealing with observational catalogues of peculiar velocities of galaxies. In previous work (Bernardeau \\& Van de Weygaert 1996, Bernardeau et al. 1997) we already adressed the issue of diluted samples. In those cases both methods yielded encouraging results. However, the true world will present problems ranging from the fact that one can measure galaxy velocities only along the line of sight to complicated selection effects like differential Malmquist bias (see e.g. Bertschinger et al. 1990, Dekel, Bertschinger \\& Faber 1990). Work on these issues is in progress, but they obviously provide a considerable complication."
    },
    "9803/astro-ph9803233_arXiv.txt": {
        "abstract": "Parallax data from the Hipparcos mission allow the direct distance to open clusters to be compared with the distance inferred from main sequence (MS) fitting.  There are surprising differences between the two distance measurements, which indicate either the need for changes in the cluster compositions or reddening, underlying problems with the technique of main sequence fitting, or systematic errors in the Hipparcos parallaxes at the 1 mas level.  We examine the different possibilities, focusing on MS fitting in both metallicity-sensitive \\bv\\ and metallicity-insensitive $V-I$ for five well-studied systems (the Hyades, Pleiades, $\\alpha$ Per, Praesepe, and Coma Ber).  The Hipparcos distances to the Hyades and $\\alpha$ Per are within 1 $\\sigma$ of the MS fitting distance in \\bv\\ and $V-I$, while the Hipparcos distances to Coma Ber and the Pleiades are in disagreement with the MS fitting distance at more than the 3 $\\sigma$ level.  There are two Hipparcos measurements of the distance to Praesepe; one is in good agreement with the MS fitting distance and the other disagrees at the 2 $\\sigma$ level.  The distance estimates from the different colors are in conflict with one another for Coma but in agreement for the Pleiades.  Changes in the relative cluster metal abundances, age related effects, helium, and reddening are shown to be unlikely to explain the puzzling behavior of the Pleiades. We present evidence for spatially dependent systematic errors at the 1 mas level in the parallaxes of Pleiades stars.  The implications of this result are discussed. ",
        "introduction": "Main sequence fitting is a basic tool used in the study of star clusters; the principle behind it is also used to estimate distances to field main sequence (MS) stars.  The Hipparcos mission (ESA 1997) has provided parallaxes for a number of open cluster stars, which permits a direct determination of the distances to the open clusters which can be compared with distances obtained from MS fitting.  There are surprising differences between distances obtained with these two methods; in this paper we explore possible explanations for them.  MS fitting relies upon the Vogt-Russell theorem:  the location of a star in the HR diagram is uniquely specified by its mass, composition, and age. This implies that we can infer the distance of a given cluster by comparing the apparent magnitudes of cluster stars with the absolute magnitudes of stars with known composition and distance.  There are several possible approaches.  Unevolved lower MS field stars with known distances or a cluster (such as the Hyades) of known distance can be used to construct an empirical MS.  The distance to the cluster is inferred from the vertical shift needed to line up the cluster MS with the empirical MS.  Clusters can also be compared with theoretical isochrones calibrated on the Sun; the latter method requires a color calibration which relates the model effective temperatures to the observed colors.  Most nearby open clusters are close to the Sun in metal abundance, which minimizes uncertainties in the distance scale from variations in composition. There is also a large database of fundamental effective temperature measurements for stars near the solar [Fe/H], so the color calibrations should be relatively reliable.  The nearby open clusters also have been extensively studied for membership, photometry, abundances, and reddening. For all of these reasons the open cluster distance scale has not been regarded as controversial, and evidence that MS fitting yields incorrect distances could have significant astrophysical importance.  The Hipparcos mission has resulted in a large increase in the number of open cluster stars with measured parallaxes.  This data allows the distance scale inferred from MS fitting to be compared with the distance scale inferred from trigonometric parallaxes.  The recently announced Hipparcos determination of the mean parallax of the Pleiades cluster gives the result $8.61 \\pm 0.23$ milliarcsec \\markcite{vh97a} (van Leeuwen \\& Hansen Ruiz 1997a).  This corresponds to a distance of $116\\pm3$ pc, or a distance modulus of $5.32\\pm0.06$ magnitude.  Traditional determinations of the Pleiades distance (e.g., \\markcite{vb84} VandenBerg \\& Bridges 1984; \\markcite{s93} Soderblom et al. 1993), comparing the cluster's main sequence to that of nearby stars, lead to a distance modulus of about 5.6 mag ($d \\sim 130$ pc; $\\pi \\sim 7.7$ mas).  Thus the Hipparcos parallax, being almost 1 mas larger than expected, suggests that the Pleiades cluster stars are systematically $\\sim 0.3$ magnitude fainter than we have thought up to now.  Parallaxes for stars in other clusters have also been measured, and the results are compared with those obtained from MS fitting in Table 1 (data taken from \\markcite{phip97} Perryman et al. 1997, \\markcite{mhip97} Mermilliod et al. 1997, \\markcite{rhip97} Robichon et al. 1997).  The standard reddening for the clusters is also indicated, along with a notation about whether or not differential reddening is present.  The second column lists the cluster [Fe/H] values from Boesgaard \\& Friel (1990) and Friel \\& Boesgaard (1992); we have adopted their abundance scale for the clusters in the present study (see Section 4).  Mermilliod et al. 1997 and Robichon et al. 1997 concluded that there is no simple explanation for the discrepancies between the MS fitting and Hipparcos distances, and that all of the possible classes of solutions appeared unsatisfactory.  \\begin{deluxetable}{lcccccc} \\tablenum{1} \\tablecaption{Open Cluster Parameters} \\tablehead{ \\colhead{Cluster} & [Fe/H] & \\colhead{$m-M$} & \\colhead{$(m-M)_o$} & \\colhead{$(m-M)_o$} & \\colhead{$(m-M)_o$} & \\colhead{$E(B-V)$}\\\\ \\colhead{ } & & \\colhead{Apparent} & \\colhead{Lynga} & \\colhead{Hipparcos} & \\colhead{This paper} & \\colhead{mag}} \\startdata Hyades        & $+0.13$ & 3.01  & 3.01  & 3.33${\\pm}0.01$ & 3.34${\\pm}0.04$ & 0.00\\nl Coma Ber      & $-0.07$ & 4.49  & 4.49  & 4.73${\\pm}0.04$ & 4.54${\\pm}0.04$ & 0.00\\nl Pleiades      & $-0.03$  & 5.61  & 5.48  & 5.33${\\pm}0.06$ & 5.60${\\pm}0.04$ & 0.04\\nl IC 2602       &  & 6.02  & 5.89  & 5.84${\\pm}0.07$ & \\nodata & 0.04\\nl IC 2391       &  & 5.96  & 5.92  & 5.83${\\pm}0.08$ & \\nodata & 0.01\\nl Praesepe      & $+0.04$ & 5.99  & 5.99  & 6.24${\\pm}0.12$ & 6.16${\\pm}0.05$ & 0.00\\nl ${\\alpha}$ Per & $-0.05$ & 6.36  & 6.07  & 6.33${\\pm}0.09$ & 6.23${\\pm}0.06$ & 0.10\\tablenotemark{a}\\nl Blanco 1      & & 6.97  & 6.90  & 7.01${\\pm}0.26$ & \\nodata & 0.02\\nl IC 4756       &  & 8.58  & 7.94  & 7.30${\\pm}0.19$ & \\nodata & 0.20\\tablenotemark{a}\\nl NGC 6475      &  & 7.08  & 6.89  & 7.32${\\pm}0.19$ & \\nodata & 0.06\\nl NGC 6633      &  & 8.01  & 7.47  & 7.32${\\pm}0.34$ & \\nodata & 0.17\\tablenotemark{a}\\nl Stock 2       & & 8.62  & 7.41  & 7.50${\\pm}0.32$ & \\nodata & 0.38\\tablenotemark{a}\\nl NGC 2516      &  & 8.49  & 8.07  & 7.71${\\pm}0.15$ & \\nodata & 0.13\\nl NGC 3532      &  & 8.53  & 8.40  & 8.10${\\pm}0.36$ & \\nodata & 0.04\\nl \\enddata \\tablenotetext{a}{Variable reddening} \\end{deluxetable}  We note that a second calculation of the distance to Praesepe has been performed by \\markcite{vh97b} van Leeuwen \\& Hansen Ruiz (1997b), and they find a distance modulus of 6.49$\\pm$0.15 - in disagreement both with MS fitting and the Mermilliod et al. Hipparcos distance.  For the purposes of this paper we have adopted the Mermilliod distance; if we were to adopt the VH97b distance to the cluster we would have to add Praesepe to the list of clusters with a significant (2 $\\sigma$) discrepancy between the MS fitting and Hipparcos distance scales.  The first column of distance moduli in Table 1 lists the values cited as ``Lynga'' by Mermilliod et al. (1997) and Robichon et al. (1997).  We note that these are {\\it apparent} distance moduli, needing considerable (up to 1.2 mag) corrections for extinction, and cannot be directly compared with the distance moduli $(m-M)_o$ calculated from the Hipparcos parallaxes. The second column in Table 1 lists the distance moduli which correspond to the cluster distances given in Lynga's (1987) Catalogue.  These distances come from a variety of sources, are still scaled to a Hyades distance modulus of 3.01 mag, and need corrections for each clusters metallicity. One motivation for our study is to place MS fitting distances for open clusters on a consistent scale.  In a paper in preparation, we have found that the MS fitting distances to some of the more distant open clusters are substantially different from the Lynga distances and in marked disagreement with the Hipparcos parallax distances.  A second question is the precision of MS fitting estimates; we will show that accuracy at the 0.05 mag level is possible for well-studied systems.  Our results for the clusters studied in this paper are in the fourth column.  Discrepancies between the MS fitting distances and the Hipparcos distances could arise from several sources.  As indicated above, one possibility is that the MS fitting distances need to be rederived on a consistent scale.  Another possibility is that some of the basic properties of well-studied open clusters, such as composition, age, or reddening, need to be revised.  If neither of these possibilities can reconcile the distance scales, then we are left with one of two important conclusions : either there are fundamental problems with MS fitting or there are unrecognized systematic errors in the Hipparcos parallaxes themselves.  These issues are important for other questions as well.  For example, recent proposed revisions to the globular cluster distance and age scales, based on Hipparcos parallaxes of subdwarfs, rely on the same MS fitting technique that gives rise to the puzzling distances to open clusters (\\markcite{r97} Reid 1997; \\markcite{g97} Gratton et al. 1997; \\markcite{c98} Chaboyer et al. 1998; but see also \\markcite{pmtv98} Pont et al. 1998).  In this paper we address the essential issues raised above.  The Pleiades, Praesepe, and $\\alpha$ Per are well-suited for a more detailed examination.  There is good membership information and multicolor photometry for all three; $\\alpha$ Per is a system with an age comparable to that of the Pleiades (50 Myr vs. 100 Myr) and therefore it provides a test of age-related effects.  We have also examined the Coma Ber star cluster, which has a low quoted error for its Hipparcos distance.  In a companion paper (Soderblom et al. 1998) we have searched for field stars with accurate parallaxes and anomalous positions in the HR diagram.  We begin by describing the theoretical models which we use and the open cluster data in section 2.  In section 3 we begin with a comparison of the Pleiades, Praesepe, and $\\alpha$ Per in different colors.  We then use the Hyades cluster to test the zero-point of our distance scale, check on the shape of the isochrones in the observational color-magnitude diagram, and to determine the sensitivity of distance estimates in different colors to changes in metal abundance. We then derive distance modulus estimates at both solar [Fe/H] and the individual abundances inferred from high-resolution spectroscopy for the Pleiades, Coma Ber, Praesepe, and $\\alpha$ Per using several different methods and both \\bv\\ and $V-I$.  The Pleiades and Coma Ber are found to be in disagreement with the Hipparcos distance scale.  We discuss the sensitivity of our results to age, composition, and reddening in section 4, and present evidence that the Hipparcos parallaxes may contain small-scale ($\\sim$1 deg) systematic effects $\\sim$1 mas in size, large enough to cause the Pleiades parallax discrepancy.  Our conclusions are in section 5. ",
        "conclusions": "The results of Section 3 indicate that it is the Hipparcos distance to the Pleiades which is in the most serious conflict with MS fitting.  In all of the other systems except Coma Ber, MS fitting in different colors yields distance results that are consistent with one another, normal helium, and [Fe/H] values from high resolution spectroscopy.  Coma Ber may have an equally serious disagreement, but the unusual behavior of the cluster in $V-I$ suggests that other problems may be contributing to the discrepancy for it.  We therefore examine in turn the various possible mechanisms that could reconcile the cluster distance scales for the Pleiades; in all cases we believe that they cannot do so.  In a companion paper we show that the same conclusions result from an examination of nearby field stars \\markcite{s98} (Soderblom et al. 1998).  We then proceed to an analysis of the Hipparcos parallaxes for the Pleiades, and show that there are indications of possible systematic errors that could be the origin of the discrepancy.  The calculations that we have presented are standard stellar models. We have therefore not included physical processes such as gravitational settling, rotational mixing, magnetic fields, internal gravity waves, or mass loss, which are surely present.  There are strong reasons for believing that these nonstandard effects will not influence the distance scale, although they could be potentially important for other issues. The single most important reason is the youth of the clusters that we have examined; detailed nonstandard calculations predict little, if any, effect for ages as young as the Pleiades.  In addition, any such process would have to affect stars with a wide range in masses to a similar extent and be different among different clusters to explain the pattern that we see.  Gravitational settling is minimal in young systems such as the Pleiades, and the degree to which helium and heavy elements sink depends strongly on the convection zone depth and thus the stellar mass.  For example, helium and heavy element diffusion are a 10\\% fractional effect in the Sun, which is almost 50 times older than the Pleiades.  The observed cluster lithium abundances require a mild envelope mixing process, and models with rotational mixing that are consistent with the lithium data predict little or no deep mixing (Pinsonneault 1997).  In addition, the observed range in rotation rates in clusters is large, and any extra mixing would produce a spread in MS properties rather than a uniform shift in the distance estimates.  Other physical processes could affect the results, but they are still subject to a variety of observational constraints which make a large effect unlikely.  We have compared different standard model calculations, and the zero-point offset is small (0.01-0.03 mag for stars between 5600 and 7000 K, for example).  The systematic errors in the standard model distance estimates is therefore also too small to explain the results that we have obtained.  We now discuss age, composition, and reddening effects.  \\subsection{Age and Stellar Activity} It is well-known that many young stars are heavily spotted; this could influence the color-temperature relationship and therefore the distance estimates for young systems such as the Pleiades and $\\alpha$ Per.  In Figures 1 and 2 we compared these two clusters at the Hipparcos distances in our two colors; the Pleiades is clearly anomalous with respect to $\\alpha$ Per if the Hipparcos distance scale is adopted.  Since $\\alpha$ Per is younger and has a larger population of rapid rotators, if anything $\\alpha$ Per should be more anomalous than the Pleiades if our MS fitting age estimates were in error because of activity.  We note that similar conclusions can be obtained by comparing young and old field stars (Soderblom et al. 1998).  The narrow width of the Pleiades MS also indicates that a wide range in stellar activity does not produce a significant effect on the color-temperature relationship.  For all of these reasons we reject the idea that youth is responsible for the difference between the distance estimates.  Another possibility is that activity could be influencing the Pleiades [Fe/H], which has been derived from LTE model atmospheres. If such a phenomenon were at play, it might lead to derived abundances being a function of line strength due to the direct effect of activity on the stronger lines formed at smaller depths in the photosphere. We have a number of high resolution spectra of Pleiades members that was originally obtained to study lithium abundances.  We have analyzed the \\ion{Fe}{1} data in the cool Pleiades dwarfs and find no such [Fe/H]-line strength correlation.  This does not exclude such a real correlation, though, given the influence of damping, which is adjusted to enforce such a lack of correlation.  To the extent that our damping assumptions seem quite reasonable compared to numerous other fine spectroscopic analyses, and are consistently applied in both the stellar and solar analyses to yield line-by-line [Fe/H] values, the analysis suggests any such trends are not substantial.  Regardless, any systematic error in the inferred mean [Fe/H] is greatly mitigated by the fact that the damping adjustments enforce consistency with the weaker lines, which are formed at deeper depths, and thus presumably are more immune from the direct effects of chromospheric activity. Activity in very young stars can manifest itself in the form of an effective veiling continuum.  Such behavior would presumably weaken the line absorption, thus leading to {\\it underestimated} line strengths and, hence, abundances.  Detailed NLTE line formation calculations to determine how the active Pleiades dwarfs' Fe and other metal abundances might be affected by activity, spots, convective flows, {\\it etc.\\/} would be of interest, but are unlikely to produce large errors for the reasons discussed above.  \\subsection{Heavy Metals}  \\subsubsection{The Cluster [Fe/H] Scale} Homogeneous Fe abundances are available for the Pleiades, Praesepe, and $\\alpha$ Per from the work of Boesgaard and collaborators.  Independent modern fine analyses of these clusters (and a few others) by other investigators are available for comparison with their work.  All the studies considered here derive self-consistent solar Fe abundances with which the stellar values are normalized.  Such a careful differential procedure can greatly reduce errors introduced by varying assumptions concerning the solar Fe abundance, model atmospheres, $gf$ values, {\\it etc}.  \\markcite{bbr88} Boesgaard {\\it et al.\\/}~(1988) determine a mean Pleiades iron abundance of [Fe/H]$=-0.03$ from analysis of 17 F stars.  The mean star-by-star reddening they use is essentially identical to the value we have adopted.  \\markcite{b88}Boesgaard (1989) determined a ``best'' Pleiades abundance by analyzing new data for 8 Pleiads; the result was [Fe/H]$=+0.02$.  Boesgaard \\& Friel (1990) used new data for 12 of the same stars in Boesgaard {\\it et al.\\/} to find a mean [Fe/H]$=-0.03$.  The single datum standard deviation in all these studies is ${\\sim}0.07$ dex.  The 1${\\sigma}$ level error in the mean is 0.02-0.03 dex, so the internal statistical uncertainties appear to be small. \\markcite{ccc88}Cayrel {\\it et al.\\/}~(1988) derive a mean Pleiades [Fe/H] of $+0.13$ from analysis of four Pleiades dwarfs, three of which are significantly cooler (mid G) than the Boesgaard F stars.  The standard deviation is 0.10 dex, which is somewhat smaller than their estimated individual errors; the error in the mean is ${\\sim}0.06$ dex.  The ${\\sim}0.1$ dex offset between the Cayrel and Boesgaard values is representative of uncertainties in reddening (which enters via photometric $T_{\\rm eff}$ determinations by Boesgaard), the $T_{\\rm eff}$ determinations (the Cayrel values are based on H$\\alpha$ profiles), and other details.  The Cayrel result is consistent with \\markcite{e86}Eggen's (1986) inference from narrow band photometry that the Pleiades [Fe/H] is near the Hyades value  In order to increase the sample of Pleiades stars with [Fe/H] determinations, some of us (\\markcite{k97}King {\\it et al.\\/}~1997) have used high quality Keck spectra of two slowly rotating very cool ($T_{\\rm eff}{\\sim}4500$ K) Pleiades dwarfs to derive Fe abundances.  Our $T_{\\rm eff}$ values are spectroscopic determinations from balancing the abundances as a function of excitation potential, and the normalized abundances are derived by comparison with similarly analyzed solar data on a line-by-line basis.  The mean abundance is [Fe/H]$=+0.06$, with estimated errors in the mean of perhaps 0.05 dex.  While comparison of the different studies indicates there may be systematic errors at the 0.1 dex level, we regard this (dis)agreement to be quite satisfactory given the ${\\sim}2000$ K range in $T_{\\rm eff}$, the disparate sources of data, and distinct methods used to derive $T_{\\rm eff}$. While a slightly sub-solar Fe abundance is often assumed for the Pleiades based on the Boesgaard \\& Friel results, the totality of the high-resolution spectroscopic evidence may be more consistent with a slightly super-solar value; our photometric [Fe/H] is consistent with solar [Fe/H].  Therefore, if anything the data suggest a distance modulus estimate larger than our MS fitting value rather than smaller.  Fe abundances for Praesepe F dwarfs have been derived by \\markcite{bb88} Boesgaard \\& Budge (1988), Boesgaard (1989), and Friel \\& Boesgaard (1992). The resulting values are $=+0.14$, $+0.10$, and $+0.05$, with star-to-star scatter of 0.06-0.07 dex, and mean uncertainties of 0.03-0.04 dex; again, the internal precision is good.  The zero-reddening assumed in their $T_{\\rm eff}$ determinations is identical to our assumption.  Other detailed studies of numerous Praesepe stars comparison are lacking.  Analysis of the primary component of the Praesepe SB2 KW367, a mid-G star which is significantly cooler than the Boesgaard F stars, by \\markcite{kh96}King \\& Hiltgen (1996) yielded [Fe/H]$=+0.01$ with an uncertainty near 0.05 dex. Again, systematic errors at the 0.1 dex are indicated by this limited comparison.  Combined with the above results, we see that [Fe/H] for Praesepe is 0.00-0.15 dex larger than for the Pleiades, with a preference for the lower middle of this range.  The results inferred from MS fitting are consistent with the upper end of the range.  Boesgaard {\\it et al.\\/}~(1988), Boesgaard (1989), and Boesgaard \\& Friel (1990) derived Fe abundances in $\\alpha$ Per F stars.  The mean [Fe/H] values are $-0.02$, $+0.00$, and $-0.05$.  The $\\alpha$ Per Fe abundance seems nearly identical to the Boesgaard Pleiades estimate.  The star-to-star scatter in the larger $\\alpha$ Per samples is 0.08-0.09 dex; mean uncertainties are ${\\sim}0.04$ dex.  The mean of the individual $\\alpha$ Per reddening values employed by Boesgaard is ${\\sim}0.03$ dex lower than the single value adopted here.  This difference might require a 0.05-0.10 dex increase in [Fe/H] for consistency with our assumptions. \\markcite{bls88}Balachandran {\\it et al.\\/}~(1988) determined Fe abundances in a very wide range (F to K type) of $\\alpha$ Per stars.  The mean abundance of the stars not considered by them to be non-members is [Fe/H]$=+0.04$ with a star-to-star scatter of 0.14 dex; the mean internal error is only 0.02 dex. Their assumed reddening is identical to our value. The results of Boesgaard {\\it et al.\\/}~and Balachandran {\\it et al.\\/}~agree to within 0.1 dex, but when adjustment is made for the slightly different reddening assumptions, the agreement is within a few hundredths of a dex if not exact.  Our photometric [Fe/H] is slightly sub-solar, at the 0.01-0.02 dex level.  It thus appears that the Fe abundance of $\\alpha$ Per is not significantly larger than for the Pleiades.  In sum, internal errors in the Fe abundances of main sequence Pleiades, Praesepe, and $\\alpha$ Per stars derived from careful homogeneous analyses employing high quality data lead to uncertainties of only 0.05-0.10 dex in relative cluster abundances.  We have seen that systematic effects due to errors in reddening, differences in the analysis methodology, {\\it etc.\\/} may approach 0.15 dex. These are small compared to the offset needed to explain the Hipparcos-based M$_V$ values for the Pleiades.  Barring fundamental failure or incompleteness in our understanding of spectral line formation and stellar atmospheres, the extant data suggests that the Fe abundances of the Pleiades, Praesepe, and $\\alpha$ Per are within $\\sim0.10$ dex of each other.  We might caution, however, that the abundances of other important atmospheric opacity contributors ({\\it e.g., Mg and Si}) are, unfortunately, unknown.  \\subsubsection{Photometric Constraints and the Binary Distance to the Pleiades} There are other factors that make a large error in the Pleiades [Fe/H] unlikely.  Colors that incorporate an infrared band are less sensitive to metallicity than \\bv.  The figures in the previous section indicate clearly that the shift in the cluster distance modulus is the same for different color indices; the Pleiades must be intrinsically subluminous if the revised distance estimate is correct.  The deviations from the high-resolution [Fe/H] values for the Pleiades are both large and inconsistent from color to color.  The spectroscopic binary HD 23642 also provides a distance of $5.61\\pm0.26$ consistent with MS fitting, albeit with a large error \\markcite{gia95}(Giannuzzi 1995).  \\subsection{CNO Abundances}  Carbon, nitrogen, and oxygen can affect stellar structure in ways other elements do not; are they anomalous in the Pleiades?  As part of his thesis, King (1993) examined the oxygen abundances of stars in several clusters over a broad range of age.  The [O/H] for the Pleiades was found to be higher than for Praesepe (+0.29 and +0.02 respectively, with errors in the mean of 0.08 for both).  However, the trustworthiness of abundances (such as these) derived from the high excitation 7774 \\AA \\ion{O}{1} lines is a matter of some debate.  Besides possible large data and analysis differences between various studies (e.g. King \\& Hiltgen 1996), there may be significant abundance corrections due to non-LTE effects on line formation in stellar atmospheres (see Garcia Lopez et al. 1995). Unfortunately, systematic errors of 0.3 dex in the cluster O abundances derived from high-excitation lines remains plausible.  In any case, the King results would act to make the Pleiades more metal-rich and therefore require a higher distance modulus estimate.  Detailed abundance studies would be useful, but deviations from the solar mixture would need to be very large to have a significant impact on the luminosity of the MS.  \\subsection{Helium}  The initial solar helium abundance can be inferred from theoretical solar models by the requirement that the model have the solar luminosity at the age of the Sun.  Modern evolution codes give estimates for the initial solar $Y$ in the range $0.26 - 0.28$; the best solar models of Bahcall, Pinsonneault, \\& Wasserburg (1995) had $Y = 0.272$ and $Y = 0.278$ with and without gravitational settling respectively.  A comparison of theoretical stellar models with the Hipparcos main sequence of the Hyades by Perryman et al. (1997) yields $Y = 0.26\\pm0.02$; for comparison, the solar $Y$ in that study was 0.266 and the solar-scaled helium for the cluster would be 0.28. This agreement between the Sun and Hyades was anticipated and reinforces the notion that stars formed in the current epoch have similar helium abundances.  Nevertheless, we consider what range of $Y$ would be needed to drop the Pleiades main sequence by 0.3 mag, and that value is about $Y = 0.37$.  Such a high value of $Y$ for the Pleiades would imply a drastic revision of chemical evolution models and, by extension, would raise the possibility that other clusters might have similar anomalies.  MS fitting would therefore require knowledge of both the metal and helium abundances; since helium can only be directly observed in young systems this would make MS fitting unreliable at the 0.3 magnitude level for the majority of clusters.  We believe that this question is best answered by direct measurements of the helium abundance in HII regions and massive stars.  We begin with a discussion of the literature on helium abundances; we have also obtained data on the relative helium abundances in the Pleiades and $\\alpha$ Per.  Neither the field star data nor our Pleiades spectra are consistent with significant variations in the initial helium abundance from the solar value.  Ignoring a deviant few percent of field stars, \\markcite{n74}Nissen's (1974) study revealed no intrinsic scatter in $Y$ greater than ${\\sim}10$\\% (compared to the 30-40\\% deviation required by the Pleiades stars) in nearby main-sequence field B stars.  \\markcite{gl92} Gies \\& Lambert (1992) found helium abundances consistent with both the Sun and the Orion nebula for a sample of 35 B dwarfs; 4 B supergiants in that sample were found to have anomalously high helium abundances.  There is evidence that evolutionary effects are responsible for helium enrichment in the most massive stars (see \\markcite{mc94}Maeder \\& Conti 1994, \\markcite{lu96}Lyubimkov 1996, \\markcite{pin97}Pinsonneault 1997 for reviews), so helium abundances from MS O stars and massive supergiants may not be reliable indicators of the initial $Y$.  The B star field data and the Orion nebula abundances are therefore our best test for the range in helium abundance at solar metal abundance, and they are consistent with only small variations in the initial helium abundance.  For Galactic clusters, however, the picture is less clear. \\markcite{ss69}Shipman \\& Strom (1969), \\markcite{ps73} Peterson \\& Shipman (1973), \\markcite{n76}Nissen (1976), and \\markcite{l77}Lyubimkov (1977) found evidence for 20\\%-30\\% variations in $Y$ among young associations, including some systems with significantly lower $Y$.  Lyubimkov suggested an increasing He abundance with {\\it increasing\\/} age amongst the young clusters/associations studied, a conclusion not supported by the subsequent field star work of Gies \\& Lambert.  \\markcite{p79}Patton (1979) determined He abundances of 60 stars in 8 young clusters and associations.  She noted that her initial abundances displayed a range in $Y$ of about 25\\%, and that this could not be explained by the the usual error sources; she also called attention to a correspondence between He abundance and cluster age.  However, Patton shows that binarity may be responsible for observed cluster-to-cluster He abundance dispersions, and the notably low He abundances (observed by others too) seen for a few stars within a given cluster/association.  Eliminating {\\it suspected\\/} (but not positively identified) binary systems from her analysis results in cluster He abundances which are identical to within the uncertainties.  This highlights the need for secure knowledge of very fundamental stellar parameters (e.g., binarity) before reliable He abundances can be derived.  With this muddled picture of main sequence stellar He abundances, one may wonder if the Pleiades He abundance could be abnormal.  Both the Pleiades and $\\alpha$ Per are young enough to have B stars, and their helium can be directly measured.  The $Y$ values from Lyubimkov (1977) agree to within ${\\Delta}Y{\\sim}0.015$, which is well within the uncertainties; the Pleiades and $\\alpha$ Per $Y$ value is 0.04 larger than the corresponding field star value, but the uncertainties are comparable to this offset.  \\markcite{kp86} Klochkova \\& Panchuk (1986) also derived B-star He abundances in both the Pleiades and $\\alpha$ Per.  They claim to find no difference between the mean abundances that is larger than the uncertainties.  However, this conclusion is not clear to us from the abundances listed in their Table II, which do demonstrate quite a very large difference.  Unfortunately, only two Pleiades stars are included in the analysis.  Therefore small number statistics and the possible effects of binarity make assessment of this difference quite difficult.  We attempted a final comparison using the ``field'' stars from Nissen (1974).  This sample includes four $\\alpha$ Per stars, and two stars (HR 5191 and 7121) which are suggested members of the purported Pleiades supercluster.  The mean $Y$ value is only 0.03 larger for the Pleiades field stars than for the $\\alpha$ Per stars; the uncertainties are probably not much smaller than this difference.  To investigate the possibility of a non-standard helium abundance in the Pleiades \\markcite{fk98} Fischer \\& King (1998) observed MS B stars in $\\alpha$ Per and Pleiades to differentially compare the helium abundances. Preliminary analysis of the lines strengths for six He lines suggests that the cluster He abundances are identical within an uncertainty of 15\\%.  Any real difference appears to be in the opposite sense of what is needed to make the Pleiades underluminous: the Pleiades line strengths are, if anything, consistently smaller than the $\\alpha$ Per counterparts.  \\subsection {Reddening and Systematic Errors in the Photometry}  Reddening will tend to make a cluster MS fainter at a given color.  If the reddening is increased the inferred distance modulus will therefore increase. The effect can be roughly estimated as follows : in the color interval that we are using for MS fitting the derivative of $M_V$ with respect to both \\bv\\ and $V-I$ is $\\sim$5.  The extinction $A_V$=3.12$E(B-V)$ and $E(V-I)_K$=1.5$E(B-V)$.  Adding these effects together an increase in $E(B-V)$ of 0.10 magnitudes would increase V at fixed \\bv\\ and fixed $V-I$ by 0.188 (0.5 mag from a shift of 0.1 in \\bv\\ - 0.312 mag from extinction) and 0.438 (0.75 mag from a shift of 0.15 in $V-I$ - 0.312 mag from extinction) magnitudes respectively.  The relative distances inferred by the two colors can therefore be affected if the reddening is incorrect. In addition the [Fe/H] abundances derived for cluster stars are sensitive to T$_{eff}$, and an increased reddening would imply a higher [Fe/H] for a given equivalent width (therefore further increasing the distance modulus). Other colors, such as $R-I$, will be less reddening-sensitive.  Neither the Hyades nor Praesepe show any evidence for reddening along the line of sight; increasing the reddening estimate for the Pleiades would worsen the discrepancy with the Hipparcos distance modulus estimate. Even changing $E(B-V)$ from 0.04 to 0 would only decrease the distance modulus by 0.08 magnitudes.  The reddening estimates for the Pleiades have been derived for a wide range of masses and from different techniques; Crawford and Barnes used Stromgren photometry to estimate $A_V$ for B, A, and early F stars in the Pleiades and Praesepe, Prosser and Stauffer used M dwarfs in the same clusters, and Breger used polarization measurements in the Pleiades. We conclude that reddening is not a significant source of uncertainty in distance estimates for the Pleiades.  Multicolor distance measurements of the type performed in this paper could be a useful check on the reddening for more heavily obscured systems.  Another possibility is that systematic errors in the photometry could cause errors in the distance estimates.  For the color range that we are considering, the slope of the MS is $\\sim$5; this would require a systematic error of 0.06 magnitudes in \\bv\\ to reconcile the Pleiades distance scales, which is unreasonably large.  The size of the systematic errors can be constrained by comparing spectroscopic temperature estimates with those based upon colors.  In the case of Coma Ber, for example, it appears that spectroscopic temperature estimates are in agreement with the \\bv\\ colors of F stars but not with the $V-I$ colors.  We note that the slope of the MS in $V-I$ is steeper for F stars than for the cooler stars, and that systematic errors in the $V-I$ photometry might explain the puzzling behavior of Coma Ber.  We have attempted whenever possible to rely upon a single source for photometry in a given color for a given cluster.  Even in the case of the $V-I$ data we see no evidence of systematic differences between the location on the color magnitude diagram of stars with colors converted to the Cousins system from the Kron system and those converted to the Cousins system from the Johnson system.  For the Pleiades, independent studies (Section 2.2) give consistent photometry for individual stars at the level of the quoted errors (0.01 - 0.02 mag).  The 0.3 mag discrepancy between the Hipparcos and  MS fitting distance distance modulii is much too large to be explained by systematic errors in the photometry.  High-resolution spectroscopy of the Pleiades is consistent with the observed colors, and the reddening is small.  For systems with higher reddening, however, care must be taken when converting between different photometric systems; the Johnson, Cousins, and Kron system I bands have different effective central wavelengths and therefore different reddening corrections.  \\subsection{Systematic Errors in the Hipparcos Parallaxes}  The final possibility is that the Hipparcos Pleiades parallaxes may contain previously undetected systematic errors.  If the MS fitting result $m-M = 5.60$ does indeed give the correct Pleiades distance, then a systematic zero-point error would need to approach the 1 mas level to produce the discordance with the Hipparcos results. Such an error seems impossibly large, in view of the extensive tests \\markcite{ar95,ar97} (Arenou et al. 1995, 1997) demonstrating the global zero-point error of the Hipparcos parallaxes to be smaller than 0.1 mas.  However, global tests have little power to reveal effects occurring on the small angular scale ($\\sim 1 \\deg$) of the Hipparcos spatial correlations (see below). Indeed, the Hipparcos parallaxes of stars in open clusters such as the Pleiades represent the first real opportunity to test for systematic effects on small angular scales.  One might well argue that it would only be prudent to consider the Hipparcos cluster results as the first direct tests for small-scale zero-point errors, rather than as cluster distance measurements.  The Hipparcos Pleiades parallax (van Leeuwen \\& Hansen Ruiz 1997a) is based on measurements of 54 cluster members, ranging in $V$ from 2.8 to 11.5 within $5\\deg$ of the cluster center, so it represents a fairly broad sampling of the cluster. Because Hipparcos observed widely separated ($\\sim 58\\deg$ apart) star fields simultaneously, the parallaxes are inherently on an absolute scale over the whole sky.  Over small regions of the sky ($\\lesssim 2\\deg$), however, the astrometric results are positively correlated because neighboring stars (within the $0.9\\deg \\times \\ 0.9\\deg$ Hipparcos field of view) tended to be observed on the same great circles the satellite swept out over the sky (Lindegren 1988, 1989).  A comprehensive discussion of the Hipparcos mission and data reductions is given in Volumes 1--3 of the Hipparcos Catalogue \\markcite{ESA97}(ESA 1997).  The spatial correlations may significantly impact the astrometric results for star clusters, whose angular size is of the same order as the Hipparcos correlation scale.  To account for this, \\markcite{vh97a} van Leeuwen \\& Hansen Ruiz (1997a) re-calculated the Pleiades mean parallax from the intermediate Hipparcos data. For this paper, one of us (R.B.H.) has re-examined the individual Pleiades parallaxes from the Hipparcos Catalogue.  Moreover, besides the spatial correlations, there is a different type of correlation affecting the Hipparcos results -- the statistical correlations among the five astrometric parameters (position, proper motion, and parallax), arising from the imperfect distribution of Hipparcos observations on the sky over time.  In classical parallax work (cf. \\markcite{va75} Vasilevskis 1975), the time distribution of observations over a star's parallactic ellipse is controlled to maximize the parallax factors and minimize the correlations between position, proper motion, and parallax.  This is easy to achieve from the ground, but Hipparcos could not do this because of the limited span of observations and the pattern of scans of the sky, as explained in Section 3.2.4 (pp. 321-325) of the Hipparcos Introduction (ESA 1997, Vol. 1).  Figures 3.2.42 to 3.2.61 of that work illustrate the patterns of the correlations over the sky; Figure 3.2.66 (p. 363) shows histograms of the 10 correlations.  The RMS values are $\\sim 0.2$, and large areas of the sky show correlations averaging 0.4 or more in size.  It must be emphasized that these correlations are substantially larger than would be considered acceptable in ground-based parallax observations.  For parallax work, the most important correlation is $\\rho_\\alpha^\\pi\\ $, between parallax and right ascension (Field H20 in the Hipparcos Catalogue). This is because, over most of the sky, most of the extent of the parallactic ellipse is in right ascension.  The Hipparcos $\\rho_\\alpha^\\pi$ correlation is shown in Fig. 3.2.44 of the Hipparcos Introduction.  Large values of $\\rho_\\alpha^\\pi$ were caused in certain areas of the sky by the unfortunate circumstance of unequal observations on both sides of the Sun, as discussed on p. 325 of the Hipparcos Introduction.  This happens to impact the Pleiades particularly badly.  The mean value of $\\rho_\\alpha^\\pi$ near the Pleiades center is +0.4; this is at the 96th percentile in the histogram in Fig 3.2.66. The question this raises is whether this large correlation, caused by the time distribution of Hipparcos observations of the Pleiades stars, has any effect on the parallax values.  We stress again that this is a different effect from the spatial correlation that exists because Hipparcos astrometric data over small ($\\sim 1 \\deg$) areas of the sky are not fully independent measurements.  In Figure 19 we plot parallax vs. the correlation $\\rho_\\alpha^\\pi$ for 49 Pleiades members verified by proper motion, radial velocity, and position in the color-magnitude diagram. (Mermilliod et al's 51 stars and van Leeuwen et al's 54 are virtually the same set as these; we rejected several additional stars on account of problems noted in Fields H30 and H59 of the Hipparcos Catalogue.)  This plot shows several interesting things.  The filled symbols are 12 bright ($V < 7$) stars within $\\sim 1 \\deg$ of the cluster center with correlations $\\rho_\\alpha^\\pi \\geq +0.34$ (the mean value for the whole sample).  Due to the spatial correlation effect, these 12 stars all have nearly the same parallax (mean 8.86 mas, RMS dispersion 0.45 mas; $\\chi^2$ too small at the 0.995 significance level). Because Hipparcos' errors are smallest for bright stars, these stars carry much of the weight of the Pleiades parallax.  There is a clear trend (slope) of parallax vs. $\\rho_\\alpha^\\pi$ correlation; a weighted least-squares solution gives a slope of $+3.04 \\pm 1.36$ mas per unit correlation.  The solid line in Fig. 19 is this slope, run through the mean point (+0.34,+8.53).  The dashed lines show $\\pm1\\sigma$ slopes.  The intercept at zero correlation is $\\pi = 7.49 \\pm 0.50$ mas, quite consistent with the MS fitting distance.  Figure 20 plots parallax vs. distance from the cluster center.  The filled symbols are the same 12 bright stars with high $\\rho_\\alpha^\\pi$ as in Fig. 16.  The open symbols are the 15 stars with $\\rho_\\alpha^\\pi < +0.25$, with no restriction on magnitude or distance.  The two sets of stars barely overlap because the brightest stars in the Pleiades are highly concentrated to the cluster center.  The low-correlation stars lie farther from the Pleiades center and show a much larger parallax scatter, reflecting (a) the larger errors for fainter stars and (b) the lack of spatial correlations on scales $\\gtrsim 1 \\deg$. Moreover, their mean parallax is smaller (reflecting the slope discussed above).  For the 15 stars with $\\rho_\\alpha^\\pi < +0.25$, the weighted mean parallax is $7.46 \\pm 0.43$ mas.  The RMS dispersion is 1.66 mas, consistent with the published parallax errors.  This exercise is not intended to be a definitive re-determination of the Pleiades parallax; that would require going back to the intermediate Hipparcos data as per van Leeuwen et al (1997), and exploring the effects of both the $\\rho_\\alpha^\\pi$ and the spatial correlations at that level. However, it is quite clear that (a) small-angular-scale systematic effects at the 1 mas level are present in the Hipparcos Pleiades parallaxes; (b) these effects are related to the high values of the $\\rho_\\alpha^\\pi$ correlation near the cluster center; (c) the bright stars within $\\sim 1\\deg$ of the center, which carry most of the weight of the mean parallax, are the most severely affected; and (d) the stars with lower $\\rho_\\alpha^\\pi$ correlations, far enough ($\\gtrsim 1 \\deg$) from the center to be unaffected by the spatial correlation, have smaller parallaxes, consistent with the MS fitting distance.  We also looked for effects of the $\\rho_\\alpha^\\pi$ correlation in the Hyades, Praesepe, $\\alpha$~Per, and Coma~Ber clusters.  In Figures 21--24 we present the parallax vs. correlation plots for those clusters.  The Hyades, Praesepe, and $\\alpha$~Per clusters also have large values of $\\rho_\\alpha^\\pi$, but the the slope (d$\\pi$/d$\\rho$) present in the Pleiades data does not occur in these clusters, where the MS fitting distances and the Hipparcos distances are in good agreement.  The data for Coma~Ber do show a slope d$\\pi$/d$\\rho \\ = \\ -4.0 \\pm 2.1$ mas, but the range of $\\rho_\\alpha^\\pi$ is small, and the mean is near zero."
    },
    "9803/astro-ph9803005_arXiv.txt": {
        "abstract": "We report on a mechanism which may lead to a spin-up of the surface of a rotating single star leaving the Hayashi line, which is much stronger than the spin-up expected from the mere contraction of the star.  By analyzing rigidly rotating, convective stellar envelopes, we qualitatively work out the mechanism through which these envelopes may be spun up or down by mass loss through their lower or upper boundary, respectively.  We find that the first case describes the situation in retreating convective envelopes, which tend to retain most of the angular momentum while becoming less massive, thereby increasing the specific angular momentum in the convection zone and thus in the layers close to the stellar surface.  We explore the spin-up mechanism quantitatively in a stellar evolution calculation of a rotating $12\\,\\Msun$ star, which is found to be spun up to critical rotation after leaving the red supergiant branch.  We discuss implications of this spin-up for the circumstellar matter around several types of stars, i.e., post-AGB stars, {\\Be} stars, pre-main sequence stars, and, in particular, the progenitor of Supernova 1987A. ",
        "introduction": "\\lSect{intro}  The circumstellar matter around many stars shows a remarkable axial symmetry.  Famous examples comprise Supernova~1987A (Plait et al. 1995; Burrows et al. 1995), the Homunculus nebula around $\\eta$~Carina and other nebulae around so called Luminous Blue Variables (Nota et al. 1995), and many planetary nebulae (Schwarz et al. 1992).  A less spectacular example are {\\Be} stars, blue supergiants showing properties which might be well explained by a circumstellar disk (Gummersbach et al. 1995; Zickgraf et al. 1996). Many of these axisymmetric structures have been explained in terms of interacting winds of rotating stars (cf. Martin \\& Arnett 1995; Langer et al. 1998; Garc\\'{\\i}a-Segura et al. 1998), which may be axisymmetric when the stars rotate with a considerable fraction of the break-up rate (Ignace et al. 1996; Owocki et al. 1996).  However, up to now only little information is available about the evolution of the surface rotational velocity of stars with time, in particular for their post-main sequence phases.  Single stars which evolve into red giants or supergiants may be subject to a significant spin-down (Endal \\& Sofia 1979; Pinsonneault et al. 1991).  Their radius increases strongly, and if the specific angular momentum were conserved in their surface layers (which may not be the case; see below) they would not only spin down but they would also evolve further away from critical rotation.  Moreover, they may lose angular momentum through a stellar wind.  Therefore, it may appear doubtful at first whether post-red giant or supergiant single stars can retain enough angular momentum to produce aspherical winds due to rotation.  However, by investigating the evolution of rotating massive single stars, we found that red supergiants, when they evolve off the Hayashi line toward the blue part of the Hertzsprung-Russell (HR) diagram may spin up dramatically, much stronger than expected from local angular momentum conservation.  In the next Section, we describe the spin-up mechanism and its critical ingredients.  In Section~3 we present the results of evolutionary calculations for a rotating $12\\,\\Msun$ star, which provides a quantitative example for the spin-up.  In Section~4 we discuss the relevance of our results for various types of stars, and we present our conclusions in Section~5. ",
        "conclusions": "\\lSect{con}  In this paper, we discussed the effect of mass outflow through the inner or outer boundary of a rigidly rotating envelope on its rotation frequency.  It causes a change of the specific angular momentum in the envelope and alters its rotation rate besides what results from contraction or expansion (cf. \\Fig{t-v}). For constant upper and lower boundaries of the envelope, which we found a good approximation for convective envelopes (cf. \\Fig{t-r-m}), a spin-down occurs for mass outflow through the upper boundary --- which corresponds, e.g., to the case of stellar wind mass loss from a convective envelope (cf. also Langer \\mbox{1998) ---,} while a spin-up results from mass outflow through the lower boundary (cf. \\Fig{panels}).  The latter situation is found in evolutionary models of a rotating $12\\,\\Msun$ star at the transition from the Hayashi-line to the blue supergiant stage.  The star increased its rotational velocity one order of magnitude above the velocity which would result in the case of local angular momentum conservation.  It would have increased its rotational velocity even further if it would not have arrived at critical rotation (cf. \\Sect{res}, \\Fig{t-v}), with the consequence of strong mass and angular momentum loss.  At this point, the specific angular momentum loss $\\Jdot/\\Mdot$ reached about $8\\,10^{19}\\,\\junit$ (cf. \\Fig{t-xx}).  The geometry of circumstellar matter around stars which undergo a red $\\to$ blue transition may be strongly affected by the spin-up.  We propose that this was the case for the progenitor of SN~1987A, the only star of which we know that it performed a red $\\to$ blue transition in the recent past. The blue supergiant in its neighborhood studied by Brandner et al. (1997), around which they found a ring nebula as well, is another candidate.  Also, {\\Be}~stars may be related with the post-red supergiant spin-up (cf. \\Sect{blue-loop}). Furthermore, the spin-up mechanism studied in this paper may be relevant for bipolar outflows from central stars of proto-planetary nebulae (\\Sect{post-AGB}), from stars in the transition phase from the red supergiant stage to the Wolf-Rayet stage (\\Sect{RSG-WR}), and from pre-main sequence stars (\\Sect{pre-MS})."
    },
    "9803/astro-ph9803280_arXiv.txt": {
        "abstract": "I present 1.5- and 8.4-GHz observations with all configurations of the NRAO VLA of the wide-angle tail source \\Ss{3C130}. The source has a pair of relatively symmetrical, well-collimated inner jets, one of which terminates in a compact hot spot. Archival {\\it ROSAT} PSPC data confirm that 3C\\,130's environment is a luminous cluster with little sign of sub-structure in the X-ray-emitting plasma. I compare the source to other wide-angle tail objects and discuss the properties of the class as a whole. None of the currently popular models is entirely satisfactory in accounting for the disruption of the jets in 3C\\,130. ",
        "introduction": "\\Ss{3C130} is a FRI radio source at redshift 0.109 (Spinrad \\etal\\ 1985). Its 178-MHz luminosity is $7.6 \\times 10^{25}$ W Hz$^{-1}$ sr$^{-1}$, slightly above the nominal FRI-FRII boundary of $\\sim 2 \\times 10^{25}$ W Hz$^{-1}$ sr$^{-1}$ (Fanaroff \\& Riley 1974, hereafter FR). Leahy (1985, 1993) and J\\\"agers and de Grijp (1985) present intermediate-resolution VLA maps of the central regions of the source, while J\\\"agers (1983) has a lower-resolution WSRT image which shows the whole source and its field; the source extends for $\\sim 1.5$ Mpc. Saripalli \\etal\\ (1996) present high-frequency maps made with the Effelsberg 100-m telescope. The host galaxy is classed as a DE2 by Wyndham (1966) and appears to lie in a cluster, although strong galactic reddening makes optical identification of the cluster members difficult. The {\\it Einstein} detection of extended X-ray emission (Miley \\etal\\ 1983), the near\\-by align\\-ed sources (J\\\"agers 1983) and the many mJy radio sources in the field at 1.5 GHz make it plausible that the object is the dominant member of a large cluster. Leahy (1985) also attempts to constrain the RM distribution of the source, but notes that it depolarizes rapidly (particularly in the S lobe) so that few good measurements are available; this could be taken as evidence for a dense magneto-ionic environment for the source (cf.\\ Hydra A, Taylor \\etal\\ 1990).  \\label{definition} 3C\\,130 is a wide-angle tail (WAT) radio source. The term WAT has been used to describe many different types of object. Here I shall use it to refer to those FRI sources which are associated with central cluster galaxies (e.g.\\ Owen \\& Rudnick 1976) and have luminosities comparable to or exceeding the Fanaroff-Riley break between FRI and FRII. I shall follow Leahy (1993) in using the behaviour of the jets at the base as another defining feature. At high resolution one or two well-collimated jets [`strong-flavour' jets, by the classification of Leahy (1993)] are seen (e.g.\\ O'Donoghue, Owen \\& Eilek 1990), extending for some tens of kpc before broadening, often at a bright flare point, into the characteristic plumes or tails. These jets are very similar to the jets seen in FRII radio galaxies, and quite different from the behaviour of jets in more typical FRIs, where a collimated inner jet, if visible at all, decollimates rapidly (on scales of a few kpc at most) and comparatively smoothly into a bright `weak-flavour' jet with a large opening angle.\\footnote{There are a few exceptions to this behaviour; \\Ss{3C66B} (Hardcastle \\etal\\ 1996) does appear to show an inner `strong-flavour' jet and a bright knot at the base of the `weak-flavour' jet. But even here the transition from strong to weak flavours occurs on scales of $\\sim 1$ kpc.} WATs, according to this definition, never have a weak-flavour jet, but make the transition between strong-flavour jet and diffuse, bent tail in a single step. The requirement that WATs be central cluster galaxies excludes objects (e.g. \\Ss{3C171}, Blundell 1996, Hardcastle \\etal\\ 1997a; \\Ss{3C305}, Leahy 1997) where the `tails' are likely to be simply ordinary FRII lobes which have been disrupted by unusual host-galactic dynamics. The condition on jet behaviour allows us to exclude objects such as the twin sources in \\Ss{3C75} (Owen \\etal\\ 1985; Hardcastle 1996) which are associated with a dominant cluster galaxy and sometimes classed as WATs but whose inner jets are similar to those of typical powerful FRIs.  Because of the requirements of this definition, wide-angle tail sources make up a small minority of the radio source population. For this reason, the detailed properties of their jets and tails have not been well studied, although a number have been imaged for studies of source dynamics (O'Donoghue \\etal\\ 1989). The only objects which have been the subject of detailed study in the radio are \\Ss{3C465} (Leahy 1984; Eilek \\etal\\ 1984) and \\Ss{3C218}, Hydra A (Taylor \\etal\\ 1990), although M87, Virgo A (e.g.\\ Biretta \\& Meisenheimer 1993) exhibits some of the properties of a WAT. In this paper I present multi-configuration, multi-frequency VLA observations of a further powerful WAT.  Throughout this paper I use a cosmology in which $H_0 = 50{\\rm\\ km\\,s^{-1}\\,Mpc^{-1}}$ and $q_0 = 0$. At the distance of \\Ss{3C130}, one arcsecond is equivalent to a projected length of 2.72 kpc. B1950.0 co-ordinates are used throughout. ",
        "conclusions": "A compact hot spot is detected at the base of one plume of the WAT \\Ss{3C130}, and the jets are shown to have longitudinal magnetic field. The source is thus very like a classical double in some respects. The data support the model in which WATs are objects whose jets make the transition from super- to sub-sonic velocities in one step, rather than decelerating gradually, by showing a bright sub-kpc structure (comparable to those seen in classical double radio sources) associated with the termination of a jet.  Archival {\\it ROSAT} PSPC observations of 3C\\,130 show it to lie in a luminous cluster with $kT \\sim 2.9$ keV. There is little sign of substructure in the X-ray, in contrast to many other WATs; this may be related to the nearly straight tails of 3C\\,130. The lack of strong substructure seems to be inconsistent with recent models for jet disruption in WATs."
    },
    "9803/astro-ph9803249_arXiv.txt": {
        "abstract": "We predict the rate at which Gamma-Ray Burst (GRB) afterglows should be detected in supernova searches as a function of limiting flux.  Although GRB afterglows are rarer than supernovae, they are detectable at greater distances because of their higher intrinsic luminosity.  Assuming that GRBs trace the cosmic star formation history and that every GRB gives rise to a bright afterglow, we find that the average detection rate of supernovae and afterglows should be comparable at limiting magnitudes brighter than $K=18$.  The actual rate of afterglows is expected to be somewhat lower since only a fraction of all $\\gamma$--ray selected GRBs were observed to have associated afterglows.  However, the rate could also be higher if the initial $\\gamma$--ray emission from GRB sources is more beamed than their late afterglow emission. Hence, current and future supernova searches can place strong constraints on the afterglow appearance fraction and the initial beaming angle of GRB sources. ",
        "introduction": "Since their discovery in the late 1960's (Klebasadel et al. 1973) through early 1997, Gamma-Ray Bursts (GRBs) had defied all attempts to determine their distance scale conclusively.  The Burst And Transient Source Experiment (BATSE) on board the Compton Gamma-Ray Observatory (GRO) showed that the burst population is highly isotropic (Meegan et al. 1993; Briggs et al. 1993), suggesting that bursts occur at cosmological distances or in an extended Galactic halo. Moreover, the cumulative number counts of faint bursts deviated from that of a uniform distribution of sources in Euclidean space and flattened at faint fluxes, consistent with the expected effect of a cosmological redshift (Fishman \\& Meegan 1995, and references therein).  Last year, with the advent of the BeppoSAX satellite (Boella et al.  1997), it became possible to localize GRB sources to within an arcminute on a timescale of hours.  Such fast, accurate localizations were quickly followed by the detection of delayed X-ray (Costa et al.  1997), optical (van Paradijs et al. 1997), and radio (Frail et al.  1997) counterparts to GRB sources.  In particular, FeII and MgII absorption lines were detected at a redshift $z=0.835$ in the spectrum of the optical counterpart to GRB970508 (Metzger et al. 1997), demonstrating conclusively that this burst occurred at a cosmological distance with a redshift $z>0.835$.  The isotropy of the burst population and the flattening of their number counts, taken in combination with the fact that the first confirmed redshift for an optical counterpart is high, provides strong evidence that GRB sources are located at cosmological distances.  Most plausible GRB models involve either the collapse of a single massive star (e.g. Usov 1992; Woosley 1993; Paczy\\'nski 1998), or the coalescence of two compact objects -- two neutron stars or a neutron star and a black hole -- in a binary system (e.g.  Paczy\\'nsky 1986; Eichler et al. 1989; Narayan et al.  1992; Mochkovitch et al. 1993; Rees 1997).  Since the lifetime of these progenitors is short compared to the Hubble time at a redshift $z\\la 5$, the cosmic GRB rate should simply be proportional to the star formation rate at these redshifts, without any appreciable delay due to the finite progenitor lifetime.  The cosmic rate of massive star formation rate has been determined from the $U$ and $B$-band luminosity density in Hubble Deep Field (Madau et al.  1996; Madau 1996; Madau, Pozzetti, \\& Dickinson 1997; Madau 1997).  The inferred star formation rate $\\dot\\rho_{\\rm s}(z)$ can then be converted to a GRB explosion rate $R_{\\rm GRB}(z)$, based on the requirement that the latter would fit the observed number count distribution of $\\gamma$--ray selected GRBs (Wijers et al. 1997).  Cosmological GRBs are at least $10^4$ times rarer than Type II supernovae (SNeII) -- possibly even $\\sim 10^6$ times rarer if GRBs occur primarily at high redshifts following the cosmic star formation history (Wijers et al. 1997).  However, at peak luminosity, the GRB afterglows are $\\sim 10^3$--$10^4$ times brighter than SNeII.  In Euclidean space, this would imply that GRBs are detected from a volume bigger by a factor $\\sim (10^{4})^{3/2}= 10^6$, roughly canceling out the factor by which they are rarer than supernovae.  Hence we expect that at some relatively bright limiting flux, the rate of afterglow detections should become comparable to that of SN detections.  Current and future supernova searches should provide information about the fraction of GRBs which produce detectable afterglows. The statistics of bursts in 1997 for which afterglows could have been identified implies that this fraction is of order tens of percent (e.g., Castro-Tirado 1998).  On the other hand, there could also be a population of afterglows without a GRB precursor.  This would occur if the source emits a jet from which the $\\gamma$--ray emission is more beamed than the subsequent optical afterglow radiation due to the deceleration of the jet by the ambient gas and the corresponding decline in its relativistic beaming with time (Rhoads 1997).  A jet geometry would imply a higher rate of afterglow detections in supernova searches.  In this {\\it Letter}, we predict the detection frequency of GRB afterglows as a function of limiting flux at various observed wavelengths, and compare this rate with the analogous predictions for SNe Type Ia and Type II at high redshifts.  We assume throughout a flat, $\\Omega =1$, $\\Lambda =0$, cosmology, with a Hubble constant $H_0=50$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "9803/astro-ph9803139_arXiv.txt": {
        "abstract": "We analyse the population of bright star clusters in the interacting galaxy pair NGC 4038/39 detected with HST WFPC1 by Whitmore \\& Schweizer (1995). Making use of our spectrophotometric evolutionary synthesis models for various initial metallicities we derive the ages of these star clusters and calculate their future luminosity evolution. This allows us to compare their luminosity function ({\\bf LF}), evolved over a Hubble time, to LFs observed for the Milky Way's and other galaxies' star cluster systems. Since effective radii are difficult to determine due to crowding of the clusters, the shape of the LF after a Hubble time may help decide whether the young clusters are young globular clusters ({\\bf GC}) or rather open clusters/OB associations. We find an intriguing difference in the shapes of the LFs if we subdivide the cluster population into subsamples with small and large effective radii. While the LF for the extended clusters looks exponential, that for clusters with small effective radii clearly shows a turn-over brighter than the completeness limit. For other possible subdivisions as to luminosity or colour no comparable differences are found. Evolving, in a first step, the LF from a common mean age of the young clusters of 0.2 Gyr to an assumed age of 12 Gyr, the LF for the subsample of clusters with small effective radii seems compatible with a Gaussian GCLF with typical parameters M$_{\\rm V_0} = -7.1$ and $\\sigma (\\rm M_{\\rm V_0}) = 1.3$ except for some overpopulation of the faint bins. These faintest bins, however, are suspected to be subject to the strongest depopulation through effects of dynamical evolution not included in our models. We also follow the colour evolution of the young star clusters over a Hubble time and compare to observations on the Milky Way and other galaxies' GC systems.  For an ongoing starburst like the one in the NGC 4038/39 system age spread effects among the young star cluster population may not be negligible. In a second step, we therefore account for age spread effects, instead of using a mean age for the young cluster population, and this drastically changes the time evolution of the LF, confirming Meurer's (1995) conjecture. We find that $-$ if age spread effects are properly accounted for $-$ the LF of the entire young star cluster population, and in particular that of the brighter subsample, after a Hubble time is in good agreement with the average Gauss-shaped LF of globular cluster systems having a turn-over at $\\langle {\\rm M_{V_0}} \\rangle = -7.1$ mag and $\\sigma({\\rm M_{V_0}}) = 1.3$ mag.  The age distribution shows that the brightest globular clusters from the interacting galaxies' original population are also observed. They make up the bulk of the red subpopulation with (V$-$I)$_0 > 0.95$. Their effective radii do not significantly differ from those of the young star cluster population, neither on average nor in their distribution.  We discuss the influence of metallicity, the effects of an inhomogeneous internal dust distribution, as well as the possible influence of internal $-$ through stellar mass loss $-$ and external dynamical effects on the secular evolution of the LF.  Referring YSC luminosities to a uniform age and combining with model M/L, we recover the intrinsic mass distribution of the YSC system. It is Gaussian in shape to good approximation thus representing a quasi-equilibrium distribution that $-$ according to Vesperini's (1997) dynamical modelling for the Milky Way GC system $-$ will {\\bf not} be altered in shape over a Hubble time of dynamical evolution, allthough a substantial number of clusters will be destroyed.  We briefly compare the young star cluster population of the Antennae to the older one in the merger remnant NGC 7252 and point out that the intercomparison of young cluster populations in an age sequence of interacting and merged galaxies may become an interesting approach to study in detail the role of external dynamical effects. ",
        "introduction": "From the fact that $-$ when normalised to the stellar mass of a galaxy $-$ the specific globular cluster ({\\bf GC}) frequency $T_{GC} := {N_{GC} \\over {M_{\\ast} / 10^9~M_{\\odot}}}$ is a factor of $\\sim 2$ higher in ellipticals than in spirals, Zepf \\& Ashman (1993) predict that if elliptical galaxies are formed from one major spiral $-$ spiral merger the number of GCs formed during the merger- induced starburst should be of the same order of magnitude as the number of GCs present in the progenitor galaxies.  The high burst strengths and star formation ({\\bf SF}) efficiencies in massive gas-rich spiral $-$ spiral mergers and in IR-ultraluminous galaxies led to expect the formation of star clusters so tightly bound that they are able to survive as GCs (Fritze $-$ v. Alvensleben \\& Gerhard 1994).  Fritze $-$ v. Alvensleben \\& Gerhard (1994) predicted the metallicity range of stars and star clusters formed in massive gas-rich (i.e. late type) spiral$-$spiral mergers on the basis of the ISM abundances of the progenitor galaxies to be $\\third ~{\\rm Z_{\\odot} \\lta Z \\lta Z_{\\odot}}$ or $-0.8 \\lta {\\rm [Fe/H]} \\lta -0.2$.  In many interacting galaxies and merger remnants, bright blue knots have by now been observed (cf. e.g. Lutz 1991, Holtzman \\etal 1992, Whitmore \\etal 1993, Hunter \\etal 1994, O'Connell \\etal 1994, 1995, Conti \\& Vacca 1994, Borne 1996, Meurer \\etal 1995). These bright blue knots, of course, immediately raised the question as to their identity: are these Young Star Clusters ({\\bf YSC}) $-$ or, at least, some of them $-$ the progenitors of GCs? And, if the latter were true, how many of them are typically formed in a merger? How many will be able to survive in the tidal field of two massive interacting spirals? Can such a higher metallicity subpopulation be identified in GC systems (hereafter {\\bf GCS}) around merger remnants and perhaps even around normal ellipticals? Could the metallicity distribution of a GCS give information about the origin of its parent galaxy (cf. Zepf \\& Ashman 1993)? Or should all of these bright blue knots be open clusters/OB associations (van den Bergh 1995) most of which will disperse within few Gyr? The discussion of the nature of these YSCs is focussed on two aspects, their effective radii R$_{\\rm eff}$ and their luminosity function. In mergers at distances of the Antennae or NGC 7252, effective radii as measured on WFPC1 images are clearly overestimated. However, it has been shown that for YSC systems close enough the mean effective radii do readily fall within the range of GC radii (Meurer \\etal 1995). Our focus in this paper is the luminosity and colour evolution of the YSC population in the Antennae and, in particular, the future evolution of the YSC's LF.  In a previous paper, we model the evolution of star clusters for different initial metallicities in terms of broad band colours and stellar metallicity indices. We find important colour differences for clusters of various metallicities, already at young ages, and showed that once the stellar metallicity is known, rather precise age dating becomes possible. Comparison with young star clusters in NGC 7252 (Whitmore \\etal 1993), the two brightest of which have spectroscopy available (Schweizer \\& Seitzer 1993), confirmed a metallicity of ${\\rm Z \\sim \\half Z_{\\odot}}$ predicted from our global starburst modelling in this Sc $-$ Sc merger remnant. The mean age of the young star cluster population was shown to agree well with the global burst age of $\\sim 1.3$ Gyr, and ages derived from solar metallicity models would differ by a factor $\\sim 2$ (see Fritze $-$ v. Alvensleben \\& Burkert 1995 for details).  \\medskip\\noindent Observationally, the best case by now to study the LF of YSCs are the Antennae with more than 700 young star clusters detected by Whitmore \\& Schweizer (1995, hereafter {\\bf WS95}), a number large enough to allow for a statistical analysis. In this paper, we will examine the LF of the young star cluster system in the Antennae. It seems clear that not all bright knots in the NGC 4038/39 system with its still ongoing starburst will probably be GCs, in particular those with large effective radii R$_{\\rm eff}$ might rather be open clusters or associations. Therefore, after age dating the clusters in Sect. 2., we subdivide Whitmore \\& Schweizer's young star cluster sample into two subsamples containing the small knots and the more extended systems, respectively (Sect. 3.). In a first step, we assume a uniform age for the YSC population and we model the evolution of the YSCs' LF over a Hubble time and compare to LFs of the Milky Way's and other nearby galaxies' GCSs (Sect.4.). In an ongoing starburst like in the Antennae, the age spread among the YSCs may not be negligible (see also Meurer 1995). To examine the age spread effects on the LF we determe individual ages for all star clusters from their (V-I) colour and discuss the star clusters' age distribution in Sect. 5. We calculate the resulting individual fading for all clusters in Sect. 6. Alternative possibilities to subdivide the YSC sample and their consequences are discussed in Sect. 7. The of a young GCS may not only change by fading but also by dynamical effects as e.g. stellar mass loss within the cluster and/or tidal interaction of a cluster with the galactic potential. For GC populations in non-interacting galaxies, these effects were studied by Chernoff \\& Weinberg (1990), their results are largely confirmed by the independent and more realistic approach of Fukushige \\& Heggie (1995). In a recent paper Vesperini (1997) shows that in the Milky Way potential an initial log-normal mas distribution represents a quasi-equilibrium state that allows to preserve both its shape and parameters during a Hubble time of dynamical evolution, even though up to 70 \\% of the initial cluster population get disrupted. In case of the Antennae, i.e. in a still uncompleted merger with its gravitational potential being highly variable both in space and in time, however, external dynamical effects seem extremely difficult to model. Referring YSC luminosities to a common age allows to recover the mass function of the YSC system when combined with model M/L. We discuss the possible influence of dynamical effects in Sect. 8. and point out the possibility to observationally approach these dynamical effects by intercomparing star cluster populations in interacting galaxies and merger remnants of various ages. Sect. 9. summarizes our conclusions. The spatial distribution of the YSCs $-$ and of their properties as derived here $-$ will be discussed in a forthcoming paper. ",
        "conclusions": "Using our method of evolutionary synthesis for various metallicities we present a first analysis of WS95's WFPC1 data on bright star clusters in the ongoing merger-induced starburst in NGC 4038/39. Assuming a metallicity Z $\\sim 0.01$ on the basis of the progenitor spirals' ISM properties and applying a uniform reddening as given by WS95 we age-date the bright cluster population from their (V$-$I) colors and, as far as available, also from their (U$-$V). It turns out that in addition to a large population of young clusters with a mean age of $2 \\cdot 10^8$ yr (consistent with  the dynamical time since pericenter) part of the original spirals' old GC population is also observed. A key question with far-reaching consequences as to the origin of elliptical galaxies is whether there are a significant fraction of young GCs among the YSC population. Two basic properties discriminate open clusters/OB associations from GCs in our Galaxy and others: the concentration parameter c = log (${\\rm R_T/R_{eff}}$) and the LF which, in contrast to that for an open cluster system, is Gaussian for {\\bf old} GCSs. Tidal radii and, consequently, concentration parameters not being accessible to observations in distant galaxies we examine the LFs of cluster subsamples with large and small effective radii.  In a first step, using a common mean age for all young clusters and a corresponding uniform fading to an age of $\\sim 12$ Gyr we find that while the LF for extended clusters at 12 Gyr is definitely not Gaussian, that for the low R$_{{\\rm eff}}$ clusters may well contain a Gaussian  (= GC) subcomponent together with a strong overpopulation of the faint bins, which themselves, however, might be expected to be severely depopulated over a Hubble time by dynamical effects not included in our models.  Since for an ongoing starburst the age spread among YSCs may be of the same order as their ages, age spread effects are expected to reshape the LF. Clusters from the bright end tend to be younger on average and fade more than clusters from the faint end. We therefore, in a second step, model the individual fading consistent with individual ages of the YSCs as derived from their ${\\rm (V-I)}$ and ${\\rm (U-V)}$ colours, and we follow the LF changing its shape over a Hubble time.  Surprisingly, accounting for these age spread effcets, we find the final LFs of large {\\bf and} small R$_{\\rm{eff}}$ cluster subsamples not to be significantly different any more. Instead, the LF of {\\bf all} YSCs evolved to a common age of 12 Gyr is well compatible with a ``normal'' GCLF. Its turn-over occurs at $\\langle {\\rm M_{V_0}} \\rangle \\sim -6.9$ mag, i.e. slightly fainter than the average value $\\langle {\\rm M_{V_0}} \\rangle \\sim -7.1$ mag for 16 galaxies. This difference is readily explained in terms of a higher metallicity of the secondary cluster population.  The number of old GCs from the spiral progenitors is consistent with the number of bright GCs expected if the progenitors had GCSs similar to the ones in the Milky Way and M31.  Strikingly, neither the mean nor the distribution of effective radii is significantly different for the old GC sample and for the YSC sample. On the basis of these WFPC1 data we tentatively conclude that the bulk of the YSC population detected in the Antennae might well be young GCs and that the open clusters/associations probably also present among the YSCs do not seem to systematically differ from young GCs in terms of R$_{\\rm{eff}}$. We are looking foreward to repeat this kind of analysis on WFPC2 data which may reach close to the old GCS's turn-over, reveal a number of fainter young objects, and will allow for more precise and definite conclusions.  Dynamical effects that eventually might further reshape the LF over a Hubble time are discussed. Referring the YSCs' luminosities to a uniform age allows to recover the intrinsic mass function of the YSC system. This mass function seems to be log-normal which, according to Vesperini (1997), represents a quasi-equilibrium distribution that is going to be preserved in shape though not in number of clusters over a Hubble time of dynamical evolution.  Dynamical effects, however, are extremely difficult to model in detail in an ongoing merger. Comparison of YSC populations in mergers/starbursts of various ages seems a promising tool in an attempt to understand these effects from an observational side.    \\vskip 1 cm  {\\sl Acknowledgements.} I am deeply indebted to B. Whitmore \\& F. Schweizer for valuable discussions, encouragement and for sending us their star cluster data in machine readable form. I am grateful to Ken Freeman, Tom Richtler, and Andreas Burkert for interesting discussions on dynamical aspects. I wish to thank Prof. Appenzeller and all the collegues from the Landessternwarte Heidelberg for their warm hospitality during a 3 months stay, when this projected was begun. My deep thanks go to the referee, G. Meurer, for his very detailed and constructive suggestions that greatly improved the paper. I gratefully acknowledge financial support from the SFB Galaxienentwicklung in Heidelberg and through a Habilitationsstipendium from the Deutsche Forschungsgemeinschaft under grant Fr 916/2-1 in G\\\"ottingen."
    },
    "9803/astro-ph9803162_arXiv.txt": {
        "abstract": "We have established a model to systematically estimate the contribution of the mid-infrared emission features between 3 $\\mu$m and 11.6 $\\mu$m to the IRAS in-band fluxes, using the results of ISO PHT-S observation of 16 galaxies by Lu et al. (1997).  The model is used to estimate more properly the $k$-corrections for calculating the restframe 12 and 25 $\\mu$m fluxes and luminosities of IRAS galaxies.  We have studied the 12-25 $\\mu$m color-luminosity relation for a sample of galaxies selected at 25 $\\mu$m. The color is found to correlate well with the 25 $\\mu$m luminosity, the mid-infrared luminosity, and the ratio of far-infrared and the blue luminosities.  The relations with the mid-infrared luminosities are more sensitive to different populations of galaxies, while a single relation of the 12-25 $\\mu$m color vs. the ratio of the far-infrared and the blue luminosities applies equally well to these different populations. The luminous and ultraluminous infrared galaxies have redder 12-25 $\\mu$m colors than those of the quasars.  These relations provide powerful tools to differentiate different populations of galaxies.  The local luminosity function at 12 $\\mu$m provides the basis for interpreting the results of deep mid-infrared surveys planned or in progress with ISO, WIRE and SIRTF.  We have selected a sample of 668 galaxies from the IRAS Faint Source Survey flux-density limited at 200 mJy at 12 $\\mu$m.  A 12 $\\mu$m local luminosity function is derived and, for the first time in the literature, effects of density variation in the local universe are considered and corrected in the calculation of the 12 $\\mu$m luminosity function.  It is also found that the 12 $\\mu$m-selected sample are dominated by quasars and active galaxies, which therefore strongly affect the 12 $\\mu$m luminosity function at high luminosities.  The ultraluminous infrared galaxies are relatively rare at 12 $\\mu$m comparing with a 25 $\\mu$m sample. ",
        "introduction": "\\label{sec:intro}  The mid-infrared (MIR) spectral region is well-suited for studying starburst and ultraluminous galaxies.  About 40\\% of the luminosity from starburst galaxies is radiated from 8-40 $\\mu$m (\\markcite{soi87}Soifer et al.  1987).  Extinction effects are small, and infrared cirrus emission is reduced at these wavelengths relative to far-infrared bands. For a fixed telescope aperture, the spatial resolution is also higher at shorter wavelengths, and the confusion limit lies at higher redshifts. All the recent and near-future infrared space missions, such as the {\\it Infrared Space Observatory (ISO)}, the {\\it Wide-Field Infrared Explorer (WIRE)}, and the {\\it Space Infrared Telescope Facility (SIRTF)} will conduct surveys in mid-infrared bands.  {\\it WIRE}, a Small Explorer mission due to launch in late 1998 (\\markcite{hac96}Hacking et al.\\ 1996; \\markcite{schemb96}Schember et al. 1996), will conduct a very deep survey at 12 and 24 $\\mu$m to study the evolution of starburst galaxies. To interpret the results of these surveys now in progress or soon to commence, it is necessary to better understand the mid-infrared properties of galaxies in the local Universe.  One of the most important tools for extracting the rate and type of galaxy evolution from a mid-infrared survey is the faint source counts. A local mid-infrared luminosity function is the basis for calculating the mid-infrared faint source counts incorporating different evolutionary scenarios, and to extract the evolution by comparing with observations.  Mid-infrared luminosity functions have been calculated at 12 and 25 $\\mu$m (Soifer \\& Neugebauer 1991) using a 60 $\\mu$m selected IRAS sample (Soifer et al. 1987).  More recently, \\markcite{rush93}Rush et al. (1993) selected a 12 $\\mu$m flux-limited sample and calculated the luminosity functions for Seyfert and non-Seyfert galaxies in the sample.  In a previous paper (\\markcite{paper1}Shupe et al. (1997), Paper I hereafter), we have presented the results of a 25 $\\mu$m luminosity function calculated from a large flux-limited IRAS sample containing 1456 galaxies.  We continue to select a flux-limited sample and calculate the luminosity function at 12 $\\mu$m in this paper.  The relation between the mid-infrared color and luminosity plays another important role in estimating various properties of galaxy evolution.  It defines distinct regions in the color-flux diagram, for example, for different types of evolution and for different populations of galaxies.  Such a relation is indicated by the 12-25 $\\mu$m vs. 60-100 $\\mu$m color-color relation or by the 12-25 $\\mu$m color vs. the far-infrared luminosity relation obtained from the IRAS survey (\\markcite{soi91}Soifer \\& Neugebauer 1991), and can be estimated from a large sample of galaxies selected at mid-infrared bands.  Emission features near 12 $\\mu$m thought to be produced by aromatic hydrocarbon molecules have been observed in many astronomical spectra (e.g., \\markcite{gill73}Gillett et al. 1973; \\markcite{russ78}Russell et al. 1978; \\markcite{sell84}Sellgren 1984; \\markcite{roch91}Roche, Aitken, \\& Smith 1991; \\markcite{oboul96}Boulade et al. 1996; \\markcite{vig96} Vigroux et al. 1996; \\markcite{met96}Metcalfe et al. 1996; \\markcite{ces96}Cesarsky et al. 1996; \\markcite{lu97}Lu et al. 1997).  These broad emission features complicate the calculations of $k$-corrections and the fluxes and luminosities at mid-infrared bands.  Fortunately, ISO observations have resulted in high-quality mid-infrared spectra in various astronomical circumstances, and the on-going surveys of IRAS galaxies using ISO can provide an especially useful handle on this problem.  In the next section we present a model to systematically calculate the contribution of the emission features in the mid-infrared bands of IRAS galaxies.  The model is then incorporated in the following sections. Section \\ref{sec:clrlum} discusses the mid-infrared color-luminosity relation obtained from the large 25 $\\mu$m-selected sample of Paper I. The population-dependency of the relation is discussed.  Then we present the calculation of the 12 $\\mu$m luminosity function in Section \\ref{sec:lumfcn}.  We first discuss a selection of galaxy sample flux-limited at 12 $\\mu$m from the IRAS Faint Source Survey in Section \\ref{sec:sample}.  Then the luminosity function is derived and corrected for density variations in Section \\ref{sec:pahlf}. In Section \\ref{sec:nopahlf} we discuss the effects of active galaxies and quasars on the 12 $\\mu$m luminosity function. We summarize our results in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Our main results are summarized as follows:  1. Quasars and Seyfert galaxies dominate the high luminosity regime in a 12 $\\mu$m flux-limited sample.  The ultraluminous infrared galaxies are relatively rare at 12 $\\mu$m (contrast to a 25 $\\mu$m sample).  2. We have a technique for differentiating between quasars and ultraluminous infrared galaxies using their 12-25 $\\mu$m color (see Figures 5-9). Qualitatively, quasars are bluer than the luminous and ultraluminous infrared galaxies at high 25 $\\mu$m luminosities.  The ultraluminous infrared galaxies also have greater far-infrared to blue luminosity ratio on average than those of the other populations.  3. A highly complete sample flux-density limited at 200 mJy at 12 $\\mu$m selected from the Faint Source Survey catalogs is used to calculate a local 12 $\\mu$m luminosity function, which is then corrected for density-variation.  We are establishing a library of galaxy SEDs as a function of luminosity for more accurate $k$-corrections, and will discuss the faint source counts based on our 12 and 25 $\\mu$m luminosity functions in a forthcoming paper (Xu et al. 1997)."
    },
    "9803/astro-ph9803024_arXiv.txt": {
        "abstract": "\\rxj\\ is an unidentified bright soft \\Xray\\ source which shows pulsations at a 8.39\\,s period and has a thermal spectrum.  We present deep B and R band images of its \\Xray\\ localization.  We find one possible counterpart in the \\Xray\\ error box, with magnitudes $B=26.6\\pm0.2$ and $R=26.9\\pm0.3$.  The very high X-ray to optical flux ratio confirms that this object is an isolated neutron star.  We discuss possible models and conclude that only two are consistent with the data and at the same time are able to draw from a large enough population to make finding one nearby likely.  In our opinion the second criterion provides a stringent constraint but appears to have been ignored so far.  The first model, suggested earlier, is that \\rxj\\ is a  weakly magnetized neutron star accreting from the interstellar medium.  The second is that it is a relatively young, highly magnetized neutron star, a ``magnetar'', which is kept hot by magnetic field decay. ",
        "introduction": "}  The population of defunct radio pulsars far exceeds that of active ones.  It is believed that the Galaxy has about $2\\,10^5$ radio pulsars.  The neutron-star birth rate is estimated to be between one per 30 yr to one per 100 yr.  Assuming a constant pulsar production rate and an age of the disk of $10^{10}$\\,yr, one infers a Galactic neutron star population of $\\sim\\!2\\,10^8$ -- three orders of magnitude larger than that of the active radio pulsar population.  It is not easy to detect old neutron stars.  While the nearest few intermediate-age pulsars can be identified by their cooling radiation, which peaks in the soft \\Xray/EUV band, the defunct pulsars will have become too cool to be observable.  A small fraction of them, however, may be in a position to accrete matter from the interstellar medium.  These will then get reheated and reappear in the \\Xray\\ sky.  Quite independent of this discussion there has been a growing recognition of a population of highly magnetized neutron stars.  The circumstantial evidence for this class comes from studies of soft gamma-ray repeaters (SGRs) and long-period pulsars in supernova remnants (Vasisht \\& Gotthelf \\cite{vasig:97}).  Thompson \\& Duncan (\\cite{thomd:95}) have introduced the term ``magnetars'' for neutron stars with field strengths significantly larger than $10^{12}$\\,G, the typical field strength inferred for radio and \\Xray\\ pulsars.  The birthrate of SGRs has been estimated to be roughly 10\\% that of the ordinary pulsars (Kulkarni \\& Frail \\cite{kulkf:93}; Kouveliotou et al.\\ \\cite{kouv&a:94}).  The relevance of magnetars to the discussion at hand is as follows. Unlike the situation for ordinary neutron stars, magnetic field decay is expected to be significant in highly magnetized neutron stars. This decay could reheat the magnetar (Thompson \\& Duncan \\cite{thomd:96}), making it hotter than an ordinary neutron star, and thus brighter in soft X rays.  Two of the best candidates for this general class of neutron stars have emerged from the ROSAT mission: \\rxjw\\ (Walter, Wolk, \\& Neuh\\\"auser \\cite{waltwn:96}) and \\rxj\\ (Haberl et al.\\ \\cite{habe&a:97}). Both are bright ROSAT objects with very soft \\Xray\\ spectra.  Walter \\& Matthews (\\cite{waltm:97}) have provided compelling evidence for the identification of a faint blue optical counterpart of \\rxjw.  In this {\\em Letter}, we present deep B and R observations of the localization of \\rxj. ",
        "conclusions": ""
    },
    "9803/astro-ph9803268_arXiv.txt": {
        "abstract": "In the context of open inflation, we calculate the probability distribution for the density parameter $\\Omega$. A large class of two field models of open inflation do not lead to infinite open universes, but to an ensemble of inflating islands of finite size, or ``quasi-open'' universes, where the density parameter takes a range of values. Assuming we are typical observers, the models make definite predictions for the value $\\Omega$ we are most likely to observe. When compared with observations, these predictions can be used to constrain the parameters of the models. We also argue that obsevers should not be surprised to find themselves living at the time when curvature is about to dominate. ",
        "introduction": "Anthropic considerations have often been used in order to justify the ``naturalness'' of the values taken by certain constants of Nature \\cite{anthropic}. In these approaches, it is assumed that the ``constants'' are really random variables whose range and ``a priori'' probabilities are determined by the laws of Physics. Knowledge of these ``a priori'' probabilities is certainly useful, but not sufficient to determine the probability for an observer to measure given values of the constants. For instance, some values which are in the ``a priori'' allowed range may be incompatible with the very existence of observers, and in this case they will never be measured. The relevant question is then how to assign a weight to this selection effect.  A natural framework where these ideas can be applied is inflation. There, the false-vacuum energy of the scalar field which drives the inflationary phase can thermalize in different local minima of its potential, and each local minimum may have a different set of values for the constants of Nature. Also, there may be different routes from false vacuum to a given minimum. In this case all thermalized regions will have the same low energy Physics constants, but each route will yield a hot universe with different large scale properties. Here, we shall be concerned with this possibility, where the fundamental constants (such as the gauge couplings or the cosmological constant) are fixed, but other cosmological parameters such as the density parameter or the amplitude of cosmological perturbations are random variables whose distribution is dynamically determined.  In this context, the most reasonable -and predictive- version of the anthropic principle seems to be the principle of mediocrity \\cite{medi,gott}, according to which we are typical observers who shall observe what the vast majority of observers would. Thus, the measure of probability for a given set of constants is simply proportional to the total number of civilizations emerging with those values of the constants. In this paper we shall use this principle in order to calculate the probability distribution for the density parameter $\\Omega$.  Standard inflationary models predict $\\Omega=1$ with ``certainty''. What this means is that these models can explain the observed homogeneity and isotropy of the universe only if the universe is flat. However, a class of ``open inflation'' models which lead to $\\Omega<1$ have received some attention in recent years \\cite{open,BGT,LM}. In these models, inflation proceeds in two steps. One starts with a scalar field $\\sigma$ trapped in a metastable minimum of its potential $V(\\sigma)$. The false vacuum energy drives an initial period of exponential expansion, and decays through quantum nucleation of highly symmetric bubbles of true vacuum. The interior of these bubbles has the geometry of an open Friedmann-Robertson-Walker universe. This accounts for the observed homogeneity and isotropy of the universe. In order to solve the flatness problem a second stage of slow roll inflation inside the bubble is necessary.  In models with a single scalar field $\\sigma$, all bubbles have the same value of $\\Omega$ which is determined by the number of e-foldings in the second period of inflation. The potential $V(\\sigma)$ in such models is assumed to have a rather special form, with a sharp barrier next to a flat slow-roll region, which requires a substantial amount of fine-tuning. Additional tunning is needed to arrange the desired value of $\\Omega$. A more natural class of models includes two fields, $\\sigma$ and $\\phi$, with $\\sigma$ doing the tunneling and $\\phi$ the slow roll \\cite{LM}.  The simplest example is \\begin{equation} \\label{coupled} V(\\sigma,\\phi)=V_t(\\sigma) + {g \\over 2} \\sigma^2 \\phi^2, \\end{equation} where $V_0(\\sigma)$ has a metastable false vacuum at $\\sigma=0$. After $\\sigma$ tunnels to its true minimum $\\sigma=v$, the field $\\phi$ would drive a second period of slow roll inflation inside the bubble. Depending on the value of $\\phi$ at the time of nucleation, the number of e-foldings of the second stage of inflation would be different.  Initially, it was believed \\cite{LM} that models such as (\\ref{coupled}) would yield an ensemble of infinite open universes, one inside each nucleated bubble, and each one with a different value of the density parameter. However, it has been recently realized \\cite{GGM} that this picture is oversimplified.  The two field models which allow for variable $\\Omega$ do not actually lead to infinite open universes, but to an ensemble of inflating islands of finite size inside of each bubble. These islands are called quasi-open universes. Within each island, the number of e-foldings of inflation decreases as we move from the center to the edges. Also, each island is characterized by a different number of e-foldings in its central region. As a result, even within the same bubble, different observers will measure a range of values of the density parameter. The picture of the large scale structure of the universe in these models is rather simple, because all bubbles have the same statistical properties.  We shall see that the quasiopen nature of inflation is of crucial importance for the calculation of the probability distribution for the density parameter.  In models of quasiopen inflation, such as (\\ref{coupled}), $\\Omega$ takes different values in different parts of the universe, while the other constants of Nature and cosmological parameters remain fixed. More general models can be constructed where other parameters can change as well, and in Section VII we give an example of a model with a variable amplitude of density fluctuations.  However, our main focus in this paper is on the models in which only $\\Omega$ is allowed to vary.  In order to apply the principle of mediocrity to our models, we will have to compare the number of civilizations in parts of the universe with different values of $\\Omega$. Of course, we cannot calculate the number of civilizations.  However, since the value of $\\Omega$ does not affect the physical precesses involved in the evolution of life, this number must be proportional to the number of habitable stars or, as a rough approximation, to the number of galaxies.  Hence, we shall set the probability for us to observe a certain value of $\\Omega$ to be proportional to the number of galaxies formed in parts of the universe where $\\Omega$ takes the specified value.  The principle of mediocrity was applied to calculate the probability distribution for $\\Omega$ in an earlier paper \\cite{VW}, which assumed the old picture of homogeneous open universes inside bubbles.  A serious difficulty encountered in that calculation was that open universes inside the bubbles have infinite volume and contain an infinite number of galaxies.  Thus, to find the relative probability for different values of $\\Omega$, one had to compare infinities, which is an inherently ambiguous task. This problem was addressed in \\cite{VW} by introducing a cutoff and counting only galaxies formed prior to the cutoff.  Although the cutoff procedure employed in \\cite{VW} has some nice properties, it is not unique, and the resulting probability distribution is sensitive to the choice of cutoff \\cite{linde}.  This cutoff dependence, which also appears in other models of eternal inflation \\cite{linde,vireg}, has lead some authors to doubt that a meaningful definition of probabilities in such models is even in principle possible \\cite{linde,gbl}.  However, this pessimistic conclusion may have been premature. According to the quasiopen picture, $\\Omega$ takes all its possible values within each bubble.  Since all bubbles are statistically equivalent, it is sufficient to consider a single bubble.  Moreover, we can restrict ourselves to a finite (but very large) comoving volume within that bubble, provided that its size is much greater than the characteristic scale of variation of $\\Omega$.  Thus, we no longer need to compare infinities, and the problem becomes well defined.  The possibility of unambiguous calculation of probabilities in the quasiopen model was our main motivation for revising the analysis of Ref.\\cite{VW}.  Also, we shall give a more careful treatment of the astrophysical aspects of the problem which were discussed rather sketchily in \\cite{VW}.  The paper is organized as follows. In Section II we review the main features of quasi-open inflation. In Section III we introduce the probability distribution for $\\Omega$. A basic ingredient in this distribution will be the anthropic factor $\\nu(\\Omega)$, which gives the number of civilizations that develop per unit thermalized volume in a region characterized by a certain value of $\\Omega$. In Section IV we evaluate $\\nu(\\Omega)$ and calculate the probability distribution for $\\Omega$ in the model (\\ref{coupled}). In Section V we extend our results to more general models with arbitrary slow roll potentials for the field $\\phi$. In Section VI we discuss observational constraints on quasiopen models due to CMB anisotropies and how these constraints restrict the class of models that give a probability distribution peaked at a non-trivial value of $\\Omega$. In Section VII we comment on the ``cosmic age coincidence'', that is, on whether it would be surprising to find ourselves living at the time when the curvature of the universe starts dominating. In Section VIII we summarize our conclusions.  Some side issues and technical details are discussed in the appendices. ",
        "conclusions": "We have calculated the probability distribution for the density parameter in models of open inflation with variable $\\Omega$. This probability is basically the product of three factors: the ``tunneling'' factor, which is related to the microphysics of bubble nucleation and subsequent expansion; the volume factor, related to the amount of slow roll inflation undergone in different regions of the universe; and the ``anthropic factor'', which determines the number of galaxies that will develop per unit thermalized volume.  It is interesting that the expression for the probability (\\ref{distributiony}) depends on the underlying particle physics model through a single dimensionless parameter $\\mu$, defined in Eq.(\\ref{defmu}).  Taking the minimum of the slow roll potential to be at $\\phi=0$, the tunneling factor tends to suppress large initial values of $\\phi$, favouring low values of $\\Omega$. However, only those regions for which $\\phi$ is large enough will inflate. Hence, there will be a competition between volume enhancement and ``tunneling'' suppression.    The most interesting situation occurs when the tunneling suppression dominates over the volume factor. In this case, the product of both would peak at $\\Omega=0$, and the anthropic factor $\\nu(\\Omega)$ becomes essential in determining the probability distribution. In an open universe, cosmological perturbations stop growing when the universe becomes curvature dominated, and for low values of $\\Omega$ structure formation is suppressed.  The effect of the anthropic factor is, therefore, to shift the peak of the distribution from $\\Omega=0$ to a nonzero value of $\\Omega$.   As a first approximation \\cite{VW,MSW}, we have taken $\\nu(\\Omega)$ to be proportional to the fraction of matter that clusters on the galactic mass scale in the entire history of a certain region. We have found that the peak of the distribution is given by the condition \\begin{equation} \\kappa \\left({1-\\Omega \\over \\Omega}\\right)_{peak} \\approx \\left({3\\over 2}\\mu-{5\\over 4}\\right)^{1/2}, \\label{mon} \\end{equation} where the coefficient $\\kappa \\sim 10^{-1}$ is defined in (\\ref{kappa}). For models with $\\mu \\sim 1$ (which can be easily constructed), the probablility distribution for the density parameter ${\\cal P}(\\Omega)$ can peak at values of $\\Omega$ such that $x=(1-\\Omega)/\\Omega\\sim 1$ (See Fig. 1). The peaks are not too sharp, with amplitude $\\Delta y \\approx 1/2$, or $\\Delta x \\approx 5$, so a range of values of $\\Omega$ would be measured by typical observers.  The analysis we presented here demonstrates that, given a particle physics model, the probability distribution for $\\Omega$ can be unambiguously calculated from first principles.  We can also invert this approach and use our results to exclude particle physics models which give the peak of the distribution at unacceptably low values of $\\Omega$. This gives the constraint $\\mu\\lesssim 3$.  An independent constraint on the model parameters can be obtained from CMB observations.  If the observed CMB anisotropies are to be explained within the same two-field model of open inflation, without adding any extra fields, then we have shown in Section IV that the corresponding constraint (if the observed value of $\\Omega$ lies in the range $.1$ to $.7$) is $\\mu \\gtrsim  10^{6} \\epsilon^2$, where $\\epsilon$ is the slow roll parameter defined in (\\ref{epsilon}). Combinig both constraints, we obtain a bound on the slow roll parameter $$ \\epsilon\\lesssim 10^{-3}. $$ This bound is somewhat restrictive. For instance, for the simple free field model (\\ref{coupled}), the slow roll parameter is of order $10^{-2}$, and so this model would contradict observations. It is easy, however, to generalize the slow roll potential in order to make $\\epsilon$ sufficiently small. If one allows some other source for CMB fluctuations (e.g., topological defects), then the CMB constraint is much less restrictive, and simple models of the form (\\ref{coupled}) are still viable.  We have advanced anthropic arguments towards explaining the ``cosmic age coincidence'', that is, whether it would be surprising to find that we live at the time when the curvature is about to dominate. We have argued that this is not unexpected. We have also discussed a three-field model in which the amplitude of density fluctuations $Q$ becomes a random variable.  We have outlined an argument explaining the observed value $Q\\sim 10^{-5}$ in the framework of this model.  While this work was being completed, Hawking and Turok \\cite{HT98}, have suggested the possibility of creation of an open universe from nothing (see also \\cite{everybody}). The validity of the instantons describing this process \\cite{alex}, and also their ability to successfully reproduce a sufficiently homogeneous universe, is still a matter of debate and needs further investigation. Clearly, the analysis presented in this paper can be easily adapted to this new framework."
    },
    "9803/astro-ph9803118_arXiv.txt": {
        "abstract": "We use the two-point correlation function to calculate the clustering properties of the recently completed SSRS2 survey, which probes two well separated regions of the sky, allowing one to evaluate the sensitivity of sample-to-sample variations. Taking advantage of the large number of galaxies in the combined sample, we also investigate the dependence of clustering on the internal properties of galaxies.  The redshift space correlation function for the combined magnitude-limited sample of the SSRS2 is given by $\\xi(s)=(s/5.85$ \\h1 Mpc)$^{-1.60}$ for separations between 2 $\\leq s \\leq$ 11 \\h1 Mpc, while our best estimate for the real space correlation function is $\\xi (r) = (r/5.36$ \\h1 Mpc)$^{-1.86}$. Both are comparable to previous measurements using surveys of optical galaxies over much larger and independent volumes. By comparing the correlation function calculated in redshift and real space we find that the redshift distortion on intermediate scales is small. This result implies that the observed redshift-space distribution of galaxies is close to that in  real space, and that $\\beta = \\Omega^{0.6}/b < 1$, where $\\Omega$ is the cosmological density parameter and  $b$ is the linear biasing factor for optical galaxies.   We have used the SSRS2 sample to study the dependence of $\\xi$ on the internal properties of galaxies such as luminosity, morphology and color.  We confirm earlier results that luminous galaxies ($L>L^*$) are more clustered than sub-$L^*$ galaxies and that the luminosity segregation is scale-independent. We also find that early types are more clustered than late types. However, in the absence of rich clusters, the relative bias between early and late types in real space, $b_{E+S0}$/$b_S$ $\\sim$ 1.2, is not as strong as previously estimated. Furthermore, both morphologies present a luminosity-dependent bias, with the early types showing a slightly stronger dependence on the luminosity. We also find that red galaxies are significantly more clustered than blue ones, with a mean relative bias of $b_R/b_B$ $\\sim$ 1.4, stronger than that seen for morphology. Finally, by comparing our results with the measurements obtained from the infrared-selected galaxies we determine that the relative bias between optical and \\iras galaxies in real space is $b_o/b_I$ $\\sim$ 1.4. ",
        "introduction": "\\subsection{Method}   The two-point correlation function $\\xi (r)$ can be computed from the data using the estimator suggested by Hamilton (1993): \\begin{equation} \\xi(r) = {DD(r) RR(r) \\over [DR(r)]^2}  - 1, \\end{equation} where $DD(r)$, $RR(r)$ and  $DR(r)$ are the number of data--data, random--random and data--random pairs, with separations in the interval between $r$ and $r + dr$.  The random catalog is generated using the same selection criteria as the galaxy sample. This estimator has the advantage that it is not too sensitive to uncertainties in the mean density, which is only a second order effect. The counts $DD(r)$, $DR(r)$ and $RR(r)$ can be generalized to include a weight $w$ which is particularly important to correct for selection effects at large distances in magnitude-limited samples: \\begin{eqnarray} DD(r) = &{\\displaystyle{ \\sum_i^{N_{\\scriptscriptstyle{gal}}} \\sum_j^{N_{\\scriptscriptstyle{gal}}}} } & w(s_j, r) w(s_i, r), \\\\ & \\scriptscriptstyle {r- \\Delta r \\leq |s_i-s_j| \\leq r+ \\Delta r} \\nonumber \\end{eqnarray} where $i$ sums over all objects in the sample and the sum over $j$ includes all particles at a distance $s$ from the origin, which in this work is taken as the centroid of the Local Group, and  $r  = | {\\bf{s}}_i -{\\bf{s}}_j|$ is the separation of the pair $(i,j)$. The galaxy-random pairs $DR(r)$ and random-random pairs $RR(r)$ are similarly weighted. The most common weighting schemes are:  equally weighted pairs $w(s_i, r)=1$; equally--weighted volumes where $w(s_i, r)=1/\\phi(s_i)$ and the minimum variance weighting given by \\begin{equation} w (s_i, r) = {1 \\over 1 +4\\pi \\bar n J_3(r) \\phi(s_i)}, \\quad J_3 (r) = \\int_0^r dr'r'^2 \\xi(r'), \\end{equation} where $\\phi (s_i)$ is the selection function at distance $s_i$ from the origin and $J_3$ is the mean number of excess galaxies out to a distance $r$ around each galaxy. Even though in the last scheme the weights depend on the unknown correlation function, in practice, it is not very sensitive to the exact form of $\\xi(r)$ (\\eg Loveday \\etal 1992; Marzke, Huchra \\& Geller 1994). In this work we adopt the minimum-variance weighting and take  $J_3 (r$ = 30 \\h1 Mpc) $\\sim 1100$, obtained by using the real-space correlation function of Davis \\& Peebles (1983). The mean densities were calculated using the estimator \\begin{equation} \\bar n = \\sum_{i=1}^{N_{gal}}w_i/\\int_{s_{min}}^{s_{max}} dV\\phi(s)w(s) \\end{equation} where again  $\\phi(s)$ is the selection function, derived from the luminosity function and $w(s)$ is the weight (\\eg Davis \\& Huchra 1982). The errors for the redshift space correlation function (\\xis) as well as for the real-space correlation discussed below, were calculated by means of bootstrap resampling (Ling, Frenk \\& Barrow 1986). For the volume-limited samples the total of bootstraps was 50, while for magnitude-limited samples 25 resamplings were calculated. As shown by Fisher \\etal (1994), bootstraping tends to overestimate the true errors, so that the estimate of the latter will in general be rather conservative.  \\subsection{Real space}  In order to estimate real space correlation functions, we follow Davis \\& Peebles (1983). For any two galaxies with redshifts {\\bf s$_1$} and {\\bf s$_2$}, we define the separation in redshift space, and the separation perpendicular to the line of sight respectively as \\begin{equation} {\\bf {s = s_1 - s_2}}, \\quad {\\bf {l}} = {1 \\over 2} \\bf{(s_1 + s_2)}, \\end{equation} in the small angle approximation. From these parameters one can derive $\\pi$, the separation between two galaxies parallel to the line of sight and $r_p$, the separation perpendicular to the line of sight using: \\begin{equation} \\pi = {{\\bf s.l} \\over |l|}, \\quad r_p = \\sqrt{|{\\bf{s}}|^2 -\\pi^2}. \\end{equation} These are then used to compute the statistic $\\xi(r_p,\\pi)$ estimated from the pair-counts as \\begin{equation} 1 + \\xi(r_p,\\pi) = {DD(r_p,\\pi) RR(r_p,\\pi) \\over [DR(r_p,\\pi)]^2}. \\end{equation}  From \\xip we define the projected function: \\begin{equation} \\omega(r_p) = 2 \\int_0^\\infty d\\pi \\quad \\xi (r_p,\\pi), \\end{equation} which is related to the real space correlation function through \\begin{equation} \\omega(r_p) = 2\\int_0^\\infty dy \\quad \\xi[(r_p^2 + y^2)^{1/2}]. \\end{equation} The inverse is the Abel integral: \\begin{equation} \\xi(r) = -{1 \\over \\pi} \\int_r^\\infty dr_p {\\omega^\\prime(r_p) \\over (r_p^2-r^2)^{1/2}, } \\end{equation} where $\\omega^\\prime(r_p)$ is the first derivative of $\\omega (r_p)$. If the real space correlation function is a power-law, the integral for $\\omega(r_p)$ can be performed analytically to give \\begin{equation} \\omega(r_p) = r_p ({r_o \\over r_p})^{\\gamma} { \\Gamma({1 \\over 2}) \\Gamma({\\gamma -1 \\over 2}) \\over \\Gamma({\\gamma \\over 2})}. \\end{equation}  \\subsection {Biasing}  The variance of galaxy counts measures the clustering amplitude at intermediate scales. It is also a useful quantity to compare models and data.  The variance in the counts is defined as \\begin{equation} \\langle (N -nV)^2 \\rangle = nV +n^2V^2\\sigma^2, \\end{equation} where nV is the mean number of galaxies in the volume V and $ n^2V^2\\sigma^2$ is the mean number of galaxies in excess of random inside a sphere of volume V.  It is related to the moment of the correlation function (Peebles 1980) \\begin{equation} \\sigma^2 = {1 \\over V^2} \\int_VdV_1dV_2 \\xi(|r_1-r_2|), \\end{equation} which can be calculated numerically. For a power law correlation function $\\xi(r) = (r/r_o)^\\gamma$,  and a spherical volume of radius R we get \\begin{equation} \\sigma^2(R) = 72(r_o/R)^\\gamma/\\lbrack 2^\\gamma(3-\\gamma)(4-\\gamma)(6-\\gamma)\\rbrack. \\end{equation} This is the expression we have used to compute \\sig8 , and which is often used to normalize theoretical models.  The relative bias between two different samples at a given separation $s$ may be estimated through (\\eg Benoist et al. 1996) : \\begin{equation} \\frac {b} {b_*}(s) = \\sqrt{\\frac{\\xi(s)}{\\xi_*(s)}} = \\sqrt{\\frac{J_3(s)}{J_3*(s)}}, \\end{equation} where the starred symbols denote a sample taken as a fiducial. The relative bias of the clustering may also be estimated through \\begin{equation} \\frac {b} {b_*}(s) = \\sqrt{\\frac{\\sigma^2(s)}{\\sigma_*^2(s)}}, \\end{equation} where $\\sigma^2$ is the variance of counts in cells described above.  These are the expressions used in this work to calculate the relative bias between galaxies of different luminosities relative to $L^*$ galaxies, as well as for different morphological types and colors.  \\section {Magnitude-limited Samples}  In order to estimate the effects due to the finite volume we are probing and to estimate the importance of cosmic variance, we compare the clustering properties of the individual SSRS2 south and north samples as well as the combined sample, with previous estimates of $\\xi$. In this analysis, we have computed \\xis taking into account all galaxies brighter than $M=-13$ in the velocity range $500 < v < 12,000$ \\kms. The correlation function was computed using the minimum-variance weighting discussed in Section 3 and a random background catalog of 10,000 points for the individual samples and 20,000 points for the combined sample. In the calculation of the selection function we have used the Schechter parameters determined for the entire SSRS2 survey by da Costa et al. (1997), which are $M^*$ = -19.55 and $\\alpha$ = -1.15. These values are virtually identical to those measured by da Costa et al. (1994) for the SSRS2 south.  We tested whether our results are affected by the presence of clusters of galaxies. For this we used a list of galaxy clusters with richness  R $\\geq$ 1 (J. Huchra, private communication). All galaxies whose positions were within one Abell radius of the central position of cluster, and that had radial velocities  within 500 \\kms of the cluster's mean radial velocity were culled from the sample. We find that the correlation parameters are virtually unchanged for the vast majority of the samples, and when there are changes, these are within the quoted errors of the complete sample. Therefore, we will not consider the removal of galaxies in clusters in this work.  The redshift space correlation function, \\xis, for the SSRS2 samples is shown in Fig. 1, where we plot the correlation function out to separations of 30 \\h1 Mpc. For the sake of clarity, in the figure we only show error bars calculated for the combined sample. One can see that beyond $\\sim$ 15 \\h1 Mpc, the errors become progressively larger, and sometimes the sample-to-sample variations are larger than the estimated errors. In general, \\xis is adequately described by a power-law on small scales. For most cases in this paper, the power-law fits were calculated in the interval $2 < s < 11$ \\h1 Mpc. The upper-limit was chosen because there is a suggestion of an abrupt change of slope in \\xis on scales $s$ $\\lsim$ 12 \\h1 Mpc. The lower-limit was chosen to minimize the effects on \\xis due to  peculiar motions of galaxies in virialized systems. The best power-law fits obtained for each sub-sample of the SSRS2 are represented as lines in Fig. 1, as explained in the caption. The correlation parameters derived from the fits are presented in Table 1, where we list:  the sample identification (column 1); the correlation length (column 2) and slope $\\gamma_s$ (column 3) obtained from the power-law fits; and in column (4) the rms variance in galaxy counts within spheres  8 \\h1 Mpc in radius, followed in columns (5) through (7) by the same parameters determined for real space, which will be discussed below.  An inspection of both Table 1 and Fig. 1 shows that the redshift correlation functions for SSRS2 sub-samples are very similar on small scales ($s < 10$ \\h1 Mpc). This also demonstrates that the sampling variations are consistent with the error estimates, at least in the range of separations for which the fits are calculated. In Fig. 2 we compare \\xis measured in this work for the combined SSRS2 with the \\xis measured in other surveys - the sparsely sampled Stromlo-APM survey (Loveday et al. 1995), the Las Campanas Redshift Survey (LCRS) (Tucker \\etal 1997), and that measured by Fisher et al. (1994) for the 1.2 Jy \\iras survey. The fit parameters calculated in these papers, as well as by other workers can be found in Table 2.  Despite small differences in amplitude, the shapes of the three optical surveys are remarkably similar. It is important to note that the volumes of the Stromlo-APM ( 2.5 $\\times$ 10$^6$ $h^{-3}$ Mpc$^3$) and the LCRS ( 2.6 $\\times$ 10$^6$ $h^{-3}$ Mpc$^3$) are about 5 times larger than that of the SSRS2 (5.2 $\\times$ 10$^5$ $h^{-3}$ Mpc$^3$), and probe different regions of space, and thus independent structures. The lower amplitude of the \\iras survey compared to the optical samples reflects the relative bias that exists between optically and infrared-selected galaxies, which will be further discussed in Section 6 below.  In Fig. 3, we compare our power-law fit parameters with equivalent measurements by other authors (see Table 2). In general, there is a good agreement between our values for the redshift space parameters and those obtained from other optical surveys, specially the Stromlo-APM and LCRS.  The effect of redshift distortions on the observed redshift correlation function has also been estimated for the SSRS2. These distortions are caused by the peculiar velocities of galaxies, which on large scales, are due to the infall of galaxies from low-density regions into high-density regions, while on small scales the correlations are smeared out by virial motions of galaxies in groups and clusters (\\eg Kaiser 1987). As described in Section 3, these effects may be accounted for by calculating the correlation function as a function of the separations parallel and perpendicular to the line of sight, which can then be used to define the \\wrp estimator, which is unaffected by redshift distortions. However, one should bear in mind that in general the calculation of the real space correlation function is much more susceptible to noise than that calculated in redshift space.  From the correlation functions \\xip computed using the minimum-variance weighting scheme, we have obtained \\wrp. From power-law fits, in the interval 2 $< r_p <$10 \\h1 Mpc,  we have derived the correlation parameters listed in Table 1. By comparing the real space fit parameters obtained in this work (Table 1) with previous measures (columns 4 and 5 in Table 2), we find a good agreement with the real space measurements of Davis \\& Peebles (1983), Loveday et al. (1995) and Marzke et al. (1995).  The fit to the real space correlation function for the combined sample is compared in Fig. 4 with the redshift space correlation function. One can see that at intermediate separations the redshift \\xis is amplified relative to the real space correlation \\xir. The small amplification suggests that the observed redshift distribution is close to the real space distribution. At separations of $\\sim$ 10 \\h1 Mpc, the ratio between the real and redshift space correlations is $\\sim$ 1.5. In the linear regime, peculiar motions on large scales cause \\xir to be amplified by a factor $\\sim$ $1 + { 1 \\over 2 } {\\beta} + { 1 \\over 5 } {\\beta^2}$ where $\\beta = {\\Omega^{0.6} \\over b}$ and $b$ is the linear biasing factor (Kaiser 1987). Therefore, a rough estimate for $\\beta$ is $ \\sim 0.6$, on scales of the order of 10 \\h1 Mpc, consistent with that determined by Loveday \\etal (1996).  \\section {The Clustering Dependence on the Internal Properties of Galaxies}  \\subsection {Luminosity}  In this work we use the combined SSRS2, as well as the SSRS2 north and south sub-samples to further explore the clustering dependence on luminosity, as was carried out by Benoist et al. (1996), but who only used the SSRS2 south. Probing independent structures in different volumes we can estimate the impact of cosmic variance. It should also be noted that the absolute magnitude limits considered in this section differ slightly from those of Benoist et al. (1996), and were chosen to compare our results with the volume-limited samples of Fisher \\etal (1994), which will be discussed in Section 6 below.  The volume-limited samples considered in this section only contain galaxies bright enough that would allow them to be included in the sample when placed at the cutoff distance. We defined samples limited at radial distances of 60, 80, 100 and 120 \\h1 Mpc. The absolute magnitude limits corresponding to these distances are  -18.39, -19.01 (both $L < L^*$) , -19.50 ($\\sim L^*$)  and -19.89 ($L > L^*$), respectively. For all galaxies in these samples, the weighting function is $w(r)=1$ and the volume densities  are simply the total number of galaxies divided by the corresponding volume.  The correlation functions obtained for the volume-limited sub-samples at the different depths are shown in Fig. 5 for $s \\leq 20$ \\h1 Mpc, where the different symbols represent different volume limits. For reasons of clarity, we only present error bars for the samples volume-limited at 60 \\h1 and 120 \\h1 Mpc. The meaning of these symbols, as well as the indication of the parent sample (SSRS2 south, north or combined) are shown in each panel. The power-law fits are represented by lines in the figure and the parameters are summarized in Table 3 where we list: in column (1) the sample; in column (2) the depth R; in column (3) the number of galaxies N$_g$; in column (4) the mean density; in columns (5) and (6) the power-law fit parameters and formal errors; and in column (7) $\\sigma_8$, the rms fluctuation of the number of galaxies in a sphere of radius 8\\h1 Mpc. The interval used in the fits is $ 2 < s < 11$ \\h1 Mpc, the same as that adopted in the previous section.  An inspection of Table 3 and Fig. 5 shows that the amplitude of \\xis tends to increase with the sample  depth, the variation being somewhat larger in the northern and combined samples.  We point out that \\xis for the SSRS2 north (panel b), is noisier because of the smaller number of galaxies, in most cases about half of those in the southern sample.  The correlation length ($s_0$) ranges from 3.8 \\h1 Mpc to 6.8 \\h1 Mpc. However, the slope varies considerably from sample to sample, though with a tendency of becoming steeper as the depth increases.  In order to evaluate the cosmic variance, we show in Fig. 6 \\xis for each of the volume-limited sub-samples, but  now plotting the results for the southern, northern and combined samples in each panel. For the samples in smaller volumes, the differences between the northern and southern samples are larger than the estimated error calculated for the combined sample, and probably reflect the amplitude of the sample to sample variation, with the north being systematically lower. For the larger volumes the samples present similar behavior, and the variations are generally consistent with the estimated errors.  To remove the effects of distortions due to motions, which may affect our estimates of the strength of clustering and the relative bias between different samples, we have also computed the real-space correlation function for the volume-limited samples. As above, we have computed \\xip for the sub-samples volume-limited at R = 60, 80, 100  and 120 \\h1 Mpc in each galactic hemisphere and for the combined sample. The resulting real space correlation parameters are listed in Table 4.  In Fig. 7 we compare \\xis  measured for each volume limit, denoted by open symbols, with the real-space correlation fits described above, represented as a solid line. For the sake of clarity, we only show the fits we measure for the combined sample, as this will be the one less affected by noise. The smearing due to motions in virialized systems for $r < 3$ \\h1 Mpc is quite noticeable for all samples, while the effect of peculiar motions is only obvious for the smaller volumes, little evidence being seen in the samples in larger volumes.  The dependence of clustering in redshift--space (as measured by $\\sigma_8$) with luminosity (as measured by the limiting absolute magnitude of each sample) is shown in Fig. 8(a), where we use as fiducial magnitude the value of $M^*$=-19.55 (see Section 4). The figure shows an overall behavior consistent with that found by Benoist et al. (1996) and which is detected in all samples, further demonstrating that this effect is unlikely to be spurious. This result supports their finding that there is a dependence of clustering on luminosity, as measured in redshift space. To further investigate its reality, we have computed \\xir in real space for the same volume limited samples. The results are shown in Fig. 8(b). Here again it is immediately apparent that the clustering amplitude increases with luminosity in the same way as seen in redshift space. On the whole, these results, using a larger sample, confirm in real space the conclusions of Benoist \\etal (1996).   In Fig. 9 we present the relative bias with scale calculated using equation (15), where we compare the the correlation function for the volume-limited samples at 60, 80 and 120 \\h1 Mpc relative to the 100 \\h1 Mpc sample. From the figure one may see that there are only minor differences between the smaller volumes. In the case of the sample volume-limited at 120 \\h1 Mpc, the relative bias is fairly constant over the range of scales we consider at $\\sim$ 1.5. This suggests that the luminosity bias is scale-independent, and that it starts to become important only for galaxies brighter than  $\\sim$ $L^*$.  \\subsection{Morphology}  Since all galaxies in the SSRS2 have morphological classifications we can also analyze the clustering dependence on morphology. With this aim, we have calculated \\xis and $\\xi(r)$ for the SSRS2 for different morphological types, dividing galaxies into broad morphological  bins - early types comprising  E, S0 and S0-a, and late types containing Sa galaxies and later. In contrast to the results of the Stromlo-APM, the luminosity function parameters used in the selection function for both samples are quite similar to those measured for the SSRS2 as a whole (Marzke et al. 1997). Furthermore, since sample-to-sample variations are within our estimated errors both for the magnitude and volume-limited samples, as shown in Section 4, in the analysis below we only consider the combined sample to improve the statistics.  The resulting correlation functions for early and late type galaxies are presented in Figure 10 panel (a) in redshift space and (b) in  real space, while the fit parameters are presented in Table 5. For the late type galaxies we find that the  correlation function is adequately described by ($s_0$ = 5.4\\mm0.2\\h1 Mpc; $\\gamma_s$=1.48\\mm0.09), while for early types we find  $s_0$ = 6.5\\mm0.2 \\h1 Mpc; $\\gamma_s$=1.86\\mm0.11. Our values for early type galaxies are close to those of Santiago \\& da Costa for the diameter-limited SSRS ($s_0$ = 6.0\\mm1.5 \\h1 Mpc, $\\gamma_s$=1.69)  and Hermit \\etal (1996) for the ORS, who measure $s_0$=6.7 and $\\gamma_s$=1.52. Our value for late types is somewhat larger than that measured by Santiago \\& da Costa (1990), while a proper comparison  with Hermit \\etal (1996) cannot be made, because we have not subdivided spirals into earlier (Sa/Sb) and later (Sc/Sd) types as they did. A comparison between the fit parameters obtained from available redshift space correlation functions, is shown in Fig. 11, where open symbols represent fits for late type galaxies and solid symbols represent early types.  Although all works agree that early types are more clustered than late types, as indicated by the larger correlation length, the scatter is large with the Stromlo-APM results yielding very extreme results. This, in turn, implies large uncertainties in the measurement of the relative bias between the two populations.  Based on the redshift space information, we estimate the relative bias between morphological types as 1.25.  However, for a proper estimate of the dependence of the correlation properties on morphology it is important to take into account the fact that redshift distortions may affect early and late type galaxies in different ways. Therefore a more meaningful comparison must be carried out in real space. The values we measure for the correlation length in real space for early types ($r_0$=6.0\\mm0.4; $\\gamma_r$=1.91\\mm0.18) show a very good  agreement with those of Loveday \\etal (1995), ($r_0$=5.9\\mm0.7; $\\gamma_r$=1.85\\mm0.13). For late types we find ($r_0$=5.3\\mm0.3; $\\gamma_r$=1.89\\mm0.15), which is somewhat larger than those measured by the same authors ($r_0$=4.4\\mm0.1; $\\gamma_r$=1.64\\mm0.05).  Our value of the correlation length is significantly smaller than that measured by Guzzo et al. (1997) ($r_0$=8.4\\mm0.8; $\\gamma_r$=2.05\\mm0.09) for early types. We should note that their sample is volume-limited at $M < -19.5$, whereas we consider galaxies down to $M = -13$.  In order to compare with these authors we consider a volume-limited sub-sample of SSRS2 galaxies with $M \\leq$ -19.5, which corresponds to maximum distance of 100 \\h1 Mpc. Using this sample, for early types we measure ($r_0$=5.7\\mm0.8; $\\gamma_r$=2.09\\mm0.49) while for late types we find ($r_0$=5.0\\mm0.5; $\\gamma_r$=2.01\\mm0.28).  For both early as well as late types, there are still discrepancies relative to the results of Guzzo et al. (1997), which could reflect the paucity of rich clusters in our sample.  By using the variance, we estimate that the relative bias between the different morphologies is $b_{E+S0}/b_S$ = 1.18$\\pm$0.15 in a sample where clusters are not important.  This value is smaller than the determination derived from the real-space correlations of Loveday et al. (1995)  $b_{E+S0}/b_S$ = 1.33 and Guzzo et al. (1997) $b_{E+S0}/b_S$ = 1.68. From these results we may conclude that the relative bias between the two populations range from roughly 1.2 to 1.7, depending on the cluster abundance in the sample, with the former value representing a lower limit.  We have also calculated the correlation function for galaxies discriminated by morphological types for volume-limited samples using the same absolute magnitude limits as in Section 5.1. This calculation was carried out both in redshift, as well as real space, and the results are presented in Tables 6 and 7 respectively.  In redshift space there is a trend of the correlation function amplitude  increasing with luminosity for both morphological classes.  The magnitude of this variation is larger for early types than for late types, although the errors are large. The same trend may be inferred from the analysis in real space, as shown in Figure 12(a), where we compare the $\\sigma_8$ values obtained for the different sub-samples. Here it may be clearly seen that there is a trend of $\\sigma_8$ increasing with luminosity, suggesting that the morphological and luminosity segregations are two separate effects.  Using equation (15) we can also examine how the relative bias varies as a function of scale. This is shown in Figure 12(b), using the real space correlation functions. In contrast to the luminosity bias we find that the morphological bias presents a small decrease from $\\sim$ 1.4 on small scales to $\\sim$ 1.0 on larger scales ($\\sim$ 8 \\h1 Mpc). Although the latter value is slightly smaller than that estimated through the $\\sigma_8$ values ($b_{E+S0}/b_S$ = 1.18 $\\pm$ 0.15), it is still within the estimated error. A similar behavior of the morphological bias changing with scale, was found by Hermit et al. (1996) but using the redshift space correlation function of the ORS, which may not be as meaningful, because of possible biases introduced by virial motions.  Taken together, the above results are consistent with the interpretation that luminosity segregation could be a primordial effect, while the morphological segregation could be enhanced by environmental effects (e.g. Loveday \\etal 1995).  \\subsection{Colors}  Another internal characteristic available in the present catalog is color. Although morphology and colors are correlated the scatter is large, and galaxies of a given type exhibit a broad range of colors, indicating different star-formation histories. On the other hand, colors are easily measured and are an objective criterion, in particular for samples of distant galaxies, whereas the morphological classification is somewhat subjective and becomes increasingly difficult to carry out as the apparent sizes of galaxies get smaller. A further evidence that morphology and colors have somewhat different distributions comes from the calculation of the luminosity function, which presents significantly different shapes for blue and red galaxies (Marzke \\& da Costa 1997), while the luminosity function calculated by separating galaxies between early and late types presents similar Schechter parameters (Marzke et al. 1997).  The few works calculating the correlation properties of galaxies divided by colors present rather conflicting results for the deep samples. Works by Infante \\& Pritchet (1993) and Landy, Szalay \\& Koo (1996) using the angular correlation function show that the correlation of redder galaxies is significantly stronger than for bluer galaxies, except for the very bluest ones (Landy et al. 1996). Carlberg et al. (1996) analyzing a redshift survey of K-band selected galaxies, find that for $0.3 \\leq z \\leq 0.9$ red galaxies are more correlated than blue galaxies by a factor of five. These results differ from those of Le F\\`evre et al. (1996) who find that at $z \\geq$ 0.5 blue and red galaxies have the same correlation properties, while for 0.2 $\\leq z \\leq$ 0.5 blue galaxies are less correlated than red ones. For nearby galaxies, Tucker et al. (1996) have calculated the correlation function and showed that at small scales ($s \\leq 10$ \\h1 Mpc) red galaxies ($[b_J - R]_0 > 1.25$) cluster more strongly than blue ($[b_J - R]_0 < 1.05$) ones, while for larger scales no evidence of color segregation is seen.  In order to make an independent estimation of the dependence of \\xis on colors, we use the the $m_B$ = 14.5 sample described in Section 2, which contains galaxies in both galactic hemispheres. As mentioned in Section 2, this bright limit was used because of incompleteness in colors, as we are restricted to galaxies with measurements in the Lauberts \\& Valentijn (1989) catalog.  In this work we adopted the restframe color cutoff as $(B_T-R_T)_0$ = 1.3  which is roughly the color of an Sbc galaxy, and was the criterion adopted by Marzke \\& da Costa (1997) in the determination of the luminosity function by colors. This value is close to the median value of $B_T-R_T$ in our sample which is $B_T-R_T$=1.2. The conversion of observed into restframe colors used the no-evolution models calculated by Bruzual \\& Charlot (1993), where we assume that the B and R measures in the Lauberts \\& Valentijn (1989) catalog are on the same system of $b_J$ and $r_F$ used by Bruzual \\& Charlot (1993). To calculate \\xis we used the following Schechter function parameters; for blue galaxies ($B_T-R_T \\leq 1.3$), $M^*$ = -19.43, $\\alpha$ = -1.46; for red galaxies ($B_T-R_T > 1.3$), $M^*$ = -19.25, $\\alpha$ = -0.73, which were obtained by Marzke \\& da Costa (1997). The sample, which only considers galaxies out to a maximum distance of 8000 \\kms, contains 387 blue and 219 red galaxies.  The results of the two-point correlation function are shown in Figure 13 (a) for redshift space while the fit  parameters may be found in Table 8.  Because of the small number of objects, the correlation function is very noisy, yet it is unquestionable that the red galaxies present a systematically higher amplitude at all separations compared to blue galaxies. In order to verify how sensitive the results may be to incompleteness, we re-calculated \\xis for the $m_B$=14.2 sample which is 92 \\% complete in colors. The fit parameters present a similar behavior, although the values differ from those measured for the 14.5 sample. The results we obtain for the samples discriminated in colors present a qualitative agreement with those of Tucker \\etal  (1996), in the sense that red galaxies are more strongly correlated than blue galaxies.  We have also calculated the real-space correlation function for the 14.5 sample and the power-law fit is presented in Fig. 13 (b), together with the redshift space correlation. The figure shows that the slopes of both power law fits are fairly  similar ($\\gamma_r$=1.99 for blue, $\\gamma_r$=2.18 for red  galaxies), though the uncertainties are rather large, in particular for the red galaxies. The observed \\xir suggests that red galaxies are probably more affected by peculiar motions than blue galaxies. Because of the relatively small size of the sample with colors, we have not been able to  investigate the dependence on luminosity, which would be dominated by errors because of the small number of objects assigned to each luminosity bin.   The relative bias estimated from $\\sigma_8$  in real space is $b_R$/$b_B$ = 1.40\\mm 0.33, and a similar result is obtained if the redshift space results are considered.  As in the case of luminosity and morphology, one may calculate the relative bias between galaxies of different colors as a function of scale, which is presented in Fig. 14. Because the observed correlation function is rather noisy, for this plot we used the fits to \\xir . Taking the results at face value they would suggest that the relative bias between red and blue galaxies on small scales is comparable to that seen for early and late type galaxies. However, it levels off more rapidly ($\\sim $ 4 \\h1 Mpc), remaining constant at $b_R/b_B \\sim 1.2$ thereafter. This behavior could be the result of evolution due to environmental effects, where early type galaxies in higher density regions lost their gas more rapidly than bluer galaxies, and thus present a much lower star formation rate. However,  because the errors are large, these results should only be considered as tentative. ",
        "conclusions": ""
    },
    "9803/astro-ph9803083_arXiv.txt": {
        "abstract": "We review our present knowledge of high-redshift galaxies, emphasizing particularly their physical properties and the ways in which they relate to present-day galaxies. We also present a catalogue of photometric redshifts of galaxies in the Hubble Deep Field and discuss the possibilities that this kind of study offers to complete the standard spectroscopically based surveys. ",
        "introduction": "For a long time models for galaxy formation and evolution advanced unhampered by observations. Nowadays, however, the rapid increase in both observational capabilities and efficiency of the selection methods (see Steidel {\\em et al.} 1995 [S95]) has converted the task of looking for distant galaxies from one of the most difficult challenges to an almost routine job, and large databases of high-$z$ galaxies are already being compiled (Dickinson 1998, this Volume). Observations can now constrain the models, and this obliges us to understand the properties of these objects in order to get a complete image of the processes involved in the formation and evolution of galaxies.  This study of the properties of high-$z$ galaxies is twofold.  We need to understand the information provided by the confirmed high-$z$ galaxies. In this way we will learn about the spectral and morphological properties of the bright end of the galaxy population, i.e., the putative progenitors of present-day large ($L>L_*$) galaxies. Second, the use of photometric redshift techniques applied to deep multi-colour images (like the HDF, Williams {\\em et al.} 1996) opens a wealth of statistical methods to study those faint objects for which we cannot obtain spectroscopic information in the near future. These studies will yield further results on the general distribution and evolution of galaxies. The main problem for both methods resides in the $z \\approx 1-2$ range, where spectroscopic identification of galaxies at optical wavelengths is made difficult by the lack of spectral features. ",
        "conclusions": "The available data allow for different interpretations. While S95, S96 and G96 support the hypothesis that the observed high-$z$ galaxies are the progenitors of present-day luminous galaxies at the epoch of formation of the first stars in their spheroidal components, T97 suggests that these objects will evolve to form the Population II components of early-type spirals.  Another interpretation (L97) maintains that these objects represent a range of physical processes and stages of galaxy formation and evolution rather than any particular class of object.  While this third interpretation might be closer to reality, we are still missing an important piece of the puzzle. Detailed IR imaging and spectroscopy is needed in order to: a) shed light on the $z=1-2$ galaxies allowing us to constrain evolutionary models; b) obtain images of the $z>2$ galaxies at optical rest-frame wavelengths to be compared with their low-$z$ counterparts and; c) perform moderate resolution spectroscopy of the $z>2$ galaxies to accurately measure their metallicities and the importance of dust corrections.  We expect that these observations, with the support of techniques like cosmological simulations and stellar population evolutionary models, will lead us closer to the long-searched-for understanding of the process by which the Universe came to be as we see it. Perhaps it is not the moment for us to ``look deeper in the Southern Sky'', but to look at it with different eyes."
    },
    "9803/astro-ph9803340_arXiv.txt": {
        "abstract": "We examine the non-linear stability of the Wisdom-Holman (WH) symplectic mapping applied to the integration of perturbed, highly eccentric ($e\\gtrsim 0.9$) two-body orbits.  We find that the method is unstable and introduces artificial chaos into the computed trajectories for this class of problems, {\\it unless} the step size is chosen small enough to always resolve periapse, in which case the method is generically stable. This `radial orbit instability' persists even for weakly perturbed systems. Using the Stark problem as a fiducial test case, we investigate the dynamical origin of this instability and show that the numerical chaos results from the overlap of step size resonances (cf. Wisdom \\& Holman 1992); interestingly, for the Stark problem many of these resonances appear to be absolutely stable.  We similarly examine the robustness of several alternative integration methods: a regularized version of the WH mapping suggested by Mikkola (1997); the potential-splitting (PS) method of Lee et al. (1997); and two methods incorporating approximations based on Stark motion instead of Kepler motion (cf. \\cite{newet97}). The two fixed point problem and a related, more general problem are used to comparatively test the various methods for several types of motion.  Among the tested algorithms, the regularized WH mapping is clearly the most efficient and stable method of integrating eccentric, nearly-Keplerian orbits in the absence of close encounters. For test particles subject to both high eccentricities and very close encounters, we find an enhanced version of the PS method---incorporating time regularization, force-center switching, and an improved kernel function---to be both economical and highly versatile. We conclude that Stark-based methods are of marginal utility in $N$-body type integrations. Additional implications for the symplectic integration of $N$-body systems are discussed. ",
        "introduction": "\\label{sec_intro}  Symplectic integration schemes have become increasingly popular tools for the numerical study of dynamical systems, a result of their often high efficiency as well as their typical long-term stability (see, e.g., \\cite{marps96} and the many references within). The Wisdom-Holman (WH) symplectic mapping in particular (\\cite{wish91}; cf. \\cite{kinyn91}) has been widely used in the context of Solar System dynamics. However, the fact that this and other symplectic methods are, by construction, ``finely tuned'' can make them susceptible to performance-degrading ailments (much as high-order methods offer little benefit if the motion is not sufficiently smooth), and the stability of these methods for arbitrary systems and initial conditions is not completely understood.  It would be prudent, therefore, to exercise caution when applying such schemes to systems entering previously unexplored dynamical states, and to ensure that adequate preliminary testing is undertaken regardless of the method's stability in previously considered problems. Recent galactic dynamics simulations by Rauch \\& Ingalls (1997), for example---which used the WH mapping---uncovered evidence of an instability in the method when applied to a particular class of problems: the integration of highly elliptical, nearly-Keplerian orbits in which the timestep is taken small enough to smoothly resolve the perturbation forces, but not so tiny as to explicitly resolve pericenter. Since particle motion in these simulations was extremely close to Keplerian near pericenter, and since the mapping itself is exact for Keplerian motion, there is no {\\it a priori} reason why the method should have performed as poorly and unstably as was found.  Recently several variations of the WH mapping have been proposed which aim to extend the range of applicability of the original method. The regularized WH mappings investigated by Mikkola (1997), for instance, appear promising in the context of elliptical motion. The potential-splitting (PS) method of Lee, Duncan, \\& Levison (1997) allows symplectic integration of close encounters between massive bodies by adding a multiple-timestep algorithm similar to that of Skeel \\& Biesiadecki (1994). Unfortunately both of these schemes have limitations of their own; the former is unable to resolve close encounters, while the latter approach (like the original mapping) appears to be unstable when orbits are eccentric (cf. \\cite{dunll97}).  In this paper, we use a series of test problems based on perturbed two-body motion to analyze the stability of the WH mapping and several of its variants. In particular, we examine the reliability of the methods for test particles whose motion is either highly eccentric or subject to close encounters with the perturbers (or both). The plan of the paper is as follows. In the following section, the performance of the WH mapping at high eccentricities is investigated using the Stark problem (e.g., \\cite{dan94}; \\cite{kir71})---for which the range of orbital eccentricities is easily controlled, and no close encounters occur---as the fiducial test case. The instability found in the integrated motion is then explained using complementary analytic and geometric arguments. In \\S~\\ref{sec_modwh}, modified forms of the original mapping are described; similarly, in \\S~\\ref{sec_sint} integrators based on Stark motion instead of Kepler motion are considered. Section~\\ref{sec_compsim} uses the two fixed point problem (e.g., \\cite{par65}) as well as a more general test problem (drawn from the area of galactic dynamics) to conduct a comparative performance analysis of the various algorithms.  Both the Stark and two fixed point problems are fully integrable and analytically soluble in terms of elliptic functions and integrals, allowing a detailed assessment of the accuracy of the numerical results to be made. Concluding discussion is given in \\S~\\ref{sec_discuss}. ",
        "conclusions": "\\label{sec_discuss}  We have shown that the WH mapping is generically unstable when applied to eccentric, nearly-Keplerian orbits whenever the step size is not small enough to resolve periapse. This `radial orbit instability' is fully explainable in terms of the overlap of step size resonances and has a simple geometric manifestation in the case of the Stark problem. Our investigation indicates that the islands of stability found in the latter problem do not exist in the more general cases we have examined; the instability therefore appears to be unavoidable in typical situations, unless one employs the brute-force approach of decreasing the timestep by the requisite amount. However, besides the fact that this is an extremely inefficient solution---it reduces the mapping to a very costly direct integration scheme---we have shown that an elegant solution to the problem is already available: the regularization approach of Mikkola (1997). In every case examined, not only was the regularized WH mapping immune to the radial orbit instability, in many cases it was also more efficient. We enthusiastically recommend its use whenever close encounters with perturbers are not of concern.  We remind the reader that our investigation has not cast into doubt all previous studies that have used the WH integrator and its variants.  In nearly every case care has been taken to use a small enough step size such that perihelion passage would be adequately sampled.  We note only one area where particular caution should be exercised.  One of the features of the long-term dynamics in mean motion resonances and secular resonance is that very high eccentricities can be developed.  These eccentricities are often large enough that physical collision with the sun is a common outcome in studies of meteorite delivery from the main asteroid belt and the long-term dynamics of ecliptic comets (\\cite{glaet97}; \\cite{morm95}; \\cite{levd97}). In those cases it is unlikely that the step size used was small enough to resolve the perihelion passage.    Although these researchers checked their results for step size dependence and reported no numerical artifacts, we suggest that further examination of those cases would be prudent.  We have demonstrated that the potential-splitting method of Lee et al. (1997) can be regularized to produce an algorithm that is robust in the face of both close encounters and highly eccentric orbits. We have also shown that force-center switching during exceptionally close encounters can be cleanly incorporated into the method and can substantially enhance the stability of the algorithm without noticeably affecting its desirable symplectic qualities. We have not, however, found a practical way to regularize around the perturber while the switch is in effect; the stability of this approach during highly eccentric {\\it encounters} is correspondingly questionable.  Our examination of Stark-based integrators indicates that they, too, are subject to the radial orbit instability unless regularized, although it tends to be less severe since the Stark approximation becomes systematically better near the origin. Unless the perturbing potential is {\\it very} well represented by a Stark potential, they also appear uncompetitive in terms of efficiency---by over an order of magnitude---due to the cost of Stark steps relative to that of Kepler steps. Among the Stark-based methods, the regularized, symplectic method (\\S~\\ref{sec_rss}) consistently outperformed the time-reversible method (\\S~\\ref{sec_trs}), in part because of the linearly growing energy error exhibited by the latter. Our conclusion is that Stark-based schemes are of marginal utility in the integration of N-body systems.  It is clear that integrators based on a two fixed point (TFP) splitting instead of a Stark or Kepler splitting are also possible; they can be constructed in the same manner as the Stark-based methods were. Such methods could be useful whenever two bodies strongly dominate the mass in the system (e.g., asteroid motion in the Sun-Jupiter system). As for Stark motion, however, the relative expense of advancing the TFP Hamiltonian is a significant handicap, and the circumstances in which its use is justified remain unclear. On the other hand, since Stark motion is a subset of TFP motion it is not unlikely that methods based on the latter splitting will generically outperform those of the former type, since their analytic solutions are of similar complexity. It would be interesting to investigate this possibility in greater detail.  Although we have confined attention to the perturbed two-body problem, the techniques employed in this paper are also applicable to general hierarchical $N$-body systems. In particular, we believe that regularization of the $N$-body version of the potential-splitting method (\\cite{dunll97}) is likely to cure the instability at high eccentricities noted by the authors. In principle force-center switching of the kind described in \\S~\\ref{sec_ps} can also be done, but we have not studied this possibility in detail.  In any event, we have found the combination of regularization and potential-splitting to be a powerful one, and to produce a remarkably versatile symplectic method for the integration of nearly-Keplerian systems."
    },
    "9803/astro-ph9803263_arXiv.txt": {
        "abstract": " ",
        "introduction": "The discovery of strong and remarkably coherent high-frequency X-ray brightness oscillations in at least sixteen neutron stars in low-mass binary systems has provided valuable new information about these stars, some of which are likely to become millisecond pulsars. Oscillations are observed both in the persistent X-ray emission and during thermonuclear X-ray bursts (see van der Klis 1997).  The kilohertz quasi-periodic oscillations (QPOs) observed in the persistent emission have frequencies in the range 325--1200~Hz, amplitudes as high as $\\sim15$\\%, and quality factors $\\nu/\\delta\\nu$ as high as $\\sim200$. Two kilohertz QPOs are commonly observed simultaneously in a given source (see Fig.~1). Although the frequencies of the two QPOs vary by  hundreds of Hertz, the frequency separation $\\Delta\\nu$ between them appears to be nearly constant in almost all cases (see van der Klis et al.\\ 1997 and M\\'endez et al.\\ 1997).   \\begin{figure}[t] % \\begin{minipage}[b]{3.3in} \\centerline{ \\psfig{file=tokyo.fig1.eps,height=7.5truecm}} \\end{minipage} \\begin{minipage}[b]{2.3in} \\caption{Power density spectrum of Sco~X-1 brightness variations, showing the two simultaneous kilohertz QPOs that are characteristic. These are two of the weakest kilohertz QPOs observed, with rms amplitudes $\\sim1$\\%. The continuum power density is consistent with that expected from photon counting noise. From van der Klis et al.\\ (1997).} \\end{minipage} \\end{figure}   The $\\sim$250--600~Hz brightness oscillations observed during type~I X-ray bursts are different in character from the QPOs observed in the persistent emission (see Strohmayer, Zhang, \\& Swank 1997). Only a single oscillation has been observed during X-ray bursts, and the oscillations in the tails of bursts appear to be highly coherent (see, e.g., Smith, Morgan, \\& Bradt 1997), with frequencies that are always the same for a given source (comparison of burst oscillations from \\fu{1728$-$34} over about a year shows that the timescale for any variation in the oscillation frequency is $\\gta 3000$~yr; Strohmayer 1997). The burst oscillations in \\fu{1728$-$34} and \\fu{1702$-$42} (see Strohmayer, Swank, \\& Zhang 1998) have frequencies that are consistent with the separation frequencies of their kilohertz QPO pairs. The burst oscillations in \\fu{1636$-$536} (Zhang et al.\\ 1997) and \\ks{1731$-$260} (Smith et al.\\ 1997) have frequencies that are consistent with twice the separation frequencies of their kilohertz QPO pairs (Zhang et al.\\ 1997; Wijnands \\& van der Klis 1997). The evidence is compelling that the burst oscillations are produced by rotation with the star of one or two nearly identical emitting spots on the surface (see Strohmayer et al.\\ 1997). The frequencies of the burst oscillations are therefore the stellar spin frequency or its first overtone.  The frequency separation $\\Delta\\nu$ between the two kilohertz QPOs observed in the persistent emission of a given star is closely equal to the spin frequency of the star inferred from its burst oscillations (see Miller, Lamb, \\& Psaltis 1998, hereafter MLP). ",
        "conclusions": ""
    },
    "9803/astro-ph9803055_arXiv.txt": {
        "abstract": " ",
        "introduction": "During the first thousand seconds in the evolution of the Universe, as it expanded and cooled from very high densities and temperatures, nuclear reactions transformed neutrons and protons into astrophysically interesting abundances of the light nuclides deuterium, helium-3, helium-4 and lithium-7. In the context of Standard, Big Bang Nucleosynthesis (SBBN; homogeneous, isotropic expansion, three flavors of non-degenerate neutrinos) these abundances depend on only one adjustable parameter, the nucleon density. Since as the Universe expands all densities decrease, it is useful to express the nucleon density in terms of a nearly constant parameter, the ratio of nucleons to photons, which has barely changed at all since the annihilation of electron-positron pairs in the early Universe.  \\begin{equation} \\eta \\equiv n_{\\rm N}/n_{\\gamma} \\ \\ ; \\ \\ \\eta_{10} \\equiv 10^{10}\\eta \\end{equation} The contribution of nucleons (baryons) to the universal mass-density may be written as the dimensionless ratio of the baryon density to the critical density (which depends on the present value of the Hubble parameter: H$_{0} = 100\\,h\\,$kms$^{-1}$Mpc$^{-1}$; $\\Omega_{\\rm B} \\equiv \\rho_{\\rm B}/\\rho_{crit}$).  \\begin{equation} \\Omega_{\\rm B}\\,h^{2} = \\eta_{10}/273 \\end{equation} SBBN is an overdetermined theory in that the observable abundances of four nuclides are predicted on the basis of one free parameter. In Figure 1 the predictions of the primordial abundances are shown for a wide range of $\\eta$.  SBBN is falsifiable in that it is possible that {\\bf no} value of $\\eta$ will be consistent with the primordial abundances inferred from the observational data.  Furthermore, consistency requires that {\\bf if} an acceptable value of $\\eta$ is found, the corresponding nucleon density at present, $\\Omega_{\\rm B}$, is in agreement with other astronomical observations.  Indeed, since there must be enough baryons to account for the visible matter in the Universe, but not too many to violate constraints on the total mass density, the {\\it interesting} range of $\\eta$ in Figure 1 is restricted to $3\\times 10^{-11} - 1\\times 10^{-8}$.  Even so, note the enormous range in the predicted abundances of deuterium and lithium.  Over this same range in $\\eta$ the predicted primordial mass fraction of $^4$He, Y$_{\\rm P}$, hardly changes at all. As we shall soon see, consistency between D and $^4$He provides a key test of SBBN.  \\begin{figure} \\centerline{\\psfig{file=fig1.ps,width=0.9\\textwidth}} \\vspace{-24pt} \\caption{SBBN-predicted abundances of the light nuclides versus $\\eta$. The $^4$He mass fraction (Y$_{\\rm P}$) is shown along with the ratio by number to hydrogen of D ($^2$H, $y_2$), $^3$He ($y_3$), and $^7$Li ($y_7$). This figure is from D. Thomas.} \\end{figure}  \\subsection{Status Quo Ante}  SBBN has provided one of the most spectacular confirmations of the standard, hot Big Bang model of cosmology.  Along with the Hubble expansion and the cosmic background radiation, SBBN is one of the pillars of the standard model.  It is the only one offerring a connection between particle physics and cosmology.  For example, Walker \\etal (1991) reanalyzed the relevant observational data to make a critical confrontation between predictions and observations.  Walker \\etal (1991) concluded that SBBN was consistent with the observational data for $\\eta_{10} = 3.4\\pm0.3$ ($\\Omega_{\\rm B}h^{2} \\approx 0.01$), making the nucleon density one of the very best determined of all cosmological parameters.  Furthermore, they noted that to preserve this consistency required that the total number of ``equivalent\", light neutrinos (particles which were relativistic at BBN), N$_{\\nu}$, should not exceed 3.4.  With the three known flavors of neutrinos (provided none has a mass comparable to MeV energies), this leaves very little room for any new (light) particles ``beyond the standard model\".  At this point it may have been tempting to declare victory for SBBN and to move on to other problems in cosmology.  However, it was still important to subject the standard model to ever more precise observational tests in order to reaffirm its consistency and to narrow even further the bounds on the nucleon density and on particle physics beyond the standard model.  To our surprise, my colleagues and I found a dark cloud looming on the horizon of the standard model (Hata \\etal 1995).  \\subsection{A Crisis For SBBN?} There had, in fact, always been a ``tension\" between the predictions of SBBN and the inferred primordial abundances of D and $^4$He (Kernan \\& Krauss 1994, Olive \\& Steigman 1995) in the sense that while deuterium favored ``high\" values of $\\eta$ (Steigman \\& Tosi 1992, 1995), helium-4 pointed towards lower values (Olive \\& Steigman 1995).  Indeed, in a reanalysis focusing on the $^4$He abundance, Olive \\& Steigman (1995) found for the best estimate of the number of equivalent light neutrinos, N$_{\\nu} = 2.2$.  Only a generous error estimate permitted consistency with SBBN.  It was, therefore, not entirely unexpected when Hata \\etal (1995) identified a ``crisis\" for SBBN in their comparison of the best estimates of the primordial abundances derived from the observational data with those predicted by SBBN.  The problem is illustrated in Figure 2 which concentrates on the key nuclides, D, $^4$He and $^7$Li.  While the $^4$He abundance is just barely consistent with the low end of the $\\eta$ range identified by Walker \\etal (1991), the deuterium abundance is only consistent with the upper end of that range.  Note that due to its ``valley\" shape and to the relatively larger uncertainties in its predicted and inferred abundances, lithium is consistent with either deuterium or helium.  Since it thus fails to discriminate between the low $\\eta$ favored by helium and the higher $\\eta$ preferred by deuterium, lithium is ignored in the following discussion. \\begin{figure} \\centerline{\\psfig{file=fig2.ps,width=0.9\\textwidth}} \\caption{SBBN predictions (solid lines) for $^4$He (Y), D ($y_2$), and $^7$Li ($y_7$) with the theoretical uncertainties (1$\\sigma$) estimated by the Monte Carlo method (dashed lines).  Also shown are the regions constrained by the observations at 68\\% and 95\\% C.L. (shaded regions and dotted lines, respectively).  This figure is from Hata \\etal (1995).} \\end{figure}  Three possible resolutions of the challenge to SBBN posed by the D -- $^4$He conflict suggest themselves.  Perhaps the primordial abundance of helium inferred from observations of extragalactic \\hii regions (see, \\eg, Olive \\& Steigman 1995 and Olive, Skillman, \\& Steigman 1997) is too small (see, \\eg, Izotov, Thuan, \\& Lipovetsky 1994 and Izotov \\& Thuan 1997).  If the primordial helium mass fraction were closer to 0.25 than to 0.23, the challenge to SBBN evaporates.  Since several dozen \\hii regions are observed, the statistical uncertainty in Y is small, typically $\\pm 0.003$ or smaller (Olive, Skillman, \\& Steigman 1997, Izotov \\& Thuan 1997). But systematic errors, such as those due to uncertainties in the corrections for unseen neutral helium, for collisional ionization, for temperature fluctuations and, especially, for underlying stellar absorption, may well be much larger.  Alternatively, it could be that our adopted primordial deuterium abundance is too small.  If the true primordial ratio (by number) of deuterium to hydrogen were a few parts in $10^4$ rather than the few parts in $10^5$ inferred from observations in the solar system and the local interstellar medium (ISM), lower $\\eta$, consistent with Y$_{\\rm P}$, is allowed (see Fig. 2).  This local estimate of the deuterium abundance requires an extrapolation from ``here and now\" (solar system, ISM) to ``there and then\" (primordial).  Any errors in this extrapolation open the door to systematic errors.  Finally, the possibility remains that our estimates of the primordial abundances are correct and the D -- $^4$He tension is a hint of ``new physics\".  For example, if the tau neutrino were massive ($\\sim 5 - 20$ MeV) and unstable (lifetime $\\sim 0.1 - 10$ sec.), the ``effective\" number of equivalent light neutrinos would be less than the standard model case of N$_{\\nu}$ = 3 (Kawasaki \\etal 1994). For N$_{\\nu} = 2.1 \\pm 0.3$, consistency among the primordial abundances may be reestablished (Hata \\etal 1995, Kawasaki, Kohri, \\& Sato 1997). Other, non-standard, particle physics solutions are conceivable; degenerate neutrinos offer one such option (Kohri, Kawasaki, \\& Sato 1997). ",
        "conclusions": "The predictions of SBBN are observationally challenged.  The primordial abundances of D and $^4$He inferred from observational data appear to be inconsistent with the predictions of SBBN.  Several options present themselves with the potential to resolve this crisis.  Perhaps the data are at fault.  The conflicting deuterium abundances derived from observations of high-redshift, low-metallicity QSO absorbers point an incriminating finger.  If these data are supplemented with solar system and ISM deuterium abundances, the lower D/H ratios are preferred.  But, is the extrapolation from here and now (solar system, ISM) to there and then (primordial) under control, or might there be unidentified systematic errors lurking?  The two sets of apparently inconsistent helium abundances suggest systematic errors at play in the extragalactic \\hii region abundance determinations. Although new data is always welcome, it is clear that a better understanding of existing data may prove even more important.  In the absence of new data and/or a better understanding of the extant data it may be worthwhile to look elsewhere for clues.  My colleagues and I (Steigman, Hata, \\& Felten 1997; SHF) have discarded the constraint on $\\eta$ from SBBN and have utilized four other observational constraints (Hubble parameter, age of the Universe, cluster gas (baryon) fraction, and effective ``shape\" parameter $\\Gamma$) to predict the three key cosmological parameters (Hubble parameter, total matter density, and the baryon density or $\\eta$).  Considering both open and flat CDM models and flat $\\Lambda$CDM models, SHF tested goodness of fit and drew confidence regions by the $\\Delta\\chi^2$ method.  In all of these models SHF find that large $\\eta_{10}$ ($\\gsim~6$) is favored strongly over small $\\eta_{10}$ ($\\lsim~2$), supporting reports of low deuterium abundances on some QSO lines of sight, and suggesting that observational determinations of primordial $^4$He may be contaminated by systematic errors."
    },
    "9803/astro-ph9803325_arXiv.txt": {
        "abstract": "We present the results of a combined study of ASCA and ROSAT observations of the distant cluster Abell 2390. For this cluster a gravitational arc as well as weak lensing shear have been previously discovered.  We determine the surface brightness profile and the gas density distribution of the cluster from the ROSAT PSPC and HRI data. A combined spatially resolved spectral analysis of the ASCA and ROSAT data show that the temperature distribution of the intracluster medium of A2390 is consistent with an isothermal temperature distribution in the range 9 to 12 keV except for the central region. Within a radius of $160 h_{50}^{-1}$ kpc the cooling time is found to be shorter than the Hubble time, implying the presence of a cooling flow. In this central region we find strong evidence for a multi-temperature structure. Detailed analysis of the combined ASCA and ROSAT data yields a self-consistent result for the spectral structure and the surface brightness profile of the cluster with a cooling flow of about $500 - 700$ M$_{\\odot}$ y$^{-1}$ and an age of about $10^{10}$ y.  From the constraints on the temperature and density profile of the intracluster gas we determine the gravitational mass profile of the cluster and find a mass of about $2\\cdot 10^{15}$ M$_{\\odot}$ within a radius of $3 h_{50}^{-1}$ Mpc. A comparison of the projected mass profiles of the cluster shows an excellent agreement between the mass determined from X-ray data and the mass determined from the models for the gravitational arc and the weak lensing results. This agreement in this object, as compared to other cases where a larger lensing mass was implied, may probably be due to the fact that A2390 is more relaxed than most other cases for which gravitational lensing mass and X-ray mass have been compared so far. ",
        "introduction": "A 2390 is one of the most prominent clusters in the redshift range around z = 0.2. It was classified by Abell (1958) and Abell, Corwin \\& Olowin (1989) only as a richness class 1 cluster. As a target of the CNOC survey (Yee \\et\\ 1996a) deep photometric and spectroscopic data were obtained for this cluster, and these authors conclude that it should more likely be classified as a richness class 3 cluster (Yee \\et\\ 1996b). This new CNOC data also provide a mean redshift for this cluster of z=0.228 (compared to the previous literature value of z=0.232 by Le Borgne \\et\\ 1991).  In X-rays it is among the ten brightest galaxy clusters known at a redshift larger than 0.18 (e.g. Ebeling \\et\\ 1996). It has been observed with the EINSTEIN observatory and showed a luminosity of $L_x \\sim 1.6 \\cdot 10^{45}$ \\egs (in the 0.7 to 3.5 keV energy band) (Ulmer \\et\\ 1986; if their result is converted to $H_0 = 50$ km s$^{-1}$ Mpc $^{-1}$ as used in this paper) and the cluster has a slightly elongated shape (McMillian \\et\\ 1989). A ``straight arc'' and several arclets were discovered in this clusters by Pello \\et\\ (1991). All these observations underline that A2390 is a very rich and massive cluster of galaxies.  Pierre \\et\\ (1996) have analyzed a deep ROSAT HRI observation and found that the X-ray emission from A2390 is very concentrated and highly peaked, indicating a strong cooling flow of about 880 \\msu y$^{-1}$. The cooling flow is centered on the giant elliptical galaxy in the cluster center. This together with the observation of the strong lensing features may indicate that the very peaked central surface brightness is probably the effect of the cooling flow as well as of a steep central gravitational potential in the cluster.  The straight arc has been modeled by Kassiola, Kovner, \\& Blandford (1992) and Narashima \\& Chitre (1993). Pierre \\et\\ (1996) have also modeled the lensing cluster with an elliptical potential model and a second clump in close consistency with the X-ray morphology. They find a projected mass within the arc radius of $M(r \\le 38'') \\sim 0.8 \\cdot 10^{14} h^{-1}$ \\msu. They compared the lensing mass with the X-ray data by taking the mass profile of the lensing model and the gas density profile from the X-ray surface brightness, and calculated the temperature profile needed to satisfy the hydrostatic equation. The bulk temperatures found by this approach are in the range 8 to 10 keV. This high temperature is consistent with the large X-ray luminosity of the cluster given the generally good correlation between X-ray ICM temperature and X-ray luminosity (e.g. Edge \\& Stewart 1991).  A weak lensing shear in A2390 was also observed recently by Squires \\et (1996). They deduced an elliptical mass distribution, elongated in the direction of the straight arc. This is qualitatively consistent with the X-ray surface brightness distribution. As we will show in this paper the mass distribution inferred from the X-ray results are in excellent agreement with both the mass deduced from the weak lensing analysis and that from the strong lensing modeling.  The detection of diffuse intracluster light was reported by V\\'ilchez-G\\'omez \\et\\ (1994). This may also be taken as a sign that the core of the cluster is relaxed and that the debris of tidal stripping of galactic halo material had enough time to settle in the gravitational potential of the cluster.  All these previous studies and the fact that the cluster is very massive and X-ray luminous makes A2390 a perfect target for more detailed X-ray observations, in particular to compare a more precise mass determination from the X-ray data with the optical and lensing results. The indication of the fair agreement of the mass in the different previous studies and the existence of the strong cooling flow suggest that the cluster is essentially relaxed and therefore ideal for the test of the various methods of mass determination.  In this paper we present a combined analysis of deep ASCA and ROSAT PSPC and HRI observations of this cluster. The ROSAT observations are discussed in Section 2 and the ASCA observations in Section 3. A combined analysis of the spectral data of both instruments which provides very interesting evidence for multi-temperature structure in the cooling flow region of A2390 is presented in Section 4. Section 5 contains the results of the cluster mass determination from the X-ray data, and compared to the lensing results in Section 6. The cooling flow structure is discussed in detail in Section 7. Section 8 provides a summary and conclusions. We use a value of $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ for the Hubble constant throughout this paper. Thus 1 arcmin at the distance of A2390 corresponds to a scale of 277 $h_{50}^{-1}$ kpc. ",
        "conclusions": "The present study of A2390 with combined use of the ROSAT PSPC and ASCA GIS data yields important new results. The two most striking results of the current study are (1) good consistency between the cooling flow rates derived independently from the spectral and imaging analyses, and (2) excellent agreement between the total mass values determined from the X-ray data and the gravitational lensing. As outlined in the introduction, the cluster shows all the apparant features of a well relaxed cluster: elliptical symmetry, strong central concentration (and a cD galaxy which may be the result of this), a cooling flow, and a large intracluster light halo around the central galaxy. A long cooling time ($\\sim 10^{10}$ y) implied from the cooling flow rate may be taken as a reinforcement of the picture that the cluster was left quite undisturbed for a long time.  In the morphological analysis of the HRI image of A2390 Pierre et al. (1996) found some indication of substructure in the cluster. They interpreted the substructural feature as a trace of substructure in the cluster potential, and such an excess potential is actually needed in the gravitational lensing model producing the observed gravitational arc. The subclump has an X-ray luminosity of only about 1/60 of that of the whole cluster and therefore its mass is less than about 1/15 of that of the cluster as concluded by Pierre et al. (1996). Such a small infalling mass component will not cause a significant disturbance on the equilibrium configuration of the cluster, and also may not influence the evolution of the cooling flow seriously. Therefore, the presence of this small substructure is not in contradiction to our finding that the cluster is generally well settled.  The large measured iron abundance of $\\sim 0.3$ of the solar value is in line with other observed results for very rich nearby clusters and some distant clusters. We should note, however, that lower abundances have been found in some rich distant clusters like CL0016+16 (Furuzawa et al. 1997) and A851 (Mushotzky \\& Loewenstein, 1997, Schindler et al. 1998).  Earlier studies have pointed out cases of striking differences between the lensing mass and X-ray determined mass (e.g. Miralda-Escud\\'e \\& Babul 1995). The reason for an excellent agreement between the two in A2390 is most probably found in that this cluster is fairly relaxed, as compared to other clusters studied by weak lensing technique and X-ray observations. For example, A2218 and A2163 show signs of recent merging (Squires et al. 1996b, 1997). The detailed study of A2218 shows a tendency that the lensing mass is higher than the X-ray mass, while in A2163 the two mass values are well consistent with each other. PKS0745 which also shows a strong cooling flow and a gravitational arc (Allen et al. 1996) may be in a similar situation to A2390, and consistency between the lensing mass and X-ray mass could be found in this system, too. Allen et al. (1996) have already stressed that the agreement of the mass determination in PKS0745 is most probably the result of the cluster being well relaxed (see also recent work in Allen 1997)."
    },
    "9803/astro-ph9803113_arXiv.txt": {
        "abstract": "Chevalier \\& Ilovaisky (1998) use {\\it Hipparcos} data to show that the X-ray binary systems LSI+61$^\\circ$ 303 and A0535+262 are a factor of ten closer (i.e. $d\\sim$ few hundred pc) than previously thought ($d\\sim2$kpc).  We present high quality CCD spectra of the systems, and conclude that the spectral types, reddening and absolute magnitudes of these objects are strongly inconsistent with the closer distances.  We propose that the {\\it Hipparcos} distances to these two systems are incorrect due to their relatively faint optical magnitudes. ",
        "introduction": "In a recent paper Chevalier \\& Ilovaisky (1998 - CI98) presented {\\it Hipparcos} distances to 17 massive X-ray binary systems.  In particular they presented results that appeared to indicate that two systems (LSI +61$^\\circ$ 303, A0535+262) were up to a factor 10 closer than previously thought.  This has profound implications for any models one constructs for these systems, for instance suggesting they need only contain white dwarfs rather than neutron stars to explain their X-ray luminosity.  In this paper we use CCD spectra to redetermine the spectral type of LSI +61$^\\circ$ 303 and A0535+262. In both cases we show that the derived spectral types and reddenings are strongly consistent with normal Be stars at the distances previously ascribed to the systems, and not with the new, closer distances.  Finally we discuss how the discrepancy between the {\\it Hipparcos} and our distances may be explained in terms of the faint nature of these particular sources. ",
        "conclusions": "We have shown that the {\\it Hipparcos} derived distances ($\\sim$ few hundred pc) to two Be/X-ray binary systems, LSI +61$^\\circ$ 303 and A0535+262 are inconsistent with the spectral types, reddenings and apparent magnitudes of the objects.  There is strong evidence that the `traditional' distances to these objects (each $\\sim 2$kpc) are in fact correct. We note here that these two objects have the worst goodness-of-fit values in the {\\it Hipparcos} catalogue (ESA 1997) of the CI98 sample (although they do lie below the maximum ``acceptable'' value of 3).  In addition they are the faintest in the sample.   This appears to indicate that the application of the simple `goodness-of-fit' criterion that anything less than 3 is a good parallax to faint objects is not reliable, and that the interpretation of {\\it Hipparcos} parallax data should always be carried out with this in mind."
    },
    "9803/astro-ph9803169_arXiv.txt": {
        "abstract": "We report results of $^{12}$CO ($J=1$-0) mapping observations of the Wolf-Rayet starburst galaxy Mrk 1259 which has optical evidence for the superwind seen from a nearly pole-on view. The CO emission is detected in the central 4 kpc region. The nuclear CO spectrum shows a blue-shifted ($\\Delta V \\simeq -27$ km s$^{-1}$) broad (FWHM $\\simeq$ 114 km s$^{-1}$) component as well as the narrow one (FWHM $\\simeq 68$ km s$^{-1}$). The off-nuclear CO spectra also show the single-peaked broad component (FWHM $\\simeq$ 100 km s$^{-1}$). The single-peaked CO profiles of both the nuclear and off-nuclear regions may be explained if we introduce a CO gas disk with a velocity dispersion of $\\sim 100$ km s$^{-1}$. If this gas disk would be extended up to a few kpc in radius, we may explain the wide line widths of the off-nuclear CO emission. Alternatively, we may attribute the off-nuclear CO emission to the gas associated with the superwind. However, if all the CO gas moves along the biconical surface of the superwind, the CO spectra would show double-peaked profiles. Hence, the single-peaked CO profiles of the off-nuclear regions may be explained by an idea that the morphology and/or velocity field  of the molecular-gas superwind are more complex as suggested by hydrodynamical simulations. ",
        "introduction": "In starburst galaxies, a large number of massive stars (e.g., $\\sim 10^{4-5}$) are formed within a short duration (Weedman et al. 1981; Balzano 1983; Taniguchi et al. 1988). Therefore, a burst of supernova explosions occurs inevitably $\\sim 10^7$ years after the onset of the starburst. Since these numerous supernovae release a huge amount of kinetic energy into the circumnuclear gas, the circumnuclear gas is thermalized and then blow out into the direction perpendicular to the galactic disk as a ``superwind'' (Tomisaka \\& Ikeuchi 1988; Heckman, Armus, \\& Miley 1990; Suchkov et al. 1994). A bubble of the ionized gas sweeps up the circumnuclear molecular gas, leading to the formation of molecular-gas superwind as well as the ionized-gas one (Tomisaka \\& Ikeuchi 1988; Suchkov et al. 1994). Thus, in order to understand the whole physical processes of superwinds, it is important to investigate the nature of molecular-gas superwinds (e.g., Nakai et al. 1987; Aalto et al. 1994; Irwin \\& Sofue 1996).  In this {\\it Letter}, we present new evidence for the molecular-gas superwind from the Wolf-Rayet starburst galaxy Mrk 1259, which shows the optical evidence for the superwind viewed from a nearly pole-on view (Ohyama, Taniguchi, \\& Terlevich 1997; hereafter Paper I). Mrk 1259 is a peculiar S0 galaxy (de Vaucouleurs et al. 1991; hereafter RC3) at a distance of 26.64 Mpc \\footnote{Paper I adopted a distance toward Mrk 1259, $D = 33.5$ Mpc. However, $V_{\\rm 3K}$ was misused instead of $V_{\\rm GSR}$ in this estimate. In this {\\it Letter}, using $V_{\\rm GSR} = 1998$ km s$^{-1}$ (RC3), with a Hubble constant $H_0$ = 75 km s$^{-1}$ Mpc $^{-1}$, we adopt a distance $D$ = 26.64 Mpc. Therefore, the HeII$\\lambda$4686 luminosity, the number of late WR (WRL) stars, the size of the superwind, and the average velocity of the superwind in Paper I should be read as $L$(HeII) = 7.0$\\times 10^{39}$ erg s$^{-1}$, $N$(WRL) $\\simeq$ 4100, $r$(superwind) $\\simeq 3.3$ kpc, and the average wind velocity $\\simeq$ 565 km s$^{-1}$, respectively.}. The logarithmic major-to-minor diameter ratio, log $R_{\\rm 25}=0.10\\pm 0.08$ (RC3), gives a nominal inclination angle, $i=37\\fdg 4^{+11.2}_{-20.1}$, and the galaxy appears to be elongated along the EW direction. If this elongation were attributed to the inclination, we would observe the rotational motion along the EW direction. However, our long slit optical spectrum along the EW direction which was analyzed in Paper I shows no hint on the rotational motion; $\\Delta V\\lesssim 50$ km s$^{-1}$, suggesting strongly that the galaxy is seen from an almost face-on view. Therefore the oval shape of Mrk 1259 may not be due to the inclination\\footnote{It seems no surprise even if an isolated galaxy shows some morphological peculiarity because any galaxy would experience some minor merger events in its life. It is also noted that minor mergers can cause nuclear starbursts (e.g., Hernquist \\& Mihos 1995; Taniguchi \\& Wada 1996).}. ",
        "conclusions": "In Table 1, we give a summary of our observational results. The integrated CO intensity was estimated by $I({\\rm CO}) = \\int T_{\\rm A}^* \\eta_{\\rm mb}^{-1} dv$ K km s$^{-1}$ where $\\eta_{\\rm mb} = 0.51$. Using a galactic conversion factor, $N_{\\rm H_{2}}/I_{\\rm CO} = 3.6\\times 10^{20}$ cm$^{-2}$ (K km s$^{-1}$)$^{-1}$ (Scoville et al. 1987), we estimate the molecular gas mass, $M_{\\rm H_2} = 5.8 \\times 10^6 I({\\rm CO}) A$, in each position where $A$ is the projected area of a 15$^{\\prime\\prime}$ HPBW in units of kpc$^2$. For the off-nuclear regions, we also give total values of $I$(CO) and $M_{\\rm H_2}$. The total molecular gas mass detected in our observations amounts to $1.2 \\times 10^9 M_\\odot$. Since we do not observe the entire disk of this galaxy, this mass is regarded as a lower limit.  \\subsection{The Nuclear CO Emission}  The nuclear CO emission shows a single-peaked profile with the evident blueward asymmetry. Applying a two-component Gaussian profile fitting (see the midst panel of Figure 1), we obtain the blueshifted broad component with FWHM $\\simeq$ 114 km s$^{-1}$ and the narrow one with FWHM $\\simeq$ 68 km s$^{-1}$. The peak velocity of the broad component is blueshifted by 27 km s$^{-1}$ with respect to that of the narrow one (Table 1). Both the intensities are nearly the same. We also mention that the red wing cannot be seen in the nuclear CO profile.  Even though there is the broad CO emission component, its width is significantly narrower than those observed for typical starburst galaxies; e.g., FWHM(CO) $\\simeq$ 200 - 250 km s$^{-1}$ for M82 (Young \\& Scoville 1984; Nakai et al. 1987), $\\sim 350$ km s$^{-1}$ for NGC 1808 (Aalto et al. 1994), and $\\sim 325$ km s$^{-1}$ for NGC 4945 (Dahlem et al. 1993). This difference can be attributed to the effect of viewing angles between Mrk 1259 and the other starburst galaxies. It is remembered that the CO line width is generally affected by the galactic rotation. Given a typical rotation velocity of a disk galaxy, $V_{\\rm rot} \\sim 200$ km s$^{-1}$, the observed full widths would amount to 2$V_{\\rm rot} \\sim$ 400 km s$^{-1}$ if seen from the edge-on view. In fact, since we observe M82, NGC 1808, and NGC 4945 from highly inclined viewing angles, their line widths are considered to be broadened by the effect of galactic rotation. On the other hand, since Mrk 1259 appears to be a nearly face-on galaxy, the observed width is not affected by the galactic rotation. Irwin \\& Sofue (1996) suggested that one of the nearby superwind galaxies, NGC 3628, has a nuclear molecular gas disk with a velocity dispersion of $\\sim$ 100 km s$^{-1}$. If Mrk 1259 has also such a nuclear gas disk, we can explain the velocity width of the nuclear CO emission. Therefore, it is suggested that the observed FWHM of Mrk 1259 is due mainly to the broadening by some dynamical effect of the starburst activity. The blueward asymmetry of the nuclear CO line profile suggests that the CO gas is affected significantly by the superwind.  \\subsection{The Off-Nuclear CO Emission}  The detection of the off-nuclear CO emission from Mrk 1259 is very intriguing from the following two points. The first point is that the host galaxy of Mrk 1259 appears to be an S0 galaxy (RC3). It is often observed that early type galaxies such as S0 and elliptical galaxies tend to have less molecular gas (e.g., Young \\& Scoville 1991) although CO emission has been detected from a number of S0 galaxies (Thronson et al. 1989; Wiklind \\& Henkel 1989; Sage 1989; Sage \\& Wrobel 1989). It is also known that the molecular gas in (non-active) S0 galaxies tends to be concentrated in the region whose diameter is typically less than one tenth of the optical diameter (Taniguchi et al. 1994). If this is also the case for Mrk 1259, the molecular gas would be concentrated within the central $12\\arcsec =0.1 D_{\\rm 0}$ region where $D_{\\rm 0}$ is the isophotal optical diameter (RC3). Therefore, the presence of the bright off-nuclear CO emission is one of very important characteristics of Mrk 1259.  The second point is that the line widths of the off-nuclear CO emission are comparable to that of the nuclear CO emission, FWHM $\\sim 100$ km s$^{-1}$. If there were an inclined off-nuclear CO disk, we may explain the wide line width because of the velocity gradient in the disk. If this is the case, we would observe that the peak velocity at 15$^{\\prime\\prime}$E is significantly different from that at 15$^{\\prime\\prime}$W. However, since our observations show that the velocity field of the off-nuclear regions is almost symmetric, this possibility is rejected. The second possibility is that there are spatially extended starburst regions and a significant amount of molecular gas is associated with them. However, radio continuum (1.5 GHz and 5 GHz) images show that the starburst region of Mrk 1259 is concentrated in the central several arcsec region (R. A. Sramek 1997, private communication). Therefore, there is no observational evidence for active star forming regions in the off-nuclear regions. As described before, we are observing the disk of Mrk 1259 from nearly a face-on view and thus the CO line width would be as narrow as $\\sim$ 10 km s$^{-1}$ if Mrk 1259 were a normal disk galaxy (Lewis 1984, 1987; Kamphuis \\& Sancisi 1993). If there were an extended molecular gas disk with a velocity dispersion of $\\sim$ 100 km s$^{-1}$ up to a radius of a few kpc, we could explain the wide line width. However, the size of the nuclear gas disk in NGC 3628 is much smaller ($\\simeq 230$ pc, or $\\sim 0.01 D_{\\rm 0}$) than the off-nuclear distance of Mrk 1259 ($\\sim 2$ kpc, or $\\sim 0.13 D_{\\rm 0}$). Although we cannot rule out the possibility that Mrk 1259 has such a very extended molecular gas disk with a large velocity dispersion, we need further detailed molecular-line observations to confirm this possibility. The third possibility is that the off-nuclear CO gas is associated with the superwind (i.e., blown out from the nuclear region). Since the ionized-gas superwind is extended to $r \\sim 3.3$ kpc (Paper I), this possibility seems to be quite high. In fact, such extended CO emission is detected in M82 at the scale of 600 pc (Nakai et al. 1987) and even at the larger scale ($\\sim 2$ kpc; Sofue et al. 1992). We discuss this possibility in detail in the next section.  \\subsection{Biconical Superwind Model for Mrk 1259}  Since the superwind of Mrk 1259 is observed from nearly the pole-on view, it is interesting to investigate both the velocity field and the geometry of the superwind. In order to perform this, we investigate the off-nuclear CO line profile using a simple biconical outflow model in which the superwind flows toward the polar directions symmetrically with its apex at the nucleus. Such a superwind geometry is expected theoretically by hydrodynamical numerical simulations (Suchkov et al. 1994) and indeed observed in M82 (e.g., Nakai et al. 1987). In our model, we assume that the molecular gas can only move along the cone surface. We assume that the axis of the cone lies along our line of sight. The full opening angle of the cone ($\\theta$) is not well constrained by the observations because of its nearly face-on viewing angle. Therefore we take this as a free parameter although Paper I has suggested as $\\theta \\lesssim 90\\arcdeg$. A mean tangential velocity of the ionized gas on the sky can be estimated as $V_{\\rm t, ion}\\simeq R_{\\rm SW}/T_{\\rm SW}\\simeq (2.3$ kpc$) /(5.5\\times 10^6$ years)$\\simeq 410$ km s$^{-1}$ where $R_{\\rm SW}$ is the projected radius of the superwind and $T_{\\rm SW}$ is the age of the superwind (Paper I). We note that the outflow velocity of the molecular gas is {\\it slower} than that of the ionized gas because the molecular gas along the cone surface is {\\it dragged} by the ionized gas, rather than directly {\\it pushed out} (Suchkov et al. 1994). For example, the model A1 of Suchkov et al. (1994) shows that the velocity of the outflowing dense gas is slower by a factor of $\\sim 5$ than that of the ionized gas at the age of 8.3 Myr. In fact, comparing the outflow velocity of the ionized gas (Heckathorn 1972) with that of the molecular gas (Nakai et al. 1987) of M82, we find that the outflow velocity of the molecular gas is slower by a factor of $\\sim 3$ than that of the ionized-gas. Thus, the mean tangential velocity of the molecular gas on the sky can be $V_{\\rm t, mol}=V_{\\rm t, ion}/\\epsilon$ where $\\epsilon$ is the decelerating factor ($\\epsilon \\simeq 3 - 5$).  We examine if the model can explain the observed CO line profiles in the off-nuclear regions. No effect of radiative transfer is included in the model calculation. We assume that the size of the cone is large enough to cover the whole off-nuclear regions. For simplicity, we also assume that the velocity field has a power-law form; i.e., $V(r) \\propto r^a$, with a boundary condition of $V_{\\rm t, ion}$ ($r = 2.3$ kpc) = 410 km s$^{-1}$. The emissivity (strength of the CO emission per a unit area) is also assumed to have a power-law form; i.e., $I(r) \\propto r^b$. Although the parameters $a$ and $b$ are not well constrained by the observations, we adopt $a = 1$ and $b = -1$ as representative values following the trend seen in M82 (Nakai et al. 1987). We calculate the model for the cases of $\\theta = 60\\arcdeg, 90\\arcdeg, 120\\arcdeg$, and $150\\arcdeg$ and $\\epsilon$ = 1, 2, 3, 4, 5, and 6. To explain the observed FWZI (Full Width at Zero Intensity) of the off-nuclear CO emission ($\\sim 200$ km s$^{-1}$; see Figure 1), we find that only models with ($\\theta = 90\\arcdeg$ and $\\epsilon \\simeq 5 - 6$) and ($\\theta = 120\\arcdeg$ and $\\epsilon \\simeq 3 - 4$) are acceptable. Models with $\\theta = 60\\arcdeg$ and $150\\arcdeg$ cannot reproduce the line width for any $\\epsilon$. Therefore, we show our results only for the cases $\\theta=90\\arcdeg$ and $120\\arcdeg$ and $\\epsilon$ = 3, 4, 5, and 6 in Figure 2. Our simple model demonstrates that the CO line has always a double-peaked profile for any combinations of the parameters. The red peak corresponds to the recessing cone while the blue one corresponds to the advancing cone. On the other hand, our observations show that the off-nuclear CO lines have smooth profiles around at the systemic velocity.  We discuss why our simple model cannot reproduce the observed off-nuclear CO profiles. One possible idea is a ``swirl''-like velocity field which is often found in the hydrodynamical numerical simulations (Tomisaka \\& Ikeuchi 1988; Suchkov et al. 1994). If the outflow actually shows such a complex geometry and/or velocity field, it is expected that some parts of the emission would contribute to the core emission and can explain the broad and smooth off-nuclear emission. In order to understand the molecular-gas superwind of Mrk 1259, detailed molecular-line observations with higher spatial resolution would be helpful.   \\vspace{0.5cm}  We would like to thank the staff of Nobeyama Radio Observatory for their kind support for our observations. We thank Naomasa Nakai for useful discussion and encouragement and Takashi Murayama and Shingo Nishiura for kind assistance of the observations. We also thank R. A. Sramek for kindly providing us his VLA data. YO was supported by the Grant-in-Aid for JSPS Fellows by the Ministry of Education, Culture, Sports, and Science. This work was financially supported in part by Grant-in-Aids for the Scientific Research (No. 0704405) of the Japanese Ministry of Education, Culture, Sports, and Science.  \\newpage  \\begin{table} \\caption{Molecular gas properties of Mrk 1259} \\begin{tabular}{llccc} \\tableline \\tableline & & & Nucleus & Off-nucleus \\\\ \\tableline rms noise & $\\delta T_{\\rm A}^*$ (K) & & 0.018 & $\\sim 0.015$ \\\\ Flux & $I$(CO)$^{a, b}$ (K km s$^{-1}$) & total & $28.0\\pm 1.1$ & $43.8\\pm 1.9$$^d$ \\\\ & & 15\\arcsec N & & $4.2\\pm 1.0$ \\\\ & & 15\\arcsec E & & $18.6\\pm 1.0$  \\\\ & & 15\\arcsec S & & $9.5\\pm 1.0$ \\\\ & & 15\\arcsec W & & $11.5\\pm 1.0$ \\\\ Mass & $M_{\\rm H_2}^{b, c}$ ($M_\\odot$) & & $(4.8\\pm 0.2) \\times 10^8$ & $(7.5\\pm 0.3) \\times 10^8$$^d$ \\\\ Line profile & & & & \\\\ & FWHM$_{\\rm narrow}$ (km s$^{-1}$) & & 68 & \\nodata \\\\ & $V_{\\rm narrow}$ (km s$^{-1}$) & & 2178 & \\nodata \\\\ & FWHM$_{\\rm broad}$ (km s$^{-1}$) & & 114 & $\\sim 100$$^e$ \\\\ & $V_{\\rm broad}$ (km s$^{-1}$) & & 2151 & $\\sim 2170$$^e$ \\\\ & $I$(broad)/$I$(narrow) & & 1.0 & \\nodata \\\\ \\tableline \\end{tabular} \\tablenotetext{a}{$I({\\rm CO}) = \\int T_{\\rm A}^* \\eta_{\\rm mb}^{-1} dv$ K km s$^{-1}$ where $\\eta_{\\rm mb}$ = 0.51.} \\tablenotetext{b}{Formal one-sigma error indicated.} \\tablenotetext{c}{A Galactic conversion factor, $N_{\\rm H_{2}}/I_{\\rm CO} = 3.6\\times 10^{20}$ cm$^{-2}$ (K km s$^{-1}$)$^{-1}$ (Scoville et al. 1987), is assumed.} \\tablenotetext{d}{Sum of the fluxes of all the four off-nuclear regions.} \\tablenotetext{e}{No profile fitting was made because of both the lower signal-to-noise ratio and the non-Gaussian profiles. The values shown here are typical ones for the off-nuclear regions except 15$^{\\prime\\prime}$N.} \\end{table}  \\newpage"
    },
    "9803/gr-qc9803068_arXiv.txt": {
        "abstract": "We propose two new classes of instantons which describe the tunneling and/or quantum creation of closed and open universes.  The instantons leading to an open universe can be considered as generalizations of the Coleman-De-Luccia solution. They are non-singular, unlike the instantons recently studied by Hawking and Turok,  whose prescription has the problem that the singularity is located on the hypersurface connecting to the Lorentzian region, which makes it difficult to remove. We argue that such singularities are harmless if they are located purely in the Euclidean region. We thus obtain new singular instantons leading to a closed universe; unlike the usual regular instantons used for this purpose, they do not require complex initial conditions.  The singularity gives a boundary contribution to the action which is small for the instantons leading to sufficient inflation, but changes the sign of the action for small $\\phi$ corresponding to short periods of inflation. ",
        "introduction": "\\label{sec-nonsing}   Suppose we have an effective potential $V(\\phi)$ with a local minimum at $\\phi_1$, and a global minimum at $\\phi=0$, where $V=0$ (see Fig. \\ref{potential}).  In an $O(4)$-invariant Euclidean spacetime with the metric \\begin{equation}\\label{metric} ds^2 =d\\tau^2 +a^2(\\tau)(d \\psi^2+ \\sin^2 \\psi \\, d \\Omega_2^2) \\ , \\end{equation} the scalar field $\\phi$ and the three-sphere radius $a$ obey the equations of motion \\begin{equation}\\label{equations} \\phi''+3{a'\\over a}\\phi'=V_{,\\phi},~~~~~ a''= -{8\\pi G\\over 3} a ( \\phi'^2 +V) \\ , \\end{equation} where primes denote derivatives with respect to $\\tau$.  \\begin{figure}[Fig0] \\hskip 1.5cm \\leavevmode\\epsfysize=4cm \\epsfbox{Potential.eps}  \\  \\caption[Fig1]{\\label{potential} Effective potential $V(\\phi) = {m^2\\over 2}(\\phi^2(\\phi- v)^2 +B \\phi^4)$ for $m^2 = 2$, $B=0.12$ and $v = 0.5$.  It has a shallow minimum at   $\\phi_0 = 0.357$ and a local maximum at $\\phi_1=0.312$. All quantities in this figure are in units of $M_{\\rm p}/\\sqrt{8\\pi}$.} \\end{figure}  These equations have several non-singular solutions, the simplest of which are the $O(5)$ invariant four-spheres one obtains when the field $\\phi$ sits at one of the extrema of its potential. In this case the first of the two equations above is trivially satisfied, and $a(\\tau) = H^{-1} \\sin H\\tau$. Here $H^2 = {8\\pi V\\over 3 M_{\\rm p}^2}$. Using the solution for which $\\phi=\\phi_1$, Hawking and Moss \\cite{HM} found the rate at which the field $\\phi$ in a single Hubble volume tunnels to the top of the potential, from which it can roll down towards the true vacuum.  For a recent discussion of this instanton and its interpretation see \\cite{ALOpen}. The main other use of these trivial instantons is to find the action of the false vacuum background solution, which must be subtracted from the bounce action to obtain a tunneling rate.  We shall consider potentials for which $V_{,\\phi\\phi} \\gg H^2$ in the region where the tunneling occurs. In this case, tunneling out of the false vacuum does not occur primarily on the scale of an entire Hubble volume via the Hawking-Moss instanton. Instead the transition will proceed via more complicated Euclidean solutions with varying field $\\phi$. These include the  Coleman-De-Luccia instanton, and related instantons which we found.   \\subsection{Bubble instantons}  A Euclidean solution which describes the creation of an open universe was first found by Coleman and De Luccia in 1980 \\cite{CL}.  It is given by a slightly distorted de~Sitter four-sphere of radius $H^{-1}(\\phi_0)$. Typically, the field $\\phi$ is very close to the false vacuum, $\\phi_0$, throughout the four-sphere except in a small region (whose center we may choose to lie at $\\tau=0$), in which it lies on the `true vacuum' side of the maximum of $V$. The behavior of the field and scale factor for the potential in Fig.~\\ref{potential} is shown in Fig.~\\ref{Colem}. The scale factor vanishes at the points $\\tau=0$ and $\\tau=\\tau_{\\rm f} \\approx \\pi/H$, which we will call the North and South pole of the four-sphere. In order to get a singularity-free solution, one must have $\\phi' = 0$ and $a'=\\pm 1$ on the poles.  This solution can be cut in half along the line $\\psi=\\pi/2$, which removes half of each three-sphere. Then one can continue analytically to a Lorentzian spacetime~\\cite{GutWei83,HT} with the time variable $\\sigma$, given by $\\psi=\\pi/2+i\\sigma$. This gives region II of the Lorentzian universe (see Fig.~\\ref{fig-regions}): \\begin{equation} ds^2 = -a^2(\\tau)\\ d\\sigma^2 + d\\tau^2 + a^2(\\tau) \\cosh^2 \\sigma\\ d\\Omega_2^2; \\end{equation} the field $\\phi$ will still depend on $\\tau$ in the same way as before, and will be independent of $\\sigma$.  This describes a shell of width $H^{-1}$, which is mostly near the false vacuum and expands exponentially. The shell separates two bubbles, regions I and III, in which the universe looks open.     \\begin{figure}[Fig1] \\hskip 1.5cm \\leavevmode\\epsfysize=9.5cm \\epsfbox{Coleman.eps}  \\  \\caption[Fig1]{\\label{Colem} The upper panel shows the behavior of the scalar field $\\phi$ for examples of the Coleman-De-Luccia ``bubble'' instanton (solid line) and the new ``double-bubble'' instanton which we have found (dashed line).  For both instantons, the field is in the domain of the true vacuum at small $\\tau$, forming a bubble.  For the bubble instanton, the field is closest to the false vacuum at the pole opposite the bubble. For the double-bubble instanton, this happens on the equator, at the moment of the maximal expansion. The behavior of the three-sphere radius $a(\\tau)$ shown in the lower panel is very similar for both instantons, though it is not identical.}  \\end{figure}  Region I is obtained by taking $\\sigma = i\\pi/2 + \\chi$ and $\\tau = it$, giving the metric \\begin{equation} ds^2 = -dt^2 + \\alpha^2(t) \\left( d\\chi^2 + \\sinh^2 \\chi d\\Omega_2^2 \\right), \\end{equation} where $ \\alpha(t) = -i\\, a[\\tau(t)]$. Its spacelike sections (defined by the hypersurfaces of constant inflaton field) are open. Thus, region I looks from the inside like an infinite open universe, which inflates while the field $\\phi$ slowly rolls down to the true vacuum. The evolution will then undergo a transition to a radiation or matter-dominated open Friedman-Robertson-Walker universe.   In region III, which is obtained by choosing $\\sigma = i\\pi/2 + \\chi$ and $\\tau = \\tau_{\\rm f} + it$, the field $\\phi$ rolls to the local minimum at $\\phi_0$, and one gets indefinite open inflation in the false vacuum.    \\begin{figure}[Fig1] \\hskip 1.5cm \\leavevmode\\epsfysize=5cm \\epsfbox{continuation.eps}  \\  \\caption[Fig1]{\\label{fig-regions} The Lorentzian de~Sitter-like spacetime obtained from the analytic continuation of Coleman-De-Luccia instantons contains three regions. In Regions I and III the hypersurfaces of constant field $\\phi$ form open spacelike sections. Region II is a shell separating the two bubbles.}  \\end{figure}   The analytic continuations we have given support the interpretation of such solutions as the spontaneous nucleation of a bubble of true vacuum on the background of de~Sitter space expanding in the false vacuum. For this reason we will call them `bubble instantons'.  The nucleation rate is given by \\begin{equation} \\Gamma = e^{-\\Delta S}, \\end{equation} where $\\Delta S$ is the difference between the action of the full Euclidean bubble solution, and the action of a Euclidean solution describing the background spacetime.  Except for near-Planckian potentials, both actions will be large and negative (about   $-2.6 \\times 10^4$ in our example).  The background solution is given by an exact Euclidean four-sphere on which the field $\\phi$ is constant and equal to $\\phi_0$, the false vacuum. Its action will be $-{3 M_{\\rm p}^4 \\over 8 V(\\phi_0)}$.  Subtracting   this from the action of the bubble solution, one obtains a positive $\\Delta S$ ($\\approx 4.9$ in our example). This means that bubble formation by tunneling is suppressed, as it should be.        One usually requires instanton solutions to interpolate between the initial and final spacelike sections (in this case, a section of pure de~Sitter space in the false vacuum and a similar section containing a bubble of true vacuum). The above description, which seems to use two disjoint instantons, is actually consistent with this formal requirement, since the instantons may be connected by virtual domain walls after small (Planck size) four-balls are removed. This will cause the background instanton to contribute to the total action with a negative sign. If one connects the background instanton to the region of the bubble instanton where $\\phi$ is closest to its false vacuum, the discontinuity in $\\phi$ will be small, so the volume contributions of the removed regions cancel almost exactly. Requiring continuous instantons, therefore, does not change the pair creation rate significantly~\\cite{BC}.  Cosmological instantons have frequently been interpreted to describe the creation of a universe from nothing, i.e.\\ without a pre-existing background. This case is considerably less well-defined than the quantum nucleation of structures on a given background solution. In particular, the sign with which the large, negative action enters the exponent in the path integral is subject to debate~\\cite{HT,ALOpen,HTnew}. Leaving such questions aside for now, we will take the position that isolated cosmological instantons are indeed related to universe creation, independently of the formalism used to assign probabilities to such processes.   \\subsection{Double-bubble instantons}  We have found a new instanton in which there are two bubbles, one on each pole. In this solution, $\\phi$ is in the domain of the true vacuum in small regions near the poles, and near the false vacuum elsewhere; this can be seen from the dashed line in Fig.~\\ref{Colem}. The geometry is still approximately a four-sphere. As before, $\\phi'$ vanishes on the poles; but now it also vanishes on the equator, at $\\tau=\\tau_{\\rm max}$. The Northern and Southern hemispheres are exactly symmetric.  Not surprisingly, the action of the double-bubble solution, after the background subtraction described above, is approximately twice that of the bubble (Coleman-De-Luccia) instanton.  For the instanton shown in Fig.~\\ref{Colem} one has $\\Delta S_2 \\approx 9.8$.  The analytic continuations will be the same as before, with a different result. Region II will be mostly in the domain of the false vacuum. Region I and III will be identical, each containing an open inflating universe in which the field rolls down to the true vacuum. Globally, therefore, we obtain two bubbles of true vacuum separated by a shell which inflates in the false vacuum.  This solution can be interpreted as the spontaneous pair-creation of bubbles of open inflation on the background of false vacuum inflation. Alternatively, one may view it as the creation from nothing of two open inflating universes separated by a metastable shell.    \\subsection{Anti-double-bubble instantons}  In addition we have found another family of instantons, two examples of which are shown in Fig.~\\ref{loop}. In these instantons, the field is in the domain of the false vacuum in two regions surrounding the poles. They are separated by a thin shell at the equator, where the field is in the true vacuum domain. These instantons have a much greater action difference to the background instanton, since the true-vacuum region is significantly larger than in the previous two cases. In particular, $\\Delta S = 93.6$ for the instanton shown by the solid line in Fig.~\\ref{loop}, and $\\Delta S = 124.7$ for the instanton shown by the dashed line.  \\begin{figure}[Fig0111] \\hskip 1.5cm \\leavevmode\\epsfysize=4.5cm \\epsfbox{NewInst.eps}  \\  \\caption[Fig1]{\\label{loop} Two examples of ``anti-double-bubble'' instantons, in which the field is in the false vacuum domain near the poles, and reaches into the domain of the true vacuum on a shell near the equator. It can be cut through the poles to describe shell nucleation, or across the equator, describing the tunneling to true-vacuum inflation in a closed universe.}  \\end{figure}  \\subsubsection{Open cut}  With the analytic continuation used for the previous two instantons, regions I and III will become open inflationary universes in which the field rolls down to the false vacuum. They are separated by the region II, which contains a shell on which $\\phi$ is in the domain of the true vacuum.  Therefore we may interpret this solution as the nucleation of a shell of true vacuum on a false vacuum inflationary background, or alternatively, as the creation of such a universe from nothing. Because of the larger action difference, spontaneous shell creation will be quite suppressed compared to bubble formation.  \\subsubsection{Closed cut}  A more intriguing application of this instanton can be found by choosing a different analytic continuation. Instead of cutting at $\\psi=\\pi/2$, we may choose to leave the three-spheres intact, and cut across the equator. Lorentzian time will be defined by $\\tau=\\tau_{\\rm max} + iT$, and we obtain a metric with closed spacelike sections: \\begin{equation} ds^2 = -dT^2 + a^2(T) d\\Omega_3^2. \\label{eq-closed} \\end{equation} The inflaton field is in the domain of the true vacuum on the nucleation surface (the equator), so it will start rolling down towards the absolute minimum. During this time, the spacelike three-spheres grow exponentially: \\begin{equation} a(T) \\approx H^{-1}(T) \\cosh \\int H(T)\\, dT. \\label{eq-cosh} \\end{equation} Thus we obtain a closed inflationary universe in which the scalar field rolls towards the true vacuum.  One could interpret this instanton as describing the creation of such a universe from nothing. But this would just add an alternative to the usual instantons on which the field is entirely in the domain of the true vacuum. A much more interesting interpretation is the one associated with a pre-existing background of false vacuum inflation. In this case, the anti-double-bubble instanton is seen to describe the spontaneous tunneling to the true vacuum in an entire Hubble-volume of de~Sitter space. Unlike the Coleman-De-Luccia bubbles, these regions will not contain an open universe. In the example we are considering, for which Hawking-Moss tunneling is not possible, this shows that one can nevertheless nucleate true vacuum bubbles containing a closed universe. ",
        "conclusions": "We have described a number of non-singular instantons leading to open inflating universes. They include the Coleman-De-Luccia solution, in which a bubble of true vacuum expands inside a universe inflating in the false vacuum. We found new solutions which contain two bubbles, or a shell of true vacuum.  We also constructed instantons with a singularity. If the singularity does not lie on the hypersurface of nucleation, it causes no problems in the Lorentzian region, and can be interpreted as a small region of Planckian density.  Such instantons can be used to describe the quantum creation of a closed inflationary universe from space-time foam without the need to use complex solutions.       \\subsection*"
    },
    "9803/astro-ph9803107_arXiv.txt": {
        "abstract": "Using time resolved 2-dimensional aperture photometry we have established that the optical candidate for PSR 1509-58 does not pulse. Our pulsed upper limits ($m_V$ = 24.3 and $m_B$ = 25.7) put severe constraints on this being the optical counterpart. Furthermore the colours of the candidate star are consistent with a main sequence star at a distance of 2-4 kpc. The probability of a chance coincidence with a normal star and the difficulty of explaining the lack of pulsed emission leads us to conclude that this object is an intermediate field star. ",
        "introduction": "Interest in the optical emission from isolated neutron stars (INSs) has been growing, as recent improvements in detector sensitivity have enabled these faint sources to be observed. Optical observations of neutron stars are important for providing an understanding of pulsar emission mechanisms and allowing direct observations of the energy spectrum of the electron pair plasma.   To date 6 INSs have been detected in the optical. Evidence suggests that it is the age and/or the period derivative of the INS, rather than its period, that determines the  optical, and indeed multiwavelength, emission from an INS (\\cite{car94a},  \\cite{gol95}).  PSR 1509-58 was initially identified by its X-ray emission (\\cite{sew82}) and shortly afterwards a radio signal, with a period of 150ms, was detected (\\cite{man82}). Later studies showed that it had a large $\\dot P$ (1.5 x 10$^{-12}$ ss$^{-1}$ ), in fact the largest known. The pulsar has had its second period derivative measured giving a breaking index ($n=\\omega {\\ddot \\omega}/{\\dot \\omega^2}$  ) of 2.83 $\\pm 0.03$ (\\cite{man85}) in close agreement with the expectations of a radiating dipole (n=3). Its magnetic field ($\\propto (P \\dot P)^{1/2}$) is the largest known and its age 1,600, second only to the Crab pulsar in youth. Its age and location make its association with SNR MSH 15-52 (\\cite{sew84}) at a distance of 4.2 kpc (\\cite{cas75}) likely, although this position has been challenged by \\cite{str94}, who has proposed a greater distance, 5.9 kpc, based upon radio dispersion and x-ray spectra of the extended emission around the pulsar.  It has also been observed in soft gamma-rays (\\cite{gun94}) and a tentative optical counterpart has been proposed with $m_V \\approx 22$ (\\cite{car94b}). When taken at a distance of 4.2 kpc the optical observations indicate an absolute magnitude of M$_V$$\\approx 4.9$ (including the effects of interstellar extinction), fainter than the Crab pulsar. However, it is much brighter than would be expected from phenomenological models which have been successful in describing the X-ray emission from PSR 1509-58 (Pacini and Savati 1987). This over-luminosity makes its behaviour similar to the older optical pulsars, PSR0656+14 (\\cite{shear97a}) and Geminga (\\cite{shear97b}). Observations of any pulsed optical component will be crucial in determining what fraction of the emission is thermal and what is magnetospheric. Only the magnetospheric emission would be expected to scale in a manner analogous to that described by Pacini and Savati. Indeed, by considering plausible emission mechanisms (Lu et al 1994) and high energy observations (\\cite{gol95}), it would be reasonable to expect that most of the emission would be non-thermal.   Given the  importance of a correct determination of the optical emision from PSR1509-58, we have made time-resolved observations in an attempt to confirm its identification and determine the optical pulsed fraction. ",
        "conclusions": "Our dervived magnitudes and colours, when combined with the lack of optical pulsations, cast doubt on the Caraveo et al (1994b) candidate being the optical counterpart of PSR 1509-58. Our data is consistent with an M type main sequence star at a distance of about 2 kpc. Alternatively, if it is the pulsar then the pulsed fraction must be anomolously low (lower than any other optical pulsar) and most of the radiation thermal. This is contrast with higher energy observations (dominated by non-thermal emission), where the pulsed fraction increases with decreasing energy (\\cite{gre95}). However if the optical emission represents the Rayleigh-Jeans tail of the neutron star's black body spectrum, then with tabulated values for extinction towards MSH15-52, the distance would have to be $\\sim$ 1 kpc assuming a surface temperature of $3~10^6$ K  in contrast to the expected distance of at least 4.2 kpc. The presence of a neutron star atmosphere (\\cite{pav96}) would not change this distance sufficiently to explain the optical excess. If the radiation is non-thermal then it is difficult to imagine a geometry and a mechanism which would give such an anomolously low pulsed fraction and still be consistent with high energy observations. Even at the distance derived from radio dispersion (5.9 kpc) the star is still too red to be explained by a thermal extrapolation alone. A final answer to the nature of this star will come from spectroscopy - its magnitude is well within the capabilities of the VLT."
    },
    "9803/astro-ph9803277_arXiv.txt": {
        "abstract": " ",
        "introduction": "The observed present-day abundance of rich clusters of galaxies places a strong constraint on cosmology: \\gs$\\Omega^{0.5} \\simeq 0.5$, where \\gs\\ is the {\\em rms} mass fluctuations on 8 \\gh\\ Mpc scale, and \\gW\\ is the present cosmological density parameter (Henry \\& Arnaud 1991, Bahcall \\& Cen 1992, White \\gE\\ 1993, Eke \\gE\\ 1996, Viana \\& Liddle 1996, Pen 1997, Kitayama \\& Suto 1997). This constraint is degenerate in \\gW\\ and \\gs; models with \\gW =1, \\gs \\gi 0.5 are indistinguishable from models with \\gW \\gi  0.25, \\gs \\gi 1. (A \\gs \\gie 1 universe is unbiased, with mass following light on large scales; \\gs \\gie 0.5 implies a mass distribution  wider than light).  The {\\em evolution} of cluster abundance with redshift, especially for massive clusters, breaks the degeneracy between \\gW\\ and \\gs\\ (see, e.g., Peebles et al.  1989, Oukbir \\& Blanchard 1992, 1997,  Eke \\gE\\ 1996, Viana \\& Liddle 1996, Carlberg \\gE\\ 1997, Bahcall \\gE\\ 1997, Fan \\gE\\ 1997, Henry 1997). The evolution of high mass clusters is strong in \\gW\\ =1, low-\\gs\\ (biased) Gaussian models, where only a very low cluster abundance is expected at $z>$0.5. Conversely, the evolution rate in low-\\gW\\ , high-\\gs\\ models is mild and the cluster abundance at $z>$0.5 is much higher than in \\gW=1 models. In Bahcall \\gE\\ (1997) and Fan \\gE\\ (1997) we used the CNOC cluster sample (Carlberg \\gE\\ 1997a,b, Luppino \\& Gioia 1995) to $z \\lesssim$ 0.5 -- 0.8 (with measured masses to $z \\lesssim 0.5$) to decouple \\gW\\ and \\gs: we found \\gW \\gie  0.3 \\gm 0.1 and \\gs \\gie 0.83 \\gm 0.15, consistent with Carlberg \\gE (1997a). The evolution rate, and the distinction among cosmological models, increases with cluster mass and with redshift: in \\gW\\ =1, low-\\gs\\  models, very massive clusters are not expected to exist at high redshifts.  In the present paper we extend the previous studies to larger mass and higher redshift clusters, using the three most massive clusters observed so far at high redshifts ($z$ \\gie 0.5--0.9) to independently constrain \\gW\\ and \\gs. The clusters discussed in this paper are the three most massive distant clusters from the EMSS/CNOC sample used above, with masses larger by a factor of $\\sim$ 2 than the mass-threshold used previously (Evrard 1989, Bahcall \\gE\\ 1997, Fan \\gE\\ 1997, Carlberg \\gE\\ 1997a). Reliably measured masses are now available for these clusters from gravitational lensing, temperatures, and velocity dispersions, not previously available in the above studies. Strong Sunyaev-Zel'dovich decrements have also been observed for these clusters, further suggesting that these are massive clusters with large amount of gas. The three clusters  have the highest masses (from weak lensing observations), the highest velocity dispersions ($\\sigma_{r} \\gtrsim $1200\\ km $s^{-1}$), and the highest temperatures (T$\\gtrsim $8 kev) in the $z >$ 0.5 EMSS survey (\\S 2). Therefore, they provide a strong constraint  on cosmology. We discuss the cluster data in \\S 2 and the cosmological implications in \\S 3. A Hubble constant of $\\rm H_{0}=100\\ h\\ km\\ s^{-1} Mpc^{-1}$  is used. ",
        "conclusions": ""
    },
    "9803/astro-ph9803195_arXiv.txt": {
        "abstract": "A measurement of the distance to Virgo Cluster by a direct method along with a realistic error analysis is important for a reliable determination of the value of Hubble Constant. Cepheid variables in the face-on spiral M100 in the Virgo Cluster were observed with the Hubble Space Telescope in 1994 under the HST Key Project on the Extragalactic Distance Scale. This work is a reanalysis of the HST data following our study of the Galactic Cepheids (in an accompanying communication). The periods of the Cepheids are determined using two independent methods and the reasons for varying estimates are analyzed. The log(period) vs $V$-magnitude relation is re-calibrated using LMC data as well as available HST observations for  three  galaxies  and the slope is found to be $-3.45 \\pm 0.15$. A prescription to compute  correction for the flux-limited incompleteness of the sample is given and a correction of 0 to 0.28 magnitude in $V$-magnitude for Cepheids in the period range of 35 to 45 days is applied. The extinction correction is carried out using {\\em period vs mean $\\vi $ color} and {\\em $V$-amplitude vs $\\vi $ color at the brightest phase} relations. The distance to M100 is estimated to be $20.4 \\pm 1.7$ (random) $\\pm 2.4$ (systematic) Mpc. ",
        "introduction": "\\label{sec:intro}  A natural scale length for the Universe is provided by the Hubble Constant ($H_0$) and undoubtedly a  determination of its value is one of the fundamental problems of cosmology.  Over the years, there has been much debate about the value of $H_0$ and the present estimates range from less than 50 \\ksm\\ to over 80 \\ksm. Most probably the major reason for the discrepancy is the conventional distance ladder involving multiple steps. Its main drawback is that analysis of the systematic errors becomes difficult when the calibrating local sample and the observed sample at the next step of the ladder are not identical. Consequently, it is believed that an accurate measurement of the distance to a galaxy cluster which is $\\sim $ 20--30 Mpc away, without involving intermediate steps, will lead to a reliable direct estimate of the value of $H_0$, provided the recession velocity of the cluster is independently known. The Virgo Cluster, which is our nearest cluster of galaxies, is fairly rich in terms of galaxy population, and an average of the distances to the individual galaxies by different methods would provide a good estimate to its mean distance.  One of the key projects of the Hubble Space Telescope (HST) was devoted to the calibration of the extragalactic distance scale for a determination of $H_0$ with reasonable accuracy. An examination of the systematic errors in the Cepheid \\plr\\ and measurement of the distance to the Virgo Cluster through Cepheid observation were among the primary aims of this key project. The nearly face-on spiral M100 in the Virgo Cluster was observed on 12 epochs over a span of $\\sim$ 57 days in 1994 with the HST using the filters F555W and F814W, which are almost equivalent to the Johnson V and Cousins I bands respectively (\\cite{freedman:94}). The advantage of choosing this particular galaxy is that being nearly face-on, the errors due to extinction and reddening are expected to be minimal, and further, it is considered to be very similar to the Milky Way in terms of age, chemical composition etc. However, its position relative to the center of the Virgo Cluster is not known accurately, and that introduces some uncertainty in the Virgo distance derived from direct distance estimation to M100. Ferrarese \\ea (1996) reported observations of 70 Cepheids in M100, and obtained a distance of $ 16.1 \\pm 1.3$ Mpc. The value of the Hubble Constant was calculated to be $88 \\pm 24$ \\ksm. On the other hand, Sandage and collaborators re-calibrated the distance to a few galaxies, where supernovae of type Ia were detected earlier, by observing the Cepheids in those galaxies with the HST. They obtained a mean absolute $B$ magnitude at peak of $-19.6$ for normal SN Ia and consequently, a value of $ 52 \\pm 9$ \\ksm\\ for the Hubble Constant (\\cite{sandage:94}; \\cite{saha:94}). However, more recent publications indicate a better reconciliation in the value of $H_0$. Freedman \\ea (1998) summarize a value of $73 \\pm 6$ (statistical) $\\pm 8$ (systematic) \\ksm, as compared to $55 \\pm 3$ (internal) \\ksm\\ quoted by Sandage's group (\\cite{saha:96}).  The present work is a re-analysis of the HST data on M100 Cepheids, based on a general calibration of Galactic Cepheids, presented in a companion paper (which we refer henceforth as Paper I). The specific problems we address here are the following: \\bei \\item Period--Luminosity relation applicable to the Cepheids of period $\\ga $ 15 days generally observed in distant galaxies. \\item Importance of the incompleteness correction and quantification of the effect. \\item Uncertainty in the periods of the Cepheids in M100 due to the phase sampling techniques applied as well as the large error in $V$-magnitude, particularly at low flux levels. \\eni  The central idea behind distance measurement with Cepheids is the \\plr. However, the values of both the slope and the intercept of this relation continue to be subjects of lively debate. There appears to be a distinct difference in the value of the slope between Cepheids of low and high periods. While applying the \\plr\\ to distant galaxy samples, where only higher period Cepheids can be detected due to flux limitation, this distinction becomes even more crucial. We address this question on the basis of our study of Galactic Cepheids (Paper I) which demonstrates a clear division between two classes of Cepheids, one with periods $\\leq 15$ days, the other at higher periods. The zero-point of the \\plr\\ is another quantity which needs to be fixed unambiguously in order to obtain reliable estimates of distance. We use the recent calibration of the local Cepheids by the Hipparcos mission (\\cite{fc:97}), rather than the distance to the Large Magellanic Cloud which is normally treated as the calibrating point for the distance scale.  A crucial aspect of our new analysis is the correction for incompleteness of the Cepheid sample. Since the Cepheid \\plr\\ has an intrinsic scatter due to the finite width of the instability strip at a given period, the Cepheids are observed to be spread over a range of luminosities. All the observed Cepheids in M100 have $V$-magnitudes between 24 and 27 mag. At such faint flux levels it is very likely that for a fixed period, the fainter Cepheids would escape detection, and only the brighter ones will appear in the surveys. This selection bias would have a systematic effect on the period--$V$-magnitude slope, especially at low periods, reminiscent of the Malmquist bias discussed in the literature. In order to counter this effect, one has to take into account the undetected Cepheids, which can be done by adequately correcting the observed magnitudes to fainter levels. Obviously, the amount of correction depends on the scatter of the \\pl\\ diagram. We have devised a formalism to correct for this incompleteness effect which we demonstrate to be present to a large extent in the M100 sample.  We have tried to estimate the correction for extinction and reddening, which again, is based on our study of Galactic Cepheids (Paper I). However, in the absence of multi-wavelength observations, this treatment is rather limited, and is based on \\pca\\ relations of Cepheids. Also, for the same reason we could not isolate the extinction correction from the incompleteness correction which ideally we should have been able to.  This paper is organized as follows. In Section~\\ref{sec:per_det} we present our methods of determination of Cepheid periods and photometric parameters. The question of choosing the correct \\plr\\ is addressed in Section~\\ref{sec:plr}. In Section~\\ref{sec:incomp}, we devise a formalism for the incompleteness correction of a biased Cepheid sample, and the mathematical aspects of compensation for flux-limited bias are described in the Appendix. Section~\\ref{sec:extcor} deals with the reddening and extinction corrections and the essentials of the numerical methods. The results and major contributions to errors are discussed in Section~\\ref{sec:results} and some remarks on the conclusions are presented in Section~\\ref{sec:summary}. ",
        "conclusions": ""
    },
    "9803/astro-ph9803126_arXiv.txt": {
        "abstract": "We revisit the problem of the flat slope of the $Mg_2$ versus $<Fe>$ relationship found for nuclei of elliptical galaxies (Faber et al. 1992; Worthey et al. 1992; Carollo et al. 1993; Davies et al. 1993), indicating that the Mg/Fe ratio should increase with galactic luminosity and mass. We transform the abundance of Fe, as predicted by classic wind models and alternative models for the chemical evolution of elliptical galaxies, into the metallicity indices $Mg_2$ and $<Fe>$, by means of the more recent index calibrations and show that none of the current models for the chemical evolution of elliptical galaxies is able to reproduce exactly the observed slope of the $<Fe>$ versus $Mg_2$ relation, although the existing spread in the data makes this comparison quite difficult. In other words, we can not clearly discriminate between models predicting a decrease (classic wind model) or an increase of such a ratio with galactic mass. The reason for this resides in the fact that the available observations show a large spread due mostly to the errors in the derivation of the $<Fe>$ index. In our opinion this fact prevents us from drawing any firm conclusion on the behaviour of Mg and Fe in these galaxies. Moreover, as already shown by other authors, one should be careful in deriving trends in the real abundances just from the metallicity indices, since these latter depend also on other physical parameters than the metallicity. This is an important point since abundance ratios have been proven to represent strong constraints for galaxy formation mechanisms. ",
        "introduction": "Elliptical galaxies do not show the presence of HII regions and it is not possible to resolve single stars in them in order to measure photospheric abundances. Therefore, most of the information on these objects is obtained from their integrated properties: abundances are derived either through colors or integrated spectra and in both cases the derived information is a complicated measure of metallicity and age (the well known age-metallicity degeneracy). The most common metallicity indicators are $Mg_{2}$ and $<Fe>$, as originally defined in Faber et al. (1977; 1985). Population synthesis techniques are adopted to analyze the integrated properties of ellipticals and to derive an estimate of their real abundances. Unfortunately, they contain several uncertainties residing either in incomplete knowledge of stellar evolution or in deficiencies in stellar libraries, as discussed in Charlot et al. (1996). In recent years more and more population synthesis models (Bruzual and Charlot, 1993; Buzzoni et al., 1992; Bressan et al. , 1994; Gibson, 1997; Bressan et al., 1996; Gibson and Matteucci, 1997; Tantalo et al., 1998) have appeared but the basic uncertainties still remain. In this paper we want to focus our attention about the comparison between theoretical model results and metallicity indicators. In this framework we will analyze the relationship between $<Fe>$ and $Mg_2$ and its implications for the mechanism of galaxy formation.   Several authors  (Faber et al. 1992; Worthey et al. 1992; Carollo et al. 1993; Davies et al. 1993; Carollo and Danziger 1994), from comparison of the observed indices with synthetic indices, concluded that the average $[<Mg/Fe>]_{*}$ in giant ellipticals must be larger than the solar value. This result was also confirmed by the analysis of Weiss et al. (1995) who made use, for the first time, of stellar evolutionary tracks calculated under the assumption of non-solar ratios.The same authors found that the $<Fe>$ versus $Mg_{2}$ relation among nuclei of giant ellipticals is rather flat and flatter than within galaxies. From the flat behavior of $<Fe>$ vs. $Mg_2$ the same authors inferred that the abundance of Mg should increase faster than the abundance of Fe among nuclei of giant ellipticals. This conclusion is at variance with the predictions of supernova-driven wind models of ellipticals (Arimoto and Yoshii, 1987; Matteucci and Tornamb\\`e 1987). In fact, Matteucci and Tornamb\\`e (1987) showed that, in the framework of the classic wind model for ellipticals, the [Mg/Fe] ratio is a decreasing function of the galactic mass and luminosity. The reason for this behavior is  clear: if Fe is mostly produced by the supernovae of type Ia, as it seems to be  the case in our Galaxy (Greggio and Renzini 1983a; Matteucci and Greggio 1986), whereas Mg is mostly originating from supernovae of type II, then the iron production is delayed relative to that of Mg and its abundance should be larger in more massive galaxies which develop a wind later than the less massive ones. All of this is valid under the assumption that after the onset of a galactic wind star formation should stop or should be negligible, which is a reasonable assumption for elliptical galaxies. Faber et al. (1992) proposed alternative scenarios to the classic supernova driven wind model, as originally proposed by Larson (1974). They suggested three different scenarios all based on the assumption that Mg is produced by type II supernovae and Fe is mostly produced by type Ia supernovae: i) a selective loss of metals, ii) a variable initial mass function (IMF) and iii) different timescales for star formation. These hypotheses were discussed by Matteucci (1994), who tested them in the context of chemical evolution models. In the hypothesis of the  different timescales for star formation Matteucci (1994) suggested that the more massive ellipticals might experience a much stronger and faster star formation than less massive ellipticals leading to a situation where the most massive objects are able to develop galactic winds before the less massive ones. She called this case ``inverse wind model''. On the other hand, in the classic wind model of Larson (1974) the efficiency of star formation was the same for all galaxies thus leading to the fact that the galactic wind in more massive systems occurs later than in less massive ones, due to their deeper potential well. In the models of Arimoto and Yoshii (1987) and Matteucci and Tornamb\\`e (1987) the efficiency of star formation was a decreasing function of galactic mass, based on the assumption that the timescale for star formation is proportional to the cloud-cloud collision timescale which, in turn, is proportional to the gas density. Therefore, since in this monolithic collapse picture the gas density decreases with the galactic mass, the galactic wind was even more delayed for the most massive systems. Matteucci (1994) proposed, as an alternative, a star formation efficiency increasing with the galactic mass and she justified this assumption by imagining that giant elliptical galaxies, instead of forming through a monolithic collapse of a gas cloud, form by merging of gaseous protoclouds. The merging process can, in fact, produce higher densities  for increasing galactic mass and/or higher cloud-cloud collision velocities resulting in a faster star formation process. In such a model the galactic wind occurs earlier in massive than in smaller ellipticals thus producing the expected trend of an increasing [Mg/Fe] as a function of galactic mass. Matteucci (1994) also showed that a variable IMF with the slope decreasing with increasing galactic mass and luminosity can produce the same effect without an inverse wind situation. The reason for that resides in the fact that a flatter IMF slope favors massive stars relative to low and intermediate masses, thus favoring Mg production over Fe production. However, Matteucci (1994) could not translate the predicted abundances into $Mg_2$ and $<Fe>$ since there were no available calibrations for [Fe/H] versus $<Fe>$ but only calibrations for $[Fe/H]$ vs. $Mg_2$. Therefore she did not compare the predicted abundances with observations.  Recently, calibrations for the iron index have become available (Borges et al., 1995; Tantalo et al., 1998) and therefore in this paper we revisit the whole problem of inferring trends on the real abundances by metallicity indices and we discuss the influence of the calibration relationships, which allow us to pass from indices to abundances, and we show that the inferred trend of Mg/Fe with galactic mass is not so clear when interpreted in terms of real abundances, thus warning us from drawing any firm conclusion on galaxy formation processes just on the basis of the observed behavior of $<Fe>$ versus $Mg_{2}$. indices. The reason for that resides partly in the large spread present in the observational data and partly in the fact that metallicity indices depend not only on the abundances of single elements but also on the ages and on the metallicity distribution (Tantalo et al. 1998) of the different stellar populations present in elliptical galaxies. \\par In Section 2 we will discuss the chemical evolution model, in Section 3 we will define the average abundances of a composite stellar population, in Section 4 we will describe the model results and transform the predicted abundances into indices by means of the most recent metallicity calibrations. Finally in Section 5 some conclusions will be drawn.  \\noindent ",
        "conclusions": "In this paper we have discussed the relation between metallicity indices, such as $Mg_2$ and $<Fe>$, and total mass in nuclei of ellipticals and its implications in terms of models of formation and evolution of elliptical galaxies. \\par In order to do that we have transformed the average abundance of Fe in the composite stellar population of the galaxy, as predicted by different models of chemical evolution, into $Mg_2$ and $<Fe>$ indices by means of the available calibrations.\\par  We have shown the results of classic wind models for ellipticals, such as those discussed by Arimoto and Yoshii (1987) and Matteucci and Tornamb\\`e (1987), as well as the results of models with variable IMF from galaxy to galaxy and with galactic winds occurring first in the more massive systems, implying that these systems are older than the less massive ones. We have found that it is not possible to establish clearly which kind of model should be preferred, first of all because of the large spread present in the data.  Moreover, little difference is found in the predicted indices of models which predict a $[<Mg/Fe>]_{*}$ either increasing or decreasing with galactic mass, although the data seem to suggest an increase of this ratio with galactic mass larger than predicted by any of the models.  On this basis, the classic wind model can not be considered worse than the other models. Actually, the classic wind model with a flat constant IMF seems to be the only one which can reproduce the whole range of the observed indices. However, if we isolate the data from Gonzalez (1993) and do not consider the others, then in order to reproduce the flat slope of the $<Fe>$ versus $Mg_2$ relation, as given by the best-fit of the data, one should assume that Fe among the nuclei of ellipticals is almost constant whereas Mg increases from less massive to more massive nuclei. This is not achieved by any of the models presented here since it would require quite ``ad hoc'' assumptions especially concerning the type Ia SNe.  From the numerical experiments performed in this paper we can say that a model which explains at the same time the mass-metallicity and the iron-magnesium relation requires an inverse wind situation, with a strong increase of the star formation efficiency with galactic mass (i.e. Model VI), rather than a variable IMF from galaxy to galaxy, and an IMF with a slope $x=0.8$. However, a model of this type is not able to reproduce the observed ranges of $<Fe>$ and $Mg_2$. We have also calculated models where amount and concentration of dark matter increases, compatibly with the formulation of the potential energy of the gas, with decreasing galactic luminous mass (see Persic et al. 1996), with the net result of obtaining an ``inverse wind'' situation. The results are very similar to those of Model III. Therefore, to obtain a better agreement with observations one should invoke also in this case an increase of the star formation efficiency with galactic mass. This would certainly flatten the $<Fe>$ vs. $Mg_2$ relation but it would further shrink the ranges of the predicted indices. In fact, both an increasing star formation efficiency and a decreasing importance of dark matter with increasing luminous galactic mass can be viable solutions to achieve the situation of more massive ellipticals being older than less massive ones.    \\par In conclusion, it is quite important to establish the value of [Mg/Fe] from the observational point of view since abundance ratios, such as [Mg/Fe], represent an important diagnostic to infer ages in galaxies, due to the different timescales for the Mg and Fe production. Generally, a high [Mg/Fe] is interpreted as a young age and the upper limit for the age is given by the time at which the chemical enrichment from type Ia SNe starts to become important. This timescale depends not only on the assumed progenitors of type Ia SNe but also on the star formation history of the galaxy considered (see Matteucci 1997) and for giant elliptical galaxies this timescale is of the order of $t_{SNeIa}\\sim 3-4 \\cdot 10^{8}$ years and in any case it can not be larger than 1 Gyr also for smaller systems. This is at variance with what stated by Kodama and Arimoto (1997) who claim, on the basis of results concerning our Galaxy (Yoshii et al. 1996), that this timescale is of the order of 1.5-2.5 Gyr. This is indeed true for our Galaxy where the star formation history has been quite different than in ellipticals and it had been already pointed out in Greggio and Renzini (1983b) and in Matteucci and Greggio (1986). This is a quite important point, both for the galactic chemical enrichment and for the predictions about SN rates at high redshift.   Therefore, an enhanced [Mg/Fe] indicates that the process of galaxy formation must have been very fast thus favoring a monolithic collapse scenario rather than a merging scenario. In this framework, a [Mg/Fe] ratio higher in more massive ellipticals than in less massive ones could be interpreted as due to their faster formation and evolution (see Matteucci 1994; Bressan et al. 1996). \\par An independent way of estimating the ages of ellipticals, where for ages we intend the time elapsed from the last star formation event, is to study the $H_{\\beta}$ index. This index is, in fact, related to the age of the dominant stellar population, since it gives a measure of the turn-off color and metallicity. It can therefore be used to solve the age-metallicity degeneracy. Bressan et al. (1996), by analyzing the $H_{\\beta}$ and other physical parameters in the sample of ellipticals observed by Gonzalez (1993), concluded that massive galaxies should have stopped forming stars before less massive ones, in agreement with the results of the inverse wind model discussed here. Finally, we would like to point out that both models with a Salpeter IMF and a variable IMF have a potential problem in reproducing high [$\\alpha$/Fe] ratios in the intracluster medium (ICM), as shown by their low average $[<\\alpha/Fe>]_*$ ratios (see Tables 7-12). Therefore, in agreement with MG95 and Gibson and Matteucci (1997) we conclude that a flat IMF is required to explain the high [$\\alpha$/Fe] ratios, as found by ASCA observations (Mushotzsky 1994)."
    },
    "9803/astro-ph9803256_arXiv.txt": {
        "abstract": "We introduce and study the distribution of an estimator for the normalized bispectrum of the Cosmic Microwave Background (CMB) anisotropy. We use it to construct a goodness of fit statistic to test the coadded 53 and 90 GHz {\\it COBE}-DMR 4 year maps for non-Gaussianity. Our results indicate that Gaussianity is ruled out at the confidence level in excess of  98$\\%$. This value is a lower bound, given all the investigated systematics. The dominant non-Gaussian contribution  is found near the multipole of order $\\ell=16$. Our attempts to explain this effect as caused by the diffuse foreground emission from the Galaxy have failed. We conclude that unless there  exists a microwave foreground emission which spatially correlates  neither with the DIRBE nor Haslam maps, the cosmological CMB anisotropy is genuinely non-Gaussian. ",
        "introduction": "We shall consider fluctuations in the CMB as a random field on the sphere, $\\frac{\\Delta T}{T}({\\bf n})$. One can expand such a field in terms of Spherical Harmonic functions: \\begin{eqnarray} \\frac{\\Delta T}{T}({\\bf n})=\\sum_{\\ell m}a_{\\ell m}Y_{\\ell m}({\\bf n}) \\label{almdef} \\end{eqnarray} For a statistically isotropic field one has \\begin{eqnarray} \\langle a_{\\ell_1 m_1}a^*_{\\ell_2 m_2}\\rangle=C_{\\ell_1} \\delta_{\\ell_1 \\ell_2} \\delta_{m_1 m_2} \\label{defiso} \\end{eqnarray} We can also define the two-point function in terms of $\\frac{\\Delta T}{T}({\\bf n})$. Isotropy implies that the correlation matrix can only depend on the angle between the two points considered. This is encoded in the 2-point correlation $C^{(2)}(\\theta)$. From (\\ref{almdef}) and (\\ref{defiso}) we find \\begin{equation}\\label{c2cl} C^{(2)}(\\theta)={\\sum _\\ell}{2\\ell+1\\over 4\\pi}C_\\ell P_\\ell(\\cos\\theta) \\end{equation} Hence the $C_\\ell$ may be regarded as a Legendre transform of the 2-point correlation function.  It is a standard lore that, barring some mathematical obstructions, one can reconstruct the probability distribution function of any random field from its moments. Isotropy imposes ``selection rules'' on these moments. For instance, the 3-point moment is given by \\begin{eqnarray} \\langle a_{\\ell_1 m_1}a_{\\ell_2 m_2} a_{\\ell_3 m_3}\\rangle= \\left ( \\begin{array}{ccc} \\ell_1 & \\ell_2 & \\ell_3 \\\\ m_1 & m_2 & m_3 \\end{array} \\right ) C_{\\ell_1\\ell_2\\ell_3} \\label{defnpoint} \\end{eqnarray} where the $(\\ldots)$ is the Wigner $3J$ symbol. The coefficients $C_{\\ell_1\\ell_2\\ell_3}$ are usually called the bispectrum. If we assume that there are no correlations between different $\\ell$ multipoles then the only non-zero component of the bispectrum is $C_{\\ell\\ell\\ell}=B_\\ell$. The collapsed 3-point correlation function $C^{(3)}(\\theta)$ (the average of a temperature squared at one point, and a temperature at another point, separated by  an angle $\\theta$) is now \\begin{eqnarray} C^{(3)}(\\theta)={\\sum _\\ell}{\\left(2\\ell+1\\over 4\\pi\\right)}^{3/2} \\left ( \\begin{array}{ccc} \\ell & \\ell & \\ell \\\\ 0 & 0 & 0 \\end{array} \\right )B_\\ell P_\\ell(\\cos\\theta) \\end{eqnarray} in analogy with (\\ref{c2cl}). Hence the $B_\\ell$ is related to the Legendre transform of $C^{(3)}$. The angular power spectrum $C_\\ell$ is often considered a more powerful tool than the correlation function $C^{(2)}(\\theta)$ for discriminating between theories, and one might argue the same way with  regard to the reduced bispectrum $B_{\\ell}$ and the 3-point function $C^3(\\theta)$.  The importance of higher order statistics for characterizing large scale structure has been stressed before (\\cite{lss}). The non-linear evolution of primordial Gaussian fluctuations has been analysed in  detail (\\cite{lss,bouch92}) and the skewness arising in such models has been shown to be consistent with current observations (\\cite{bouch93,gaz94}). \\cite{luo94} discussed the statistical properties and detectability of the bispectrum for a variety of non-Gaussian signals. \\cite{kog96a} measured the pseudocollapsed and equilateral three point function of the DMR four year data and found them to be consistent with Gaussianity. The analysis performed here should be considered complementary to that of \\cite{kog96a}: non-Gaussian signals which may be obscured in real space can become evident in $\\ell$ space.  In this letter we shall use a  general formalism for generating estimators of higher order moments on a sphere (\\cite{fergormag}). In this formalism one considers all possible tensor products of $\\Delta T_\\ell$ (each multipole component of the field) and from these one extracts the singlet (invariant) term. In the case of  bispectrum one has \\begin{eqnarray} {\\hat B}_\\ell&=&\\alpha_\\ell\\sum_{m_1m_2m_3}\\left ( \\begin{array}{ccc} \\ell & \\ell & \\ell \\\\ m_1 & m_2 & m_3 \\end{array} \\right ) a_{\\ell m_1}a_{\\ell m_2} a_{\\ell m_3} \\nonumber \\\\ \\alpha_\\ell&=&\\frac{1}{(2\\ell+1)^{\\frac{3}{2}}}\\left ( \\begin{array}{ccc} \\ell & \\ell & \\ell \\\\ 0 & 0 & 0 \\end{array} \\right )^{-1} \\label{bispec} \\end{eqnarray} Note that only even values of $\\ell$ lead to nonzero values of the ${\\hat B}_\\ell$ due to the symmetries of the Wigner 3-J coefficients. In practice it is essential to factor out the power spectrum from our statistic. We also wish to define statistics which are invariant under parity transformations, and not just rotations. Therefore we define $I^3_\\ell$ to be \\begin{eqnarray} I^3_\\ell &=&\\left| { {\\hat B}_\\ell\\over ({\\hat C}_\\ell)^{3/2}} \\right| \\label{defI} \\end{eqnarray} where ${\\hat C}_\\ell=\\frac{1}{2\\ell+1}\\sum_m|a_{\\ell m}|^2$. Our statistics are dimensionless and are normalized so that a cylindrically symmetric multipole has $I^3_\\ell=1$.   The $\\ell=2$ case was discussed  and given a physical interpretation in \\cite{mag1}. The quadrupole has 5 degrees of freedom. Of these only 2 are rotationally invariant. One is the quadrupole intensity $C_2$, and tells us how much power there is in the quadrupole. The other is essentially $I^3_2$ and tells us how this power is distributed among the different $a_{2 m}$ but only as far as there is a rotationally invariant meaning to the concept. For instance if $I^3_2=1$ then there is a frame in which all the power is concentrated in the $m=0$ mode. Such a quadrupole is cylindrically symmetric, but of course the symmetry axis orientation is uniformly distributed, to comply with statistical isotropy. If $I^3_2=0$ then on the contrary cylindrical symmetry is maximally broken. The probability distribution function of $I^3_2$  is uniform in Gaussian theories (\\cite{mag1}). ",
        "conclusions": "The result that we have obtained raises a number of questions which we shall attempt to answer. From Fig.~\\ref{fig1} it is clear that $I^3_{16}$ is far in the tail of the Gaussian ensemble and it dominates the statistic.  One would like to understand the importance of both cosmic variance and noise to this measurement. We would also like to assess the extent to which a galactic foreground contaminant could be responsible for this result.  In order to answer the first question we look for Bayesian estimates for the $I^3_\\ell$ as they are {\\it for our sky}. To do this we first estimate what the temperature fluctuations $T_i=\\frac{\\Delta T}{T}({\\bf n}_i)$ in each pixel $i$ in our dataset are likely to be, given DMR observations $O_i$, and noises $\\sigma_i^2$. We construct the posterior $P(T_i|O_i)$ assuming uniform priors in $T_i$, and also that a priori no correlations exist between the $T_i$. The latter assumption is often used in image restoration algorithms, such as maximum entropy methods. We then produce an ensemble of skies with the distribution $P(T_i|O_i)$. From it we infer $P(I^3_\\ell|O_i)$, the distributions for what the $I^3_\\ell$ for our sky are likely to be given DMR observations and noise. This procedure will allow us to assess the importance of noise in each of our measurements. However note that this analysis is totally decoupled from the result in the previous section where all we need to know are the observed $I^3_\\ell$, not their estimates for our sky.   In Fig.~\\ref{fig2} we plot in dotted lines $P(I^3_\\ell|O_i)$ for our data set. We also plot in solid lines the cosmic variance distribution of $I^3_\\ell$ in skies with the same galactic cut. The vertical line is the observed invariant $I^3_\\ell(O_i)$. As expected we see that, as $\\ell$ gets larger, the spread in $P(I^3_\\ell|O_i)$ due to noise becomes more important,  at $\\ell=18$ dominating the distribution function. On the other hand we clearly have succeeded in making measurements for $\\ell=4,6,8,12,14,16$. For them $P(I^3_\\ell|O_i)$ are peaked and clearly different from the cosmic variance distribution. The fact that $P(I^3_{16}|O_i)$ does not peak at $I^3_{16}(O_i)$ is merely a failure of the prior. The measurement of $I^3_{16}(O_i)$ is therefore a signal and is not dominated by noise. We have further checked that the signal to noise in power at $\\ell=16$ is of order 1.   Next we wish to know if galactic emissions could be blamed for this result. We can proceed in three ways. Firstly we may use instead the DMR cosmic emission maps, where a linear combination of the various DMR channels is used to separate out the foreground Galactic contamination. In these maps the noise level is considerably higher. Plotting the counterpart of Fig.~\\ref{fig2} for this case we find that the distributions of the actual $I^3_\\ell$ for our sky, given noise induced errors, are very similar to their cosmic variance distributions. The measurement is therefore dominated by noise and inconclusive. We find $X^2_{COBE}=.4$, consistent with Gaussianity, but this is a mere check of the Gaussianity of noise. Hence this approach towards foregrounds turns into a dead end, but serves to show how large angle Gaussian tests is a field constrained by noise, not cosmic variance.   As an alternative approach we may subject galactic templates to the same analysis. At the observing frequencies the obvious contaminant should be foreground dust emission. The DIRBE  maps (\\cite{boggess92}) supply us with a useful template on which we can measure the $I^3_\\ell$s. We have done this for two of the lowest frequency maps, the $100$ $\\mu$m  and the $240$ $\\mu$m maps. The estimate is performed in exactly the same way as for the DMR data (i.e. using the extended Galaxy cut). We performed a similar exercise with the Haslam 408Mhz (\\cite{haslam}) map. We display their values in Fig.~\\ref{fig3}. As expected the two maps have consistent values for the $I^3_\\ell$. However they do not have a non-Gaussian value at $\\ell=16$. Indeed for all $\\ell$ the $I^3_\\ell$ are within Gaussian cosmic variance error bars. This is not surprising. DIRBE maps exhibit structures on very small scales. These should average into a Gaussian field when subject to a $7^\\circ$ beam.  As a third alternative we may use foreground corrected maps.  In these one corrects the coadded 53 and 90 Ghz maps for the DIRBE correlated emission. We have considered corrected maps in ecliptic and galactic frames, and also another map made in the ecliptic frame but with the DIRBE correction forced to have the same coupling as determined in the galactic frame. As shown in Fig.~\\ref{fig3}, in all of these the non-Gaussian signal at $\\ell=16$ is enhanced, although we observe large variations in $I^3_\\ell$ at $\\ell=4-8$ (a phenomenon noticed before  when estimating $C_\\ell$-s). In fact the corrected maps exclude Gaussianity at the confidence level of 99.5\\%.  It would be interesting to relate our result to the curious dip in power at $\\ell\\approx 16$ provided by the maximum likelihood estimates in \\cite{gorski97}. These show that, {\\it assuming a Gaussian signal}, the power in signal and noise is unusually low at $\\ell\\approx 16$. One wonders how this would be affected if non-Gaussian degrees of freedom were allowed into the estimation (\\cite{fergormag}).  We have also subjected our work to a variety of numerical tests. Arbitrary rotations of the coordinate system affects results to less than a part in $10^5$. More importantly, comparing data pixelized in the ecliptic and galactic frames, we found that our results were very robust, indeed  more so than the power spectrum estimation (see the bottom pannel of Fig.~\\ref{fig3}). We also tried different galactic cuts, and found that although the non-Gaussian signal gets transferred to other $\\ell$, one does not fully erase it until a cut of $\\pm 40^\\circ$ is applied. Finally we checked the effect of varying the offset in the cut map. We found that for any other prescription than the one used the effect is enhanced, often leading to rejecting Gaussianity at more than the 99.5\\% confidence level.   To conclude, we have not been able to attribute our result to a known contaminating source or a systematic. Indeed the confidence level quoted refers to the worst result obtained within the set of effects explored. Of course it is always possible that this non-Gaussian signal comes from some yet unmapped foreground, which cannot be separated from the CMB anisotropy signal in the {\\it COBE}-DMR data --- the poorly known free-free emission from the Galaxy comes to mind here.  If indeed our results  are due to a foreground contamination one should note the following two points. First, we would have demonstrated that DMR data is more contaminated by foregrounds than thought before. Second, Galactic emissions on the scales considered are often assumed to be Gaussian. In fact this assumption is used in subtraction algorithms based on the idea of optimal filtering. The discovery of a distinctly non-Gaussian galactic emission would in the very least require a rethinking of the foreground subtraction algorithms.  If, on the other hand, the CMB signal itself is demonstrably non-Gaussian, we would not  need to over-emphasise the epistemological implications of our findings."
    },
    "9803/astro-ph9803074_arXiv.txt": {
        "abstract": "\\noindent A phase transition in the nature of matter in the core of a neutron star, such as quark deconfinement or Bose condensation, can cause the spontaneous spin-up of a solitary millisecond pulsar. The spin-up epoch for our model lasts for $2\\times 10^7$ years or 1/50 of the spin-down time (Glendenning, Pei and Weber in Ref. \\cite{glen97:a}). The possibility exists also for future measurements on X-ray neutron stars with low-mass companions for mapping out the tell-tale ``backbending'' behavior of the moment of inertia. Properties of phase transitions in substances such as neutron star matter, which have more than one conserved charge, are reviewed. ",
        "introduction": "Neutron stars have a high enough interior density as to make phase transitions in the nature of nuclear matter a distinct possibility. Examples are hyperonization, negative Bose condensation (like $\\pi^-$ and $K^-$) and quark deconfinement. According to the QCD property of asymptotic freedom, the most plausible is the quark deconfinement transition. From   lattice QCD simulations, this phase transition is expected to occur in very hot ($T\\sim 200$ MeV) or cold but dense matter. In this work we will use the deconfinement transition as an example, but in principle, any transition that is accompanied by a sufficient softening of the \\eos and occurs at or near the limiting mass star, can produce a similar signal.  The paper is organized as follows. We discuss first the physical reason why a rapidly rotating pulsar, as it slows down over millions of years  because of angular momentum loss through the weak electromagnetic process of magnetic dipole radiation, will change in density due to weakening centrifugal forces  and possibly encounter, first at its center, and then in a slowly expanding region, the conditions for a phase transition. Conversely, an accreting star will be spun up from low to high frequency by accretion from a low-mass companion. This too will have a very long time-scale because accretion is regulated by the radiation pressure of the star's surface, heated by infalling matter.  After having discussed the reasons why we might see signals of phase changes, both in rapidly rotating stars that are spinning down because of angular momentum loss to radiation and stars that are spinning up due to the input of angular momentum by accretion, we discuss some aspects of phase transitions that are common to all first order transitions in neutron star matter, or more generally in isospin asymmetric matter. ",
        "conclusions": ""
    },
    "9803/astro-ph9803242_arXiv.txt": {
        "abstract": "We have developed 1D time-dependent numerical models of accretion discs, using an adaptive grid technique and an implicit numerical scheme, in which the disc size is allowed to vary with time. The code fully resolves the cooling and heating fronts propagating in the disc. We show that models in which the radius of the outer edge of the disc is fixed produce incorrect results, from which probably incorrect conclusions about the viscosity law have been inferred. In particular we show that outside-in outbursts are possible when a standard bimodal behaviour of the Shakura-Sunyaev viscosity parameter $\\alpha$ is used. We also discuss to what extent insufficient grid resolutions have limited the predictive power of previous models. We find that the global properties (magnitudes, etc. ...) of transient discs can be addressed by codes using a high, but reasonable, number of fixed grid points. However, the study of the detailed physical properties of the transition fronts generally requires resolutions which are out of reach of fixed grid codes. It appears that most time-dependent models of accretion discs published in the literature have been limited by resolution effects, improper outer boundary conditions, or both. ",
        "introduction": "The thermal-viscous accretion-disc instability model is more than 15 years old (see Cannizzo \\shortcite{can93b} for a historical overview). It is widely accepted that it provides the correct description of dwarf-nova outbursts and of (`soft') X-ray transient events. When, however, observations of these systems are compared with predictions of the model, the agreement is far from perfect (e.g. Lasota \\& Hameury \\shortcite{lh98} and references therein). It is sometimes also unclear what the predictions of the model are. One of the reasons for these uncertainties is the existence of various, different, versions of the model. From the very beginning these versions of the disc instability model differed in assumptions about viscosity and boundary conditions; they differed in the amount of matter accreted during the outburst, the shapes of light-curves, etc. (see Cannizzo \\shortcite{can93b}). At that time these differences seemed to be less important than the differences between the disc instability model and the competing, mass-transfer instability model \\cite{bp81}. The exponentially decaying tails of theoretical light curves predicted by the mass-transfer model were thought to contradict observations \\cite{can93b} and the outer disc radius behaviour during and after outbursts seemed to favour the disc instability model \\cite{io92}. The demise of the mass-transfer instability model was, however, caused by the lack of a physical mechanism which would trigger it.  With one model left it became important to establish just what its predictions are, and not merely whether it is better (or worse) than the competing model \\cite{pvw86}. The first systematic study of the disc instability model was presented by Cannizzo \\shortcite{can93a}, who analysed the importance of various terms in the disc evolution equations and the influence of the numerical grid resolution on the outburst properties. Ludwig, Meyer-Hofmeister \\& Ritter \\shortcite{lmr94} studied general properties of disc outbursts, such as the location of the instability that triggers them. Recently Ludwig \\& Meyer \\shortcite{lm98} analysed non-Keplerian effects which may arise during front propagation. The general conclusions of this group of studies were that non-Keplerian effects are negligible, that a few hundred grid points provide a sufficient resolution for the calculation results to be independent of the number of grid points, and, finally, that with the usual assumption (see Smak (1984b) of a jump in the value of the viscosity parameter $\\alpha$, the model produces only `inside-out' outbursts, i.e. outbursts starting in the inner disc regions. This last conclusion, if true, would entail changing the standard viscosity law (in which the $\\alpha$ parameter is constant in the hot and cool branch of the $\\Sigma$ - $T_{\\rm eff}$ curve) because outburts starting in the outer disc regions are clearly observed in classical dwarf-nova system SS Cyg \\cite{mau96}. This law already had to be modified when it was found \\cite{sma84b} that in order to get lightcurves similar to those observed in dwarf novae, $\\alpha$ in outburst had to be larger than $\\alpha$ in quiescence.  The absence of `outside-in' outbursts in these studies, however, is just the result of keeping the outer disc radius constant in the calculations (Section 4.1) \\cite{io92,sma84b}; from this point of view, there is no reason to modify the viscosity prescription. This does not in itself prove that changes in viscosity are correctly described by the bimodal behaviour of the $\\alpha$-parameter (see e.g. Gammie \\& Menou 1998), but the reasons given for preferring other versions (the exponential decay from outburst being the principal one) are not compelling, and these versions involve more fundamental changes in the disc physics (see Lasota \\& Hameury 1998 for discussion and references).  For example, Cannizzo, Chen \\& Livio \\shortcite{ccl95} use the formula $\\alpha = \\alpha_0 \\left(H/R \\right)^n$, but to make the model work they have to `switch off' convection. It is interesting, therefore, to recall that Faulkner, Lin \\& Papaloizou (1983) found dwarf nova outburst with $\\alpha$ constant, but their model was criticized \\cite{can93b} because they claimed that convection has only a minor influence on the energy transport in the disc. These are not formal problems because the constant $\\alpha$ models predict optically thin quiescent discs, whereas in bimodal $\\alpha$ models the quiescent disc is optically thick. There is observational evidence that dwarf nova discs in quiescence are optically thin (see Horne, 1993 and references therein).  Conclusions about the number of grid points required to get resolution-invariant results seem, on inspection, too optimistic, especially because fronts are not resolved, a point which is particularily worrying for the heating fronts.  The present situation of the disc instability model seems to be confused. Various versions are based on different assumptions about the physical processes in the disc and numerical codes suffer either from incorrect boundary conditions or from insufficient resolution or from both. Quite often, in the case of explicit codes the resolution is limited by the required computer time.  In this article we describe a numerical model of time-dependent accretion discs, using an adaptive grid technique and an implicit numerical scheme, in which the disc size is allowed to vary with time. This numerical scheme allows rapid calculations of disc outburst cycles at very high resolution. These properties allow an easy comparison with other versions of the model and a systematic study of its various assumptions.  In the near future we will use our code to model various properties of dwarf novae and X-ray transients. The model was alread used to model properties and outbursts of the dwarf-nova WZ Sge \\cite{lhh,hlh} and the rise to outburst of the X-ray transient GRO J1655-40 \\cite{hlmn97}.  In \\S 2 we discuss the time-dependent equations describing the disc radial structure and the implicit method used to solve them with a high spatial resolution. The vertical structure of the disc, and hence the heating and cooling terms that enter the time-dependent energy equation, are considered in \\S 3. In \\S 4 we present the results of our calculations and we discuss the importance of having sufficient numerical resolution and a correct boundary condition at the outer edge of the disc. ",
        "conclusions": "We have constructed a numerical code which can calculate, in a reasonable amount of computer time and at very high spatial resolutions, long cycles of accretion disc outbursts. This code works efficiently in the most general framework of the disc instability model and does not require special assumptions about viscosity or outer or inner radii. Its validity is of course limited by the way physical processes such as turbulent viscosity, convection, radiative transfer etc. are treated. Of course it is also a 1D code modeling a fundamentally 2D (or even 3D) situation.  Since the mass of the central object enters the disc equations only in $\\Omega_K$, we expect most of our results to be valid for BH disc models as well ({\\i.e. similar}, but at a slightly smaller radius), as long as irradiation and general relativistic effects can be neglected.  In future work we intend to include effects of irradiation (Dubus et al. 1998) and to apply the code to a systematic study of dwarf nova outbursts and X-ray transient events.   \\subsection*"
    },
    "9803/astro-ph9803132_arXiv.txt": {
        "abstract": "We have studied the poor southern cluster of galaxies S639. Based on new Str\\\"{o}mgren photometry of stars in the direction of the cluster we confirm that the galactic extinction affecting the cluster is large. We find the extinction in Johnson B to be $\\AB = 0.75\\pm 0.03$. We have obtained new photometry in Gunn $r$ for E and S0 galaxies in the cluster. If the Fundamental Plane is used for determination of the relative distance and the peculiar velocity of the cluster we find a distance, in velocity units, of $(5706\\pm 350)\\kms$, and a substantial peculiar velocity, $(839\\pm 350)\\kms$. However, the colors and the absorption line indices of the E and S0 galaxies indicate that the stellar populations in these galaxies are different from those in similar galaxies in the two rich clusters Coma and HydraI. This difference may severely affect the distance determination and the derived peculiar velocity. The data are consistent with a non-significant peculiar velocity for S639 and the galaxies in the cluster being on average 0.2 dex younger than similar galaxies in Coma and HydraI. The results for S639 caution that some large peculiar velocities may be spurious and caused by unusual stellar populations. ",
        "introduction": "The relation known as the Fundamental Plane (FP) may be used for determination of relative distances to E and S0 galaxies (e.g., Dressler et al.\\ 1987; J\\o rgensen, Franx \\& Kj\\ae rgaard 1996, hereafter JFK96; Baggley 1996; Hudson et al.\\ 1997). The FP relates the effective radius, $\\re$, the mean surface brightness within this radius, $\\Ie$ and the (central) velocity dispersion $\\sigma$, in a relation, which is linear in logarithmic space (Djorgovski \\& Davis 1987; Dressler et al.\\ 1987).  The FP has a low scatter ($15-20$\\% in $\\re$) and is therefore a valuable tool for studies of peculiar velocities of galaxies and clusters (e.g., Baggley 1996; Hudson et al.\\ 1997). The use of the FP for determination of distances and peculiar velocities relies on the assumption that the FP is universally valid. Several authors have investigated possible differences in the FP related to the cluster environment (e.g., Burstein, Faber \\& Dressler 1990; Lucey et al.\\ 1991; de Carvalho \\& Djorgovski 1992; JFK96; Baggley 1996). Only de Carvalho \\& Djorgovski find that the environment has significant effects on the FP. These authors find field galaxies to be brighter than cluster galaxies of similar effective radii and velocity dispersions. de Carvalho \\& Djorgovski also find field galaxies to be bluer and have weaker $\\Mgtwo$ line indices than cluster galaxies with similar velocity dispersions. This is in general agreement with studies that show that E and S0 galaxies in the outer parts of clusters have weaker $\\Mgtwo$ and $\\Fe$ indices than those in the central parts of clusters (Guzm\\'{a}n et al.\\ 1992; JFK96; J\\o rgensen 1997).  In this paper we study the poor cluster of galaxies S639, previously studied by JFK96. The cluster identification is from Abell, Corwin \\& Olowin (1989). The cluster has a radial velocity in the Cosmic Microwave Background (CMB) frame of $cz_{\\rm CMB}=6545\\kms$ and is located $\\approx 28$$\\degr$ from the direction to the large mass-concentration known as the ``Great Attractor'' (Faber \\& Burstein 1988). The velocity dispersion of the cluster is $456_{-74}^{+83}\\kms$ (JFK96). Its richness is 14 measured as the number of galaxies with magnitudes between $m_3$ and $m_3+2$ (Abell et al.\\ 1989). $m_3$ is the magnitude of the third ranked galaxy. S639 has a smaller velocity dispersion and is poorer than clusters like the Coma cluster and the HydraI cluster. Coma and HydraI have velocity dispersions of $1010_{-44}^{+51}\\kms$ and $608_{-39}^{+58}\\kms$, respectively (Zabludoff, Huchra \\& Geller 1990). The richnesses given by Abell et al.\\ (1989) is 106 for Coma and 39 for HydraI.  Using the FP for 10 E and S0 galaxies in S639 JFK96 found a large peculiar velocity of the cluster, $v_{\\rm pec}=(1295 \\pm 359)\\kms$ relative to the CMB frame. Further, JFK96 found that the galaxies in the cluster follow a $\\Mgtwo$-$\\sigma$ relation offset from the relation established for their full sample of 11 clusters. The galaxies in S639 had on average weaker $\\Mgtwo$ indices, see also J\\o rgensen (1997). JFK96 tried to correct the derived peculiar velocity of the cluster for the offset in the $\\Mgtwo$ indices by including a $\\Mgtwo$ term in the FP. The result was $v_{\\rm pec}=(879\\pm 392)\\kms$. However, the coefficient for the $\\Mgtwo$ term is not well determined, cf.\\ JFK96.  S639 is located at low galactic latitude, ($l$,$b$) = ($280\\degr$,$11\\degr$). Thus, the galactic extinction is large and uncertainties in the adopted value may severely affect the precision of the derived distance and peculiar velocity for the cluster.  The main issue discussed in this paper is whether the large peculiar velocity of S639 found by JFK96 is real or the result was caused either by incorrect correction for the (large) galactic extinction, by selection effects, or by unusual stellar populations. In order to reach conclusions about these issues we have obtained additional photometry of galaxies in S639, giving a sample of 21 E and S0 galaxies with available photometric and spectroscopic parameters. We have also obtained Str\\\"{o}mgren $uvby$-$\\beta$ photometry for stars in the direction of the cluster. This photometry is used to determine the galactic extinction affecting the cluster.  The sample selection for the E and S0 galaxies and the available data are briefly described in Sect.\\ 2. The determination of the galactic extinction is covered in Sect.\\ 3. In Sect.\\ 4 the FP is discussed and used for determining the distance to the cluster. The importance of the stellar populations is investigated in Sect.\\ 5. The conclusions are summarized in Sect.\\ 6.  The relations between the parameters for the galaxies established in this paper are determined by minimization of the sum of the absolute residuals perpendicular to the relations. This fitting technique has the advantage that it is rather insensitive to a few outliers, and that it treats the coordinates in a symmetric way. The uncertainties of the coefficients are derived by a bootstrap method. See also JFK96 for a discussion of this fitting technique. ",
        "conclusions": "The galactic extinction in the direction of the poor cluster of galaxies S639 has been determined from Str\\\"{o}mgren $uvby$-$\\beta$ photometry for stars in the direction of the cluster. Further, we have tested the consistency of the derived galactic extinction by using the $(B-r)$-$\\Mgtwo$ relation for E and S0 galaxies in the cluster. Our best estimate of the galactic extinction in the direction of the cluster is $\\AB = 0.75\\pm 0.03$.  The FP for S639 has been established based on a sample of 21 E and S0 galaxies. The coefficients for the FP for this cluster are not significantly different from those of the FP for the Coma and the HydraI clusters, and are also in agreement with previous results for other nearby clusters (e.g., JFK96). Under the assumption that the FP (coefficients and zero point) is universal we find a distance, in velocity units, to S639 of $(5706\\pm 350) \\kms$. This implies a peculiar velocity for the cluster of $(839\\pm 350)\\kms$.  The E and S0 galaxies in S639 have significantly smaller $\\Mgtwo$ indices, larger $\\HbG$ indices and are bluer than E and S0 galaxies of similar velocity dispersions in Coma and HydraI. The offset in the FP for S639 relative to Coma and HydraI may be due to a difference in the stellar populations, rather than a large peculiar velocity for S639. The data are consistent with a zero peculiar velocity and mean ages of the S639 galaxies 0.2 dex younger than the mean ages of similar galaxies in Coma and HydraI. Alternatively, the $\\Mgtwo$ indices and the $(B-r)$ colors are consistent with a metallicity difference of 0.1 dex, with galaxies in S639 having a lower metallicity than those in Coma and HydraI. In this case the peculiar velocity of S639 is $\\approx 490\\kms$. However, this interpretation is not consistent with the strong $\\HbG$ indices measured for the four galaxies, for which we have measurements of this index.  We conclude that the peculiar velocity of S639 is most likely overestimated if the FP is used as a distance determinator for this cluster. Even though many studies have shown that the FP (and the $\\Dn$-$\\sigma$ relation which is a projection of the FP) to a large degree is universally valid (e.g., Burstein, Faber \\& Dressler 1990; JFK96; Baggley 1996), our results for S639 caution that there may be exceptions (see also Gregg 1992). When distance determinations are attempted the best approach will be to obtain colors and line indices together with the other required data. This will give the possibility of identifying clusters (and galaxies), which deviate strongly from the mean relations between the various global parameters. These clusters can also be expected to deviate from the FP otherwise valid for the bulk of the E and S0 galaxies. One may attempt to include a $\\Mgtwo$ term in the FP in order to correct for the effects caused by differences in the stellar populations. However, the coefficient for such a term is not well determined, cf.\\ JFK96, and the peculiar velocities derived with this method may not be accurate enough for investigations of large-scale flows.  \\vspace{0.5cm} Acknowledgements: Lars Freyhammer is thanked for obtaining part of the observations used for this research. The Danish Board for Astronomical Research and the European Southern Observatory are acknowledged for assigning observing time for this project and for financial support. Support for this work was provided by NASA through grant number HF-01073.01.94A to IJ from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. HJS acknowledges financial support from the Carlsberg Foundation, Denmark."
    },
    "9803/astro-ph9803304_arXiv.txt": {
        "abstract": "\\B observed several galactic binary X--ray pulsars during the Science Verification Phase and in the first year of the regular program. The complex emission spectra of these sources are an ideal target for the \\B instrumentation, that can measure the emission spectra in an unprecedented broad energy band.\\\\ Using this capability of \\B a detailed observational work can be done on the galactic X--ray pulsars. In particular the 0.1--200 keV energy band allows the shape of the continuum emission to be tightly constrained. A better determination of the underlying continuum allows an easier detection of features superimposed onto it, both at low energy (Fe K and L, Ne lines) and at high energies (cyclotron features).\\\\ We report on the spectral properties of a sample of X--ray pulsars observed with \\B comparing the obtained results.\\\\ Some ideas of common properties are also discussed and compared with our present understanding of the emission mechanisms and processes. ",
        "introduction": "The instrumentation aboard \\B \\cite{sax,lecs,mecs,hp,pds} is particularly well suited to study the X--ray emission from X--ray pulsars. This class of sources is composed by binary systems in which a magnetized rotating neutron star accretes matter from a less evolved mass--donor star. The mass--donor may be a OB supergiant star as in the case of Vela X--1, a Be main--sequence or near--main--sequence star as in the case of transient recurrent pulsars like A0535+26, a low mass star as in the case of 4U1626--67. The type of mass donor star strongly affects the temporal behaviour on the medium (days) to long (years) time scales. The transient behaviour is almost completely restricted to the subclass of X--ray pulsars that have Oe or Be counterparts.  \\B has observed some persistent pulsars and one transient pulsar during the first year of its operative life. We report results from the observations of some of these sources, emphasizing the commonalities and the differences. In particular we discuss the observational evidence on cyclotron line feature, comparing the observed results, also in terms of possible correlations, with the expected ones on the basis of the available theoretical models. ",
        "conclusions": "The correlation showed in Figure 1 was ``qualitatively'' predicted by M\\'esz\\'aros and Nagel \\cite{mn} (see also \\cite{pulmod}). The model predicts a width of the cyclotron feature proportional to its energy and to the square root of the electron temperature of the atmosphere \\begin{equation} \\label{eq:fw} \\Delta \\omega_B \\simeq \\omega_B \\left( 8 \\times \\ln(2) \\times \\frac{\\rm kT_e}{\\rm m_ec^2} \\right)^{\\frac{1}{2}} |\\cos\\theta| \\label{eq13} \\end{equation} In this equation $\\Delta\\omega_B$ is the line width, $\\omega_B$ is the cyclotron line frequency, $T_e$ is the electron temperature and $\\theta$ is the viewing angle with respect to the magnetic field axis. A better insight on the properties of the cyclotron lines can be obtained with pulse--phase resolved spectroscopy, as equation \\ref{eq:fw} suggests that there may be a dependence on the viewing angle of the observed line width (see also \\cite{pphspe}).  However Araya and Harding (1996) \\cite{araya} caution that, in the limit of a single scattering, the line width is not related to the electron temperature.  This ambiguity in the interpretation of these observational data points out the need of a more detailed and quantitative model for the line properties and for the broad band continuum emission of X--ray pulsars"
    },
    "9803/astro-ph9803148_arXiv.txt": {
        "abstract": "The Goddard High Resolution Spectrograph (GHRS) of the {\\it Hubble Space Telescope (HST)} has been used to observe the boron 2500 \\AA\\ region of \\bd-13.  At a metallicity of [Fe/H]=$-3.00$ this is the most metal-poor star ever observed for B. Nearly 26 hours of exposure time resulted in a detection.  Spectrum synthesis using the latest Kurucz model atmospheres yields an LTE boron abundance of log \\eps(B)$= +0.01\\pm0.20$.  This value is consistent with the linear relation of slope $\\sim$1.0 between log \\eps(B$_{\\rm LTE}$) and [Fe/H] found for 10 halo and disk stars by Duncan \\etal\\ (1997).  Using the NLTE correction of Kiselman \\& Carlsson (1996), the NLTE boron abundance is log \\eps(B)$= +0.93\\pm0.20$. This is also consistent with the NLTE relation determined by Duncan \\etal\\ (1997) where the slope of log \\eps(B$_{\\rm NLTE}$) vs. [Fe/H] is $\\sim$0.7.  These data support a model in which most production of B and Be comes from the spallation of energetic C and O nuclei onto protons and He nuclei, probably in the vicinity of massive supernovae in star-forming regions, rather than the spallation of cosmic ray protons and alpha particles onto CNO nuclei in the general interstellar medium. ",
        "introduction": "The light elements lithium, beryllium, and boron are of great interest out of proportion to their very low abundances, having implications in Big Bang Nucleosynthesis and stellar structure, as well as in constraints on models of galactic chemical evolution.  The ``canonical'' theory of the origin of the elements Li, Be, and B was first presented by Reeves, Fowler, \\& Hoyle (1970) and further developed by Meneguzzi, Audouze, \\& Reeves (1971), and then Reeves, Audouze, Fowler, \\& Schramm (1973). In this model, most light element formation can be accounted for by galactic cosmic rays (GCR) impinging on the interstellar medium (ISM), assuming a constant flux of GCRs through the life of the Galaxy and making reasonable assumptions about CR confinement by the Galactic magnetic field.  Meneguzzi \\etal\\ (1971) also introduced the idea of a large (up to three orders of magnitude) increase in the low energy (5-40 MeV nucleon$^{-1}$) CR flux; since CRs in this energy range are mostly shielded from the Solar System by the solar wind, they are not detectable.  This additional CR flux increased the production of all light elements and matched the isotopic ratios and total abundances to the accuracy known at the time.  Reeves \\& Meyer (1978) added the additional constraint that models should match not only present-day abundances but their evolution throughout the life of the Galaxy.  Their conclusions were similar to MAR, except that they had to introduce infall of light-element-free matter into the Galactic disk to match the evolution with time.  In retrospect, it can be seen that the data they were fitting were sparse and not very precise.  With the launch of the the {\\it Hubble Space Telescope} ({\\it HST}) and the availability of uv-sensitive CCD detectors (B is usually observed at $\\lambda$2500 and Be at $\\lambda$3130), data are now much more numerous and accurate, and abundances can be traced from the epoch of formation of the Galactic halo until the present day.  In the past several years, the evolution of Li, Be, and B has been used as a test of different models of the chemical and dynamical evolution of the Galaxy.  For example, in the models of Vangioni-Flam \\etal\\ (1990), Ryan \\etal\\ (1992), and Prantzos \\etal\\ (PCV; 1993), light element production depends on the intensity and shape of the GCR spectrum, which in turn depends on the supernova (SN) and massive star formation rates.  It also depends on the rise of the (progenitor) CNO abundances and the decline of the gas mass fraction, which is affected by rates of infall of fresh (unprocessed) material and outflow, e.g. by SN heating.  Other things being equal, at early times when target CNO abundances were low, light element production would be much lower for a given CR flux than presently, when the ISM abundances are higher.  PCV found that even with these numerous adjustable parameters, no time-independent CR spectrum can reproduce the evolution of light element abundances.  By assuming a very particular form of time variation of the CR flux (greatly enhanced at early epochs), they were able to (barely) fit the evolution of the abundances.  The present investigation supports a different solution to the problem of the origin of the light elements.  Duncan \\etal\\ (1997) present B abundances in a large number of stars ranging in metallicity from $\\sim$solar to [Fe/H]~$\\sim -2.8$.  They find that (LTE) B follows metals in direct proportion from the earliest times (very metal-poor stars) to the present, with little if any change of slope between halo and disk metallicities.  A straightforward interpretation of this is that the rate of production of B and Be does {\\it not} depend on the CNO abundances in the ISM, and that the production site is associated with the production site for metals.  This would be true if the spallation process most important for light element production is not primarily protons and $\\alpha$ particles colliding with CNO nuclei in the ISM but rather C and O nuclei colliding with ambient protons and $\\alpha$ particles, probably in regions of massive star formation (cf. Vangioni-Flam \\etal\\ 1996 and Ramaty \\etal\\ 1997).  This paper focuses specifically on the B measurement in \\bd-13; it is consistent with (and was used to help determine) the relationships between LTE and NLTE B and [Fe/H] seen in Duncan \\etal\\ (1997). The data suggest that B production at the lowest metallicities occurs in the same way as today, and thus support the new description of galactic light element production.  It is possible that recent GRO satellite observations of gamma rays from the Orion Nebula (Bloemen \\etal\\ 1994; Bykov \\& Bloemen 1994) provide direct evidence of C and O spallation occurring today, even though not all instruments on GRO detected evidence of such spallation (Murphy \\etal\\ 1996).  The light element data alone, however, are the strongest evidence in support of a new model of their production. ",
        "conclusions": "Fig.~4 shows the LTE abundances as a function of [Fe/H] from stars analyzed in the larger investigation of Duncan \\etal\\ (1997), with the point for \\bd-13 from the present investigation emphasized.  It can be seen that there is an approximately linear relation between $\\log$ \\eps(B$_{\\rm LTE}$) and [Fe/H] over both disk and halo metallicities, and that \\bd-13\\ is consistent with this relationship.  A least-squares fit to all the data of Fig.~4 (allowing for errors in both coordinates) yields a slope of 0.96$\\pm$0.07 and a reduced chi square, $\\chi_{\\nu}^2$, of 0.71, indicating an excellent fit.  If NLTE abundances are used, as shown in Fig.~5, the slope is 0.70$\\pm$0.07 and $\\chi_{\\nu}^2 = 1.63$.  Although it is true that \\bd-13\\ is used to determine this line, it is important to note that this star is quite consistent with the trend defined by the other stars.  Inspection and $\\chi_{\\nu}^2$ tests confirm this.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{figs/fig4.eps}} \\caption{LTE B abundances from Duncan \\etal\\ (1997) with the program star highlighted.} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{figs/fig5.eps}} \\caption{NLTE B abundances from Duncan \\etal\\ (1997) with the program star highlighted.} \\end{figure}   \\subsection{Comparison to standard models and a new model for light element production}  The slope of close to 1 suggesting a primary process is {\\it not} expected from canonical models of CR spallation in the ISM, which predict a secondary process and thus a steeper relation.  In a secondary process the rate of light element production depends on the product of the abundance of target CNO nuclei and the CR flux, both of which vary with time.  If SNe are the source of the target nuclei and the ISM is well-mixed, the ISM metallicity is proportional to the integral (total) number of SN up to a given time.   If, as is commonly supposed, SNe also seed the acceleration mechanism which produces CRs, the CR flux is proportional to the SN rate.  The result is light element abundances which vary quadratically with the metallicity of the ISM, or a logarithmic slope of 2 (Prantzos \\etal\\ 1993). Figs.~4 and 5 show that such a slope is certainly not consistent with our data. Duncan \\etal\\ (1997) discuss these issues in greater detail.  As was discussed by Duncan, Lambert, and Lemke (1992), the data for the first three metal-poor stars observed for B already seemed to show a linear (in the log) relationship with [Fe/H], suggesting some primary production mechanism rather than the secondary mechanism in the ISM described above. This idea has been modelled in detail by Cass\\'e \\etal\\ (1995), Ramaty \\etal\\ (1995 and 1997), Lemoine \\etal\\ (1997), and Vangioni-Flam \\etal\\ (1996).  In the new scenario, B and Be are primarily produced by the spallation of C and O onto protons and $\\alpha$ particles.  Such a process could occur near massive star SNe, where the particle flux would be very non-solar in composition; depleted in H and He and especially enriched in O and C.  Vangioni-Flam \\etal\\ find that a composition matching either winds from massive (Wolf-Rayet; WR) stars in star-forming regions or massive star SNe produce a flux of O and C which, after further acceleration, can reproduce both the magnitude and slope of B production seen in Figs.~4 and 5 through collisions with protons and $\\alpha$ particles.  As Ramaty \\etal\\ point out, production of some additional $^{11}$B by the neutrino process (Woosley \\etal\\ 1990) is not ruled out, and may be favored on energetic grounds.  Nevertheless the bulk of the B and Be would be produced from the spallation process.  Although the NLTE correction to the B abundance of \\bd-13 is relatively large and tends to raise the B abundance above the line in Fig.~5, two other effects not included here would tend to move it closer to the curve.  One is the effect of the blending Co line discussed above, which could reduce the B abundance as much as $\\sim$0.15 dex.  The other is the fact that if the spallation producing light elements is caused by O (and to a lesser extent C), $\\log$ \\eps(B) should be plotted against [O/H] rather than [Fe/H].  As is well-known, very metal-poor stars are overabundant in O compared to Fe (moving points representing the most metal-poor stars to the right in a figure with O on the x-axis).  Duncan \\etal\\ (1997) make such a plot, and demonstrate that although the measurement errors in O are greater than those for Fe, when all the metal-poor stars are considered together a straight line of slope 1.10$\\pm$0.14 fits the LTE abundances, and one of slope 0.82$\\pm$0.10 the NLTE abundances. However, oxygen abundance measurements are also surrounded by greater uncertainty than are iron measurements, and systematic errors in the oxygen abundance which depend on metallicity will affect the derived slope."
    },
    "9803/astro-ph9803181_arXiv.txt": {
        "abstract": "Following an approach developed by Paczy\\'nski \\& Stanek, we derive a distance to the Large Magellanic Cloud (LMC) by comparing red clump stars from the {\\em Hipparcos}\\/ catalog with the red clump stars observed in two fields in the LMC that were selected from the ongoing photometric survey of the Magellanic Clouds to lie in low extinction regions. The use of red clump stars allows a single step determination of the distance modulus to the LMC, $\\mu_{0,LMC} = 18.065\\pm 0.031\\pm 0.09\\;$mag (statistical plus systematic error), and the corresponding distance, $R_{LMC}= 41.02\\pm 0.59\\pm 1.74\\;kpc$. This measurement is in excellent agreement with the recent determination by Udalski et al., also based on the red clump stars, but is $\\sim 0.4\\;$mag smaller than the generally accepted value of $\\mu_{0,LMC} = 18.50\\pm 0.15\\;$mag. We discuss possible reasons for this discrepancy and how it can be resolved. ",
        "introduction": "The generally accepted distance modulus to the Large Magellanic Cloud (LMC) is $\\mu_{0,LMC} \\approx 18.5 \\pm0.15\\;$mag (for recent discussion see Westerlund 1997, Madore \\& Freedman 1998).  However, there is a long standing $\\sim 0.3\\;$mag discrepancy between the ``long'' distance determined using Cepheids (e.g. Laney \\& Stobie 1994) and the ``short'' distance determined using RR Lyr stars (e.g. Walker 1992, Layden et al.~1996). A similar discrepancy is present in the distance to the LMC derived with the supernova SN1987A ($\\mu_{0,LMC}< 18.37\\;$mag, Gould \\& Uza 1998; $\\mu_{0,LMC} = 18.56\\;$mag, Panagia et al.~1997).  Recently Udalski et al.~(1998) used red clump stars observed in the LMC by the OGLE 2 project (Udalski et al.~1997) and obtained a value of $\\mu_{0,LMC} = 18.08\\pm 0.03 \\pm 0.12 \\;$mag (statistical plus systematic error).  This distance modulus is $\\sim 0.4\\;$mag smaller than the ``long'' distance modulus used, for example, by the {\\em HST}\\/ Extragalactic Distance Scale Key Project team (e.g.~Rawson et al.~1997 and references therein).  Because errors in the distance to the LMC can propagate into errors in such key quantities as distances, luminosities, masses, and sizes of extragalactic objects, it is important to check the result of Udalski et al.~(1998) using independent data, in order to investigate possible systematic errors.  Red clump stars are the metal rich equivalent of the better known horizontal branch stars, and theoretical models predict that their absolute luminosity only weakly depends on their age and chemical composition (Seidel, Demarque, \\& Weinberg 1987; Castellani, Chieffi, \\& Straniero 1992; Jimenez, Flynn, \\& Kotoneva 1998).  Indeed, the absolute magnitude-color diagram from {\\em Hipparcos}\\/ data (Perryman et al.~1997, their Figure~3) clearly shows a compact red clump -- the variance in the $I$-band magnitude is only $\\sim 0.15\\;$mag (Stanek \\& Garnavich 1998; Udalski et al.~1998).  Despite their large number and the theoretical understanding of their evolution, red clump stars have seldom been used as distance indicators.  However, Stanek (1995) and Stanek et al.~(1994, 1997) used these stars to map the Galactic bar.  Paczy\\'nski \\& Stanek (1998) used the red clump stars observed by the OGLE project (Udalski et al.~1993) to obtain the distance to the Galactic center.  Stanek \\& Garnavich (1998) used red clump stars observed by the {\\em HST}\\/ in M31 to obtain a one-step distance to this galaxy.  In this paper we follow the approach of Paczy\\'nski \\& Stanek (1998) and present an estimate of the distance to the LMC based on the comparison between the red clump giants observed locally by the {\\em Hipparcos}\\/ (Perryman et al.~1997) satellite and those observed in the LMC by the $UBVI$ digital photometric survey of the Magellanic Clouds (Zaritsky et al.~1997).  In Section 2 we describe the data used in this paper and select low extinction regions for further analysis.  In Section 3 we analyze the red clump distribution in the LMC and derive the distance to this galaxy. In Section 4 we discuss the possible reasons for the discrepancy with the Cepheid distance to the LMC and how it can be resolved. ",
        "conclusions": "As with all distance-ladder techniques, our analysis includes the assumption that the calibrating and target objects being compared are intrinsically similar. In our red clump analysis, this assumption is manifested by the assertion that the $I$-band brightness of red clump stars is independent of the age, chemical composition, and mass differences that may exist between the red clump stars near the Sun and those in the LMC.  Indeed, the LMC red clump is systematically bluer than the local one, indicating the somewhat different properties of these stars. However, Paczy\\'nski \\& Stanek (1998), Stanek \\& Garnavich (1998) and Udalski et al.~(1998) found that the $I$-band peak magnitude of the red clump depends very weakly on their $(V-I)_0$ color in the range $0.7<(V-I)_0<1.4$, and therefore is independent of the metallicity (Jimenez et al.~1998).  This is confirmed in this paper as well, by comparing the peak brightness of the red clump for two color ranges $0.55<(V-I)_0 <0.8$ and $0.8<(V-I)_0 <1.25$ (see the previous Section). The fact that the observed red clump distributions are so narrow ($\\sigma_{RC}\\approx 0.15\\;$mag) indicates that the age dependence of the red clump $I$-band peak luminosity is also small ($\\lesssim 0.1\\;$mag).  Otherwise, in a system with a complex star formation history, such as the LMC (Holtzman et al.~1997; Geha et al.~1998), the resulting red clump should have considerable width. Stanek \\& Garnavich (1998) compared three different lines-of-sight that probe a large range of M31 galactocentric distances and locations, and hence a range of metallicities and possibly ages and star formation histories. The fact that the derived distance moduli for their three fields varied by only $\\sim 0.035\\;$mag indicates that the red clump is a potentially stable standard candle. The mostly empirical support for using the red clump stars as a distance indicator should also be verified using modern theory of the stellar structure and evolution. In particular, $I$-band predictions are seldom given by such theoretical calculations.  So why does the red clump distance to the LMC disagree with the Cepheid distance (Madore \\& Freedman 1998)?  As usual, there are several possible answers. Contrary to our arguments given above, there might still be something ``unusual'' about the red clump population in the LMC.  Although the red clump and Cepheid distances to the LMC are discrepant, the distances to M31 derived from the two methods are in excellent agreement ($m-M = 24.471\\pm0.035\\pm0.045$ from Stanek \\& Garnavich~1998 and 24.44$\\pm 0.13$ from Freedman \\& Madore~1990). Another possibility is that the Cepheid distance to the LMC is simply poorly determined, as there are few Cepheids with well determined parallaxes in the {\\em Hipparcos}\\/ catalog.  In their recent study, Madore \\& Freedman (1998) find $\\mu_{0,LMC}= 18.44\\pm 0.35\\;$mag, from a sample of 19 Cepheids observed by {\\em Hipparcos}\\/ with good $BV$ data, and $\\mu_{0,LMC}= 18.57\\pm 0.11\\;$mag, from a sample of only 7 Cepheids with good $BVIJHK$ data. Yet a third possibility, as discussed by Madore \\& Freedman (1998), is that there are other effects on the Cepheid PL relation (e.g. extinction, metallicity, and statistical errors), which are as significant as any reassessment of the zero point based on {\\em Hipparcos}. The metallicity effect on the Cepheid PL relation, determined by Kennicutt et al.~(1998) ($\\delta(m - M)_0/\\delta[O/H] = -0.24 \\pm 0.16 $\\magdex), reduces the discrepancy between the red clump and Cepheid distances to the LMC by $\\sim 0.1\\;$mag, while the somewhat larger metallicity dependence found by Sasselov et al.~(1997) and Kochanek (1997) reduces it by $\\sim 0.15\\;$mag.  To illustrate the effect of the assumed reddening on the derived distance modulus, we note that the value of the LMC distance modulus, $\\mu_{0,LMC}= 18.54\\;$mag, derived recently by Salaris \\& Cassisi (1998) and based on the $V$-band brightness of the RR Lyr stars, becomes $\\mu_{0,LMC}= 18.22\\;$mag if their assumed reddening of $E(B-V)=0.10\\;$mag is increased to $E(B-V)=0.20\\;$mag, corresponding to the mean reddening found by Harris et al.~(1997). It is disturbing that the distance to a key calibrator of the entire distance scale is uncertain by up to 20\\%.   As described by Udalski et al.~(1998), the $\\sim 0.4\\;$mag discrepancy between their (and now our as well) ``short'' distance to the LMC and the ``long'' distance to the LMC from the Cepheids can be resolved by using detached eclipsing binaries as a direct distance indicator (Paczy\\'nski 1997). The Cepheids in the LMC can also be used to get a direct distance estimate through a modified Baade-Wesselink method (e.g. Krockenberger 1996; Krockenberger, Sasselov, \\& Noyes 1997). Both these methods require no intermediate steps in the distance ladder, therefore avoiding the propagation of errors usually crippling the distance scale.  With the 6.5--8 meter telescopes now being built in the Southern Hemisphere the necessary spectroscopy of the detached eclipsing binaries and Cepheids can be quite easily obtained for these 14--18 magnitude stars. It is worth mentioning here that the effort to obtain direct distances with the detached eclipsing binaries and Cepheids to the M31 and M33 galaxies is already under way and the first results look promising (project DIRECT: Kaluzny et al.~1998, Stanek et al.~1998, Krockenberger et al.~1998, Sasselov et al.~1998).  To summarize, among the various stellar distance indicators the red clump giants might be the best for determining the distance to the LMC and other nearby galaxies because there are so many red clump stars. In particular, {\\em Hipparcos}\\/ provided accurate distance determinations for almost 2,000 such stars, but unfortunately $I$-band photometry is available for only $\\sim 30\\%$ of them, so it is important to obtain $I$-band photometry for all {\\em Hipparcos}\\/ red clump giants.  We also need to test the metallicity dependence of {\\em Hipparcos}\\/ red clump giant absolute luminosities and colors. There are many stars within $100\\;pc$ of the Sun for which very high-resolution spectroscopy is possible."
    },
    "9803/astro-ph9803238_arXiv.txt": {
        "abstract": "Recent work by Pringle and by Maloney, Begelman \\& Pringle has shown that geometrically thin, optically thick, accretion disks are unstable to warping driven by radiation torque from the central source. This work was confined to isothermal (\\ie surface density $\\Sigma\\propto R^{-3/2}$) disks. In this paper we generalize the study of radiation-driven warping to include general power-law surface density distributions, $\\Sigma\\propto R^{-\\delta}$. We consider the range $\\delta=3/2$ (the isothermal case) to $\\delta=-3/2$, which corresponds to a radiation-pressure-supported disk; this spans the range of surface density distributions likely to be found in real astrophysical disks. In all cases there are an infinite number of zero-crossing solutions (\\ie solutions that cross the equator), which are the physically relevant modes if the outer boundary of the disk is required to lie in a specified plane. However, unlike the isothermal disk, which is the degenerate case, the frequency eigenvalues for $\\delta\\neq 3/2$ are all distinct. In all cases the location of the zero moves outward from the steady-state (pure precession) value with increasing growth rate; thus there is a critical minimum size for unstable disks. Modes with zeros at smaller radii are damped. The critical radius and the steady-state precession rate depend only weakly on $\\delta$.  An additional analytic solution has been found for $\\delta=1$.  The case $\\delta=1$ divides the solutions into two qualitatively different regimes. For $\\delta \\ge 1$, the fastest-growing modes have maximum warp amplitude, $\\beta_{\\rm max}$, close to the disk outer edge, and the ratio of $\\beta_{\\rm max}$ to the warp amplitude at the disk inner edge, $\\beta_o$, is $\\gg1$. For $\\delta < 1$, $\\beta_{\\rm max}/\\beta_o\\simeq 1$, and the warp maximum steadily approaches the origin as $\\delta$ decreases. This implies that nonlinear effects {\\it must} be important if the warp extends to the disk inner edge for $\\delta \\ge 1$, but for $\\delta < 1$ nonlinearity will be important only if the warp amplitude is large at the origin. Because of this qualitative difference in the shapes of the warps, the effects of shadowing of the central source by the warp will also be very different in the two regimes of $\\delta.$ This has important implications for radiation-driven warping in X-ray binaries, for which the value of $\\delta$ characterizing the disk is likely to be less than unity.  In real accretion disks the outer boundary condition is likely to be different from the zero-crossing condition that we have assumed. In accretion disks around massive black holes in active galactic nuclei, the disk will probably become optically thin before the outer disk boundary is reached, while in X-ray binaries, there will be an outer disk region (outside the circularization radius) in which the inflow velocity is zero but angular momentum is still transported. We show that in both these cases the solutions are similar to the zero-crossing eigenfunctions. ",
        "introduction": "Evidence for warped, precessing accretion disks in astrophysical systems ranging from X-ray binaries to active galactic nuclei has steadily accumulated over the last two decades (see Maloney \\& Begelman 1997a, and references therein). The origin and maintainence of such warped disks has until recently stood as an unsolved theoretical problem. While it is possible, for example, to generate non-planar modes with $m=1$ symmetry in thin, relativistic disks (\\cite{k90}; \\cite{kh91}), these modes only exist at small radii ($R\\lessapprox 10$ Schwarzschild radii), since they rely on trapping of the modes in the non-Newtonian region of the potential. However, an important clue was provided by \\cite{pet77}, who pointed out that in an optically thick disk with a central source of luminosity, the pressure resulting from re-radiation of the intercepted flux will produce a net torque if the disk is warped.  Almost twenty years were to pass before it was recognized that radiation pressure torque actually leads to a warping instability.  \\cite{pri96} (P96) showed that, for the special case in which the disk surface density $\\Sigma\\propto R^{-3/2}$ (corresponding to an isothermal disk in the usual $\\alpha-$disk formalism, with disk viscosity $\\nu\\propto R^{3/2}$), even an initially planar disk is unstable to warping by this mechanism. Pringle solved the linearized twisted disk equations in this case using a WKB approximation.  Maloney, Begelman \\& Pringle (1996, Paper I, hereafter MBP) found exact solutions to the linearized twist equations, and demonstrated the importance of the outer boundary condition for determining the growth rates of the unstable modes.  These previous works all specialized to the case of an isothermal disk, which simplifies the twist equations. While this may be a reasonable approximation for some astrophysical disks (\\eg the masing molecular disk in NGC 4258; see MBP), there are many other systems, such as accretion disks in X-ray binary systems, where this is likely to be a poor assumption. In this paper we extend the work of MBP by considering disks with power-law surface density profiles, $\\Sigma\\propto R^{-\\delta}$. We consider the range $-3/2\\le\\delta\\le 3/2$: the lower limit corresponds to a radiation-pressure-supported disk (\\eg \\cite{fkr92}), while the upper limit is the isothermal value (MBP). Within the limitations of assuming a constant power-law for the surface density, this spans the probable range of surface density laws relevant to real astrophysical accretion disks. For example, the standard Shakura-Sunyaev gas pressure-supported disk is characterized by $\\delta=0.75$ (\\cite{ss73}).  In \\S 2 we discuss the twist equation, including the effect of radiation torque, and cast it into a more convenient form. We solve the equation numerically in \\S 3 and discuss both the time-dependent and steady-state solutions. As in the isothermal case, the outer boundary condition is crucial for determining the stability of the disk and the growth rates of the unstable modes. In \\S 4 we discuss the important issue of the appropriate outer boundary condition for accretion disks around stellar-mass objects and AGN. Finally, in \\S 5 we discuss the implications of the results and their application to real accretion disks. ",
        "conclusions": "Earlier work on the radiation-driven warping instability discovered by Pringle (P96; MBP) considered only the isothermal, $\\delta=3/2$ case. In this paper we have considered more general power-law disk density distributions, from the isothermal disk to $\\delta=-3/2$, corresponding to a radiation-pressure supported disk; this spans the range that is likely to be relevant to astrophysical disks. Although the shapes of the eigenfunctions do change with decreasing $\\delta$, the most important features of the instability are generic. Most importantly, the instability exists over the entire range of surface density index that we have considered, and the critical radius above which disks are unstable to radiation-driven warping changes only by a factor of $\\simeq 6$ from $\\delta=3/2$ to $\\delta=-3/2$. Similarly, the growth and precession rates (in dimensional units) do not depend strongly on $\\delta$ (see the discussion after equation [18] and below). Evaluating equation (15) for the critical radius, \\bea R_{\\rm cr}&=&\\left(5.9\\times10^8\\;\\,\\quad -\\quad 3.5\\times 10^9\\right)\\; \\left({\\eta\\over \\epsilon_{0.1}} \\right)^2\\left({M\\over\\msol}\\right)\\;{\\rm cm} , \\nonumber \\\\ &=& \\left(5.9\\times10^{16}\\quad - \\quad 3.5\\times 10^{17}\\right)\\, \\left({\\eta\\over \\epsilon_{0.1}} \\right)^2\\left({M\\over 10^8 \\msol}\\right)\\;{\\rm cm} , \\eea where the range in numerical values is for $\\delta=3/2$ to $\\delta=-3/2$ and $\\epsilon=0.1\\epsilon_{0.1}$. The only warping modes with zeros at $R < R_{\\rm cr}$ are damped, so that disks that are smaller than $R_{\\rm cr}$ are stable against warping. In consequence of the $\\epsilon^{-2}$ scaling of $R_{\\rm cr}$, accretion disks in systems with very low radiative efficiency will not be unstable to radiation-driven warping unless they are implausibly large. For this reason, this mechanism cannot provide an explanation for the warp in the thin maser disk of NGC 4258 (\\eg \\cite{miy95}, \\cite{her97}) if the inner disk is advection-dominated with $\\epsilon\\sim 10^{-3}$ (\\cite{las96}; see the discussion in MBP), since the maser disk would be far too small for instability in this case. This also indicates that radiation-driven warping generally will not be important in cataclysmic variables or protostellar disks dominated by accretion-powered luminosity, since the radiative efficiency is limited to small values as the stellar surfaces are at $R_*\\gg R_s$ (but see \\cite{arm97} for a discussion of the possible action of the instability in the protostellar case).  To evaluate the typical precession timescales, we need to evaluate the viscous inflow timescale at $R_{\\rm cr}$. Letting $\\nu_1=\\alpha c_s H$, where $c_s$ is the isothermal sound speed and $H$ is the scale height, we can write the viscous timescale as \\be t_{\\rm visc}\\sim {2\\over 3} {R\\over V_\\phi}\\alpha^{-1} (H/R)^{-2} \\ee where $V_\\phi$ is the rotational velocity (assumed to be Keplerian) and $H/R$ is evaluated at the radius in question. Since $t_{\\rm visc}\\propto R^{3/2}$, and $R_{\\rm cr}/R_o=x_{\\rm cr}^2$, \\be t_{\\rm visc}(R_{\\rm cr})\\sim {1\\over 3}\\left({\\eta\\over\\epsilon} \\right)^3 {R_s\\over\\alpha c} x_{\\rm cr}^{3/2}\\left(H/R\\right)^{-2} \\ee where $H/R$ is now evaluated at $R_{\\rm cr}$. Taking the precession timescale $t_{\\rm prec}=2\\pi/\\sigma_r$, where $\\sigma_r$ is given by equation (18), and evaluating the constants, we find \\be t_{\\rm prec}\\sim 12\\, {\\eta^2\\over\\epsilon_{0.1}^3} {M/\\msol\\over \\alpha_{0.1}}\\left({H/R\\over 0.01}\\right)^{-2}_{R_{\\rm cr}}\\;{\\rm days} \\ee with only weak dependence on $\\delta$: the numerical coefficient only varies by a factor of two over the whole range of $\\delta$. Thus the precession timescales for X-ray binary systems (the only systems in which precession can actually be observed) are expected to be of the order of weeks to months.  This is of course the precession timescale for the steady-state modes from linear theory. As discussed in \\S 3, real disks will ordinarily be unwarped beyond some maximum radius, either the physical edge of the disk or where the disk becomes optically thin. This outer boundary, which will not in general correspond to the critical radius, will determine the warp growth rate. We expect that the warp will eventually saturate at some amplitude (but see Pringle 1997). Assuming that the disk does reach a steady state, what will the precession rate be? There is reason to suspect it may not be very different from the linear theory result. Figure 6 shows that, except for growth rates very close to the maximum, the real part of the eigenvalue $\\tilde\\sigma_r$, \\ie the precession rate, is nearly independent of the growth rate. In the isothermal case, in fact, $\\tilde\\sigma_r$ is independent of $\\tilde\\sigma_i$. This suggests that, however different modes may couple in reaching the final state, the precession rate will be similar to the linear steady-state result.  Implicit throughout this paper has been the assumption that the disks are optically thick to both absorption and re-emission, so that they are subject to the radiation-driven warping instability. This requirement imposes a minimum mass accretion rate that must be exceeded for the disk to be optically thick. In Appendix D, we derive this critical mass accretion rate for three different possible sources of opacity in astrophysical disks (electron scattering, dust absorption, and Kramer's opacity) and show that it does not in general place any significant limitations on occurrence of the instability.  As discussed in \\S 3.2, there is one very important systematic change in the nature of the instability with $\\delta$. The difference in the behavior of the growing modes for $\\delta \\ge 1$ and $\\delta < 1$ is of fundamental importance for the evolution of disks warped by radiation pressure. For $\\delta \\ge 1$, the fast-growing modes all have their maximum warp (\\ie tilt $\\beta$) close to the outer edge of the disk, and the amplitude $\\beta_{\\rm max}$ is much greater than $\\beta_o$, the tilt at the origin. This immediately implies that the warp must reach the nonlinear regime when the tilt at small radius is negligible. In this case the evolution of the disk at radii interior to the warp maximum is almost certainly driven by the nonlinear evolution of the outer warp (\\eg \\cite{pri97}), so that nonlinear effects {\\it must} be important if the warp extends to the disk inner edge.  For $\\delta < 1$, the behavior is qualitatively different, as $\\beta_{\\rm max}/\\beta_o$ is always of order unity. In this regime, nonlinearity will be important only if the warp has grown out of the linear regime at the origin. Furthermore, because the shapes of the growing warps in these two regimes are so dissimilar, the effects of shadowing of the central source by the warping of the disk will be very different. These distinctions are liable to be crucial for X-ray binaries such as SS 433 and Her X-1, which show evidence for a {\\it global} precessing warp.  One final point regarding X-ray binary systems must be mentioned. In one of the best-studied systems, Her X-1, the direction of precession of the warp is inferred to be retrograde with respect to the direction of rotation (\\eg \\cite{ger76}) and this has also been suggested for SS 433 (\\cite{leib84}; \\cite{bri89}). As shown in Appendix B, in the absence of external torques the direction of precession of the warp must be prograde. However, the qualification on this statement is extremely important: as pointed out in \\S 4.2, and discussed in detail by Maloney \\& Begelman (1997b), including the quadrupole torque from a companion star allows retrograde as well as prograde solutions to exist.  The zero-crossing outer boundary condition that we have imposed will not be strictly correct in real astrophysical disks. However, as discussed in \\S 4, the solutions that obey the likely realistic outer boundary conditions -- the optically thin outer boundary for accretion disks in active galactic nuclei, and a flat outer boundary for disks in X-ray binaries -- are in all important respects similar to the zero-crossing solutions.  Radiation-driven warping and precession offers a robust mechanism for producing tilted, precessing accretion disks, in accreting binary systems such as Her X-1 and SS 433, and in active galactic nuclei. Because radiation-driven warping is an inherently global mechanism, it avoids the difficulties inherent in other proposed mechanisms for producing warping and precession, \\eg communicating a single precession frequency through a fluid, differentially-rotating disk. This mechanism can thus explain the simultaneous precession of inner disks (as evidenced by the jets of SS 433 and the pulse profile variations of Her X-1) and outer disks (as required to match the periodicities in X-ray flux and disk emission in these same objects).  A full understanding of the nature of the radiation-driven warping instability will require nonlinear simulations of the type presented in Pringle (1997), which will not only allow for inclusion of the nonlinear terms but also inherently nonlinear effects such as shadowing. This will be the subject of future work."
    },
    "9803/astro-ph9803328_arXiv.txt": {
        "abstract": " ",
        "introduction": "A reconstruction of density perturbation spectrum is a key problem of the modern cosmology. It made a dramatic turn after detecting the primordial CMB anisotropy by DMR COBE (Smoot et al\\cite{1}, Bennet et al\\cite{2}) as the signal found at $10^0$ $\\Delta T/T = 1.06 \\times 10^{-5}$ appeared to be few times more than the expectable value of $\\Delta T/T$ in the most simple and developed cosmological model -- standard CDM one (SCDM\\footnote {$\\Omega_M = 1$, $\\Omega_b = 0.06$ (Walker at al\\cite{3}), $\\Omega_{CDM} = 0.94$, $h = 0.5$, no cosmological gravitational waves.}).   Currently there are a lot of experimental data (such as the spatial distributions of galaxies, clusters of galaxies and quasars, bulk velosities, CMB anisotropy, and others) which can be used to reconstruct the density perturbation spectrum. Characteristic scales of data are different and vary from $\\sim 10$ Mpc which is a scale of nonlinearity to the horizon scale. However, it seems now the most crucial tests are large-scale CMB anisotropy and the number of galaxy clusters in {\\it top-hat} sphere with radius $R = 8 h^{-1}$ Mpc = 16 Mpc. The former can be easy related to the amplitude of density perturbations through the SW effect (Sachs \\& Wolfe\\cite{4}): $$ \\frac{\\Delta T}{T}(\\vec e) = \\frac{H_0^2}{2(2\\pi)^{3/2}} \\int \\limits_{-\\infty}^{\\infty} \\frac{1}{k^2} \\delta_{\\vec k} e^{i\\vec k\\vec x}d^3\\vec k,\\;\\;\\;\\; \\vec x \\simeq \\frac{2\\vec e}{H_0}, $$ where $\\delta_{\\vec k}$ is a Fourrier transform of density contrast $\\delta(\\vec x) \\equiv \\delta \\rho / \\rho$, $H_0$ is the Hubble constant, $\\langle \\delta_{\\vec k} \\delta_{\\vec k'} \\rangle = P(k) \\delta(\\vec k -\\vec k')$, $P(k) = A k^{n_S} T^2 (\\Omega_\\nu,k)$ is a power spectrum of density perturbations, $A $ is the normalization constant, $T(\\Omega_\\nu,k)$ is a transfer function. The latter determines the value of biasing parameter $b^{-1} \\equiv \\sigma_R$ for spatially flat Universe: $$ \\sigma_R^2 = \\frac{1}{(2\\pi)^3}\\int_{-\\infty}^\\infty P(k) W^2(kR)d^3 \\vec k, $$ where $W(kR) = \\frac{3}{(kR)^3}(\\sin kR - kR\\cos kR)$ is the Fourrier transform of hop-hat window function.   Obviously, both normalizations are model-dependent, the $\\Delta T/T$ normalization depends on the amplitude of cosmological gravitational wave spectrum on large scale and, therefore, is related to the model of inflation, the $\\sigma_R$ normalization does depend on the nature of dark matter. Here we prefer to fix $\\sigma_{16}$ to consider the relative contribution of gravitational waves at COBE scale T/S as an additional calculable parameter (instead of considering some inflationary model).   Below, we report results based on P\\&S formalism which deals with abundance of gravitationally bounded halos of dark matter. ",
        "conclusions": ""
    },
    "9803/astro-ph9803058_arXiv.txt": {
        "abstract": "We apply a unique gas fraction estimator to published X-ray cluster properties and compare the derived gas fractions of observed clusters to simulated ones.  The observations are consistent with a universal gas fraction of $0.15\\pm 0.01 h_{50}^{-3/2}$ for the low redshift clusters that meet our selection criteria.  The fair sampling hypothesis states that all clusters should have a universal constant gas fraction for all times.  Consequently, any apparent evolution would most likely be explained by an incorrect assumption for the angular-diameter distance relation.  We show that the high redshift cluster data is consistent with this hypothesis for $\\Omega_0<0.63 $ (95\\% formal confidence, flat $\\Lambda$ model) or $\\Omega_0<0.60$ (95\\% formal confidence, hyperbolic open model).  The maximum likelihood occurs at $\\Omega_0=0.2$ for a spatially flat cosmological constant model. ",
        "introduction": "It has been proposed that clusters of galaxies should be a fair sample of baryonic matter (White \\etal 1993).  Rich clusters form through the gravitational collapse of the matter within a 15-30 Mpc diameter volume, where the force of gravity acts equally on all non-relativistic forms of matter.  The richest clusters have temperatures above 10 keV, which corresponds to velocity dispersions in excess of 1000 km/sec.  Most non-gravitational processes do not appear to affect the bulk of matter at comparable energies, and so we would expect the gas and dark matter to collapse into objects where they are fairly represented.  This is confirmed in simulations using a large variety of techniques (Frenk \\etal 1998), where it is found that within the virial radius the gas and dark matter are indeed equally represented with deviations of only 10\\%.  Since clusters of galaxies are observed at cosmological distances and are spatially resolved objects, this opens the possibility of directly measuring angular diameter distances if the gas to dark matter ratio were known in advance (Pen 1997).  In this paper we compare observed and simulated cluster properties and we estimate the errors in our methods. ",
        "conclusions": "Comparing a catalog of local cluster properties with simulations, we find that the data is consistent with a universally constant gas fraction of $f_g=0.15\\pm 0.01 h_{50}^{-3/2}$.  With the present day uncertainty in $0.025 \\lesssim \\Omega_b h_{50}^2 \\lesssim 0.1$ (Schramm and Turner 1997) and some errors in the Hubble constant, we obtain no useful constaint on $\\Omega_0$ using the low redshift clusters.  Independent conservation of baryons and dark matter, however, allows us to constrain $\\Omega_0$ from the apparent evolution of the cluster gas fraction.  We have shown that the 3 clusters with measured gas fractions at $z>0.5$ are inconsistent with an $\\Omega=1$ universe and the fair sampling hypothesis.  For a spatially flat cosmological constant dominated universe, we obtain a bound of $\\Omega_0<0.63$ (95\\% formal confidence) with a best fit value of $\\Omega_0=0.2$.  For a spatially hyperbolic universe with only matter, we find $\\Omega_0<0.60$ (95\\% formal confidence) with the maximum likelihood at $\\Omega_0=0$.  The errors are dominated by the temperature uncertainties in the high redshift clusters, and future observations could reduce the errors by a factor of two.  This work was supported by the National Science Foundation through REU grant AST 9321943, NASA ATP grant TBD and the Harvard Milton Fund. Computing time was provided by the National Center for Supercomputing Applications.  We would like to thank Bill Forman and Christine Jones for providing the cluster X-ray tables."
    },
    "9803/astro-ph9803022_arXiv.txt": {
        "abstract": "The issue of the approximate isotropy and homogeneity of the observable universe is one of the major topics in modern Cosmology: the common use of the Friedmann--Robertson--Walker [FWR] metric relies on these assumptions.  Therefore, results conflicting with the ``canonical'' picture would be of the utmost importance.  In a number of recent papers it has been suggested that strong evidence of a fractal distribution with dimension $D\\simeq2$ exists in several samples, including Abell clusters [ACO] and galaxies from the ESO Slice Project redshift survey [ESP].  Here we report  the results of an independent analysis of the radial density run, $N(<R) \\propto R^D$, of the ESP and ACO data.  For the ESP data the situation is such that the explored volume, albeit reasonably deep, is still influenced by the presence of large structures.  Moreover, the depth of the ESP survey ($z \\la  0.2$) is such to cause noticeable effects according to different choices of k-corrections, and this adds some additional uncertainty in the results.  However, we find that for a variety of volume limited samples the dimensionality of the ESP sample is $D \\approx 3$, and the value $D = 2$ is always excluded at the level of at least five (bootstrap) standard deviations. The only way in which we reproduce $D \\approx 2$ is by both unphysically ignoring the galaxy k--correction and  using Euclidean rather than FRW cosmological distances.  In the cluster case the problems related to the choice of metrics and k--correction are much lessened, and we find that ACO clusters have $D_{ACO} = 3.07 \\pm 0.18$ and $D_{ACO} = 2.93 \\pm 0.15$ for richness class $\\cR \\geq 1$ and $\\cR \\geq 0$, respectively. Therefore $D=2$ is excluded with high significance also for the cluster data. ",
        "introduction": "Recently Pie\\-tro\\-ne\\-ro and collaborators [hereafter P\\&C] (Pie\\-tro\\-ne\\-ro et al. \\cite{Pprinc}, Sylos Labini et al. \\cite{SLGMP}, Baryshev et al. \\cite{SLPvistas}) argued that a large number of samples give strong evidence that the distribution of clusters and galaxies is a simple fractal of dimension $D \\simeq 2$ \\footnote{It must be noticed that this value is different from the value $D\\simeq 1.6$, supported by the same authors in the past years (Coleman \\& Pie\\-tro\\-ne\\-ro \\cite{CP}), and $D=2$ has been often discussed in the literature, because it would naturally arise \\emph{locally} from a geometrical dominance of planar structures, such as ``pancakes'', see f.i. Guzzo et al. (\\cite{gigi+91}).} on scales of several hundreds of Mpc, quite different from the ``canonical'' value of $D=3$.  This is an important claim because of its far reaching implications, and this issue, which dates back to the beginning of the century (Charlier, 1908, see Peebles \\cite{peebles}), can be properly addressed with long and careful analyses which, however, demand much deeper and better samples than those presently available in order to be able to explore an adequate range of scales (see e.g. Mc~Cauley \\cite{McCauley97}, Hamburger et al. \\cite{Hamburger+96}). While 2D constraints come from fluctuations in cosmic backgrounds (see e.g. Peebles 1993), historically much work on this subject has been done for decades in analyzing galaxy counts which, in a non evolving Euclidean Universe, should follow $N(m) \\propto 10^{0.6 m}$.  Indeed, this behaviour has been observed in an intermediate magnitude range, $m \\approx 15 \\div 17$, (see e.g. Sandage \\cite{sandage}). However, on the bright side one has to deal with magnitude errors given from saturation of photographic plates and/or the relatively small volumes sampled which also reflect in uncertainties in locally the derived space galaxy density (cf Loveday et al. \\cite{SAPM}, Zucca et al. 1997, Maddox \\cite{maddox}), while on the faint side cosmological curvature and evolutionary effects become dominant and difficult to disentangle (Koo \\& Kron \\cite{kookron}, Ellis \\cite{ellis}). Therefore one really needs redshift information in order to avoid effects of time--space projections.  In this paper we limit ourselves to an independent check on two of the samples discussed by P\\&C, the ESP galaxies (Vet\\-to\\-la\\-ni et al. \\cite{paolo+97}) and the ACO clusters (Abell, Corwin \\& Olowin \\cite{ACO}), without touching upon the very general issue of a possible fractal distribution of the matter in the universe and its consequences: the interested reader can consult Peebles (\\cite{peebles}), Coleman \\& Pie\\-tro\\-ne\\-ro (\\cite{CP}), Stoeger et al. (\\cite{SEH}), Ehlers \\& Rindler (\\cite{ER}), Szalay \\& Schramm (\\cite{SS}), Luo \\& Schramm (\\cite{LS}), Mc~Cauley (\\cite{McCauley97}), Buchert (\\cite{Buchert97}), Guzzo (\\cite{gigi97}) and references therein.  However, if the evidence were present in the data at the highly significant level claimed by P\\&C, the simple analysis presented here should be more than adequate in confirming the $D=2$ claims.  In section 1 we present the formalism, in section 2 the results from ESP galaxies, in section 3 the results from the ACO clusters, and finally in section 4 the conclusions. ",
        "conclusions": "On the basis of the present analysis we do not find any evidence in the ESP and ACO samples for a fractal exponent $D\\approx 2$ as claimed by P\\&C. For the ESP sample we showed that $D\\approx 2$ can be obtained only by neglecting the k--correction term: in our opinion to neglect this term would be both unphysical and unjustified. Moreover, also the cluster sample, which is not significantly affected by this particular uncertainty, clearly excludes the value $D=2$.  On the contrary, we find evidence that within the errors both samples show a behaviour on large scales both consistent with and supporting the canonical value, namely $D=3$. It must be stressed that our results use direct depth information differently from indirect arguments such as scaling of angular correlation functions or fluctuations of cosmic backgrounds (see e.g. Peebles \\cite{peebles}).  The above conclusions suggest the need of further and more careful and sophisticated studies on the claims of a strong evidence for a fractal matter distribution from the data sets we analyzed. Also further independent checks should be done and possibly on other, better suited samples."
    },
    "9803/astro-ph9803214_arXiv.txt": {
        "abstract": "We present the serendipitous discovery of 16 ms pulsed X-ray emission from the Crab-like supernova remnant \\snr\\ in the Large Magellanic Cloud. This is the fastest spinning pulsar associated with a supernova remnant (SNR). Observations with the {\\it Rossi} X-ray Timing Explorer (\\xte), centered on the field containing SN1987A, reveal an X-ray pulsar with a narrow pulse profile.  Archival \\asca\\ X-ray data confirm this detection and locate the pulsar within $1^{\\prime}$ of the supernova remnant \\snr, $14^{\\prime}$ from SN1987A. The pulsar manifests evidence for glitch(es) between the \\xte\\ and \\asca\\ observations which span 3.5 years; the mean linear spin-down rate is ${\\dot P} = 5.126 \\times 10^{-14} \\ {\\rm s \\ s^{-1}}$. The background subtracted pulsed emission is similar to other Crab-like pulsars with a power law of photon index of $\\sim 1.6$. The characteristic spin-down age ($\\sim 5000$ years) is consistent with the previous age estimate of the SNR.  The inferred $B$-field for a rotationally powered pulsar is $\\sim 1 \\times 10^{12}$ Gauss.  Our result confirms the Crab-like nature of \\snr ; the pulsar is likely associated with a compact X-ray source revealed by ROSAT HRI observations. ",
        "introduction": "Crab-like supernova remnants (SNRs) play a critical role in our understanding of young pulsars and their pulsar wind nebulae. These rare SNRs which contain central pulsars are distinguished by their centrally-filled morphologies and non-thermal X-ray spectra.  The X-ray emission of such SNRs comes predominantly from a synchrotron nebula powered by an embedded young pulsar (e.g., Seward 1989). Only three confirmed members of this class were previously known: the Crab Nebula with its famous 33 ms pulsar PSR B0531$+$21, the SNR B0540$-$693 in the Large Magellanic Cloud (LMC) with its 50 ms pulsar \\lmcpsr, and the nebula around the 150 ms pulsar PSR B1509$-$58.  Several candidate Crab-like SNR have been reported. These have similar morphologies and spectral characteristics but so far lack a detected pulsar.  Young Crab-like SNRs are expected to evolve into composite SNRs, in which the thermal emission from shock-heated gas will rival or exceed the diminishing X-ray radiation from the pulsar-powered synchrotron nebulae. This evolutionary sequence is not yet firmly established, though, mainly because of the limited number of samples.  We report in this Letter the discovery of an ultra fast pulsar in \\snr\\ (also known as NGC\\sp 2060, SNR 0538$-$69.1, \\& 30\\sp Dor\\sp B; Henize 1956), confirming the Crab-like nature of this remnant (Wang \\& Gotthelf 1998 and ref. therein). The pulsar was first detected serendipitously with \\xte\\ data while searching for a pulsed signal from nearby SN1987A (Marshall \\etal\\ 1998). We used archival \\asca\\ data of the region to confirm the detection of the pulsed emission, to measure its spin-down rate, and to locate the emission to within the boundary of \\snr.  Based on the \\asca\\ position, we will refer to the pulsar as \\psr.  In the following, we present our discovery of \\psr\\ and discuss its properties in the context of other young pulsars. ",
        "conclusions": "We have found 16 ms pulsed X-ray emission from the 30 Doradus region of the LMC.  We have shown that these pulsations are associated uniquely with the X-ray emission from the SNR \\snr.  This remnant is an important new addition to the class of Crab-like SNRs.  It provides a rare laboratory for studying both an unusually rapidly-rotating pulsar and its relativistic wind, as well as the structure and evolution of a neutron star. Future observations are important to constrain the age of the pulsar, the period second derivative, and to measure putative glitches."
    },
    "9803/astro-ph9803164_arXiv.txt": {
        "abstract": "s#1#2#3{{ \\centering{\\begin{minipage}{4.5in}\\baselineskip=10pt\\footnotesize \\parindent=0pt #1\\par \\parindent=15pt #2\\par \\parindent=15pt #3 \\end{minipage}}\\par}}  \\def\\keywords#1{{ \\centering{\\begin{minipage}{4.5in}\\baselineskip=10pt\\footnotesize {\\footnotesize\\it Keywords}\\/: #1 \\end{minipage}}\\par}}  \\newcommand{\\bibit}{\\nineit} \\newcommand{\\bibbf}{\\ninebf} \\renewenvironment{thebibliography}[1] {\\frenchspacing \\ninerm\\baselineskip=11pt \\begin{list}{\\arabic{enumi}.} {\\usecounter{enumi}\\setlength{\\parsep}{0pt} \\setlength{\\leftmargin 12.7pt}{\\rightmargin 0pt} % \\setlength{\\itemsep}{0pt} \\settowidth {\\labelwidth}{#1.}\\sloppy}}{\\end{list}}  \\newcounter{itemlistc} \\newcounter{romanlistc} \\newcounter{alphlistc} \\newcounter{arabiclistc} \\newenvironment{itemlist} {\\setcounter{itemlistc}{0} \\begin{list}{$\\bullet$} {\\usecounter{itemlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  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\\def\\runninghead#1#2{\\pagestyle{myheadings} \\markboth{{\\protect\\footnotesize\\it{\\quad #1}}\\hfill} {\\hfill{\\protect\\footnotesize\\it{#2\\quad}}}} \\headsep=15pt  \\font\\tenrm=cmr10 \\font\\tenit=cmti10 \\font\\tenbf=cmbx10 \\font\\bfit=cmbxti10 at 10pt \\font\\ninerm=cmr9 \\font\\nineit=cmti9 \\font\\ninebf=cmbx9 \\font\\eightrm=cmr8 \\font\\eightit=cmti8 \\font\\eightbf=cmbx8 \\font\\sevenrm=cmr7 \\font\\fiverm=cmr5  \\newtheorem{theorem}{\\indent Theorem}  \\newtheorem{lemma}{Lemma}  \\newtheorem{definition}{Definition} \\newtheorem{corollary}{Corollary}  \\newcommand{\\proof}[1]{{\\tenbf Proof.} #1 $\\Box$.}     \\textwidth=5truein \\textheight=7.8truein  \\def\\qed{\\hbox{${\\vcenter{\\vbox{\t\t\t% \\hrule height 0.4pt\\hbox{\\vrule width 0.4pt height 6pt \\kern5pt\\vrule width 0.4pt}\\hrule height 0.4pt}}}$}}  \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}\t%  \\begin{document}  \\runninghead{Anisotropy in the Propagation of Radio Wave Polarizations $\\ldots$} { Anisotropy in the Propagation of Radio Wave Polarizations $\\ldots$}  \\normalsize\\textlineskip \\thispagestyle{empty} \\setcounter{page}{1}  \\copyrightheading{}\t\t\t%  \\vspace*{0.88truein}  \\fpage{1} \\centerline{\\bf ANISOTROPY IN THE PROPAGATION OF RADIO POLARIZATIONS  } \\vspace*{0.035truein} \\centerline{\\bf FROM COSMOLOGICALLY DISTANT GALAXIES} \\vspace*{0.37truein} \\centerline{\\footnotesize PANKAJ JAIN} \\vspace*{0.015truein} \\centerline{\\footnotesize\\it Physics Department, I.I.T. Kanpur} \\baselineskip=10pt \\centerline{\\footnotesize\\it  Kanpur, India - 208016} \\vspace*{10pt} \\centerline{\\footnotesize JOHN P. RALSTON} \\vspace*{0.015truein} \\centerline{\\footnotesize\\it Department of Physics and Astronomy, Kansas University} \\baselineskip=10pt \\centerline{\\footnotesize\\it Lawrence, KS-66045, USA} \\vspace*{0.225truein} \\publisher{(received date)}{(revised date)}  \\vspace*{0.21truein} \\abstracts{ Radiation traversing the observable universe provides powerful ways to probe anisotropy of electromagnetic propagation. A controversial recent study claimed a signal of dipole character.  Here we test a new and independent data set of 361 points under the null proposal of {\\it statistical independence} of linear polarization alignments relative to galaxy axes, versus angular positions.  The null hypothesis is tested via maximum likelihood analysis of best fits among numerous independent types of factored distributions.  We also examine single-number correlations which are parameter free, invariant under coordinate transformations, and distributed very robustly.  The statistics are shown explicitly not to depend on the uneven distribution of sources on the sky. We find that the null proposal is not supported at the level of less than 5\\% to less than 0.1\\% by several independent statistics. The signal of correlation violates parity, that is, symmetry under spatial inversion, and requires a statistic which transforms properly. The data indicate an axis of correlation, on the basis of likelihood determined to be $[{\\rm R.A.}=(0^{\\rm h},9^{\\rm m}) \\pm (1^{\\rm h},0^{\\rm m})$, ${\\rm Decl.} = -1^o\\pm 15^o]$. }{} {}  \\vspace*{1pt}\\textlineskip \\textheight=7.8truein \\setcounter{footnote}{0} \\renewcommand{\\thefootnote}{\\alph{footnote}} ",
        "introduction": "\\noindent The orientation of linear radio polarizations emitted by cosmologically distant galaxies has a consistent relation with the galaxy symmetry axis. Exceedingly small physical effects accumulate during propagation, which conventional measurements can directly probe.  Thus electromagnetic radiation traversing the observable universe can detect subtle forms of cosmological anisotropy. A signal with dipole character was claimed recently$^1$ from an analysis of published radio data. Analysis found an ``anisotropy axis\" $\\vec s_{\\rm NR}=(21^{\\rm h}\\pm 2^{\\rm h}, 0^o\\pm 20^o$) governing orientation of polarization of the radio signals varying in a coherent way across the dome of the sky. The origin of this behavior is not clear, and may or may not indicate a fundamental anisotropy on a scale larger than previously found in cosmology.   There is a long history of puzzling observations. Beginning in the 1960's observers noticed that Faraday-subtracted polarizations were distributed in peculiar ways relative to the source axes. In 1982 Birch$^2$ empirically observed a coherent angular anisotropy in the off-sets of the polarization and galaxy axes, using a data set of 137 points.  Birch's statistical methods were questioned, but more sophisticated studies$^{3,4}$ confirmed surprisingly strong signals in Birch's data.  The statistics were not consistent with isotropy at 99.9\\% and 99.98\\% confidence levels, respectively.  One of the same groups$^{4,5}$ went on to create an independent set of 277 points and simultaneously introduced a different statistical measure. They obtained no signal in this set and dismissed Birch's results. This left unresolved the puzzling fact that his data had contained a signal at such a high level of statistical significance. When Nodland and Ralston,$^1$ initially unaware of Birch's$^2$ work, independently found a statistically significant signal in an independent set of 160 points, criticisms focused on proposing different statistical baselines$^{6,7,8}$ and again claimed to find no signal of anisotropy.  The question of systematic bias in such data had been raised by the authors$^1$ (henceforth $NR$) and earlier$^9$ regarding Birch.   Here we report analysis of a considerably larger data set which contains 361 points.  We have taken into account criticisms and experience from earlier work, and used the most robust statistical methods available. New progress has been made by paying close attention to the symmetries of the problem. The usual expectation of {\\it independence} of the polarization and sky angular coordinates, or ``uncorrelated isotropy'', happens to represent a definite symmetry, which is that the distribution factors.  The classic scientific method becomes applicable: we can test isotropy as a clean hypothesis and see if it can be ruled out, which is immensely powerful. We use generic methods to represent the correlations, emphasising the symmetry that they are {\\it odd} in the polarization variable at hand, which is a consequence of parity (spatial inversion) symmetry.$^{10}$ This simple point resolves many apparent discrepancies between the previous studies. Rather than being at odds with one another, all the facts are now found to be consistent; we know of nothing in contradiction to our conclusions.    The data collects variables from cosmologically distant galaxies, as compiled in the literature.$^{2,4,5,11,12,13}$ The data set by $NR$ reproduced that of Carroll et al$^{13}$ except for a half-dozen corrections from the original literature. The compilation of Eichendorf and Reinhardt,$^{11,12}$ available on the NASA-ADC archives, contain numerous sources for which the position angle of the source is listed. We obtained the polarization angle for these sources from Simard-Normandin et al$^{14}$  for all the sources for which they were available. We compiled a total of 152 data points in this fashion. Taking these as our primary data set we added any distinct data points contained in Bietenholz$^5$, making a total of 313 points.  Data points were regarded as distinct if they had different Right Ascension, and differed in Declination by more than one degree (which can be attributed to change in convention). This set was further combined with the $NR$ and remaining distinct points of the Birch data, in that order, making a total of 361 data points. In combining these different data sets, we verified that the polarization off-set values for points with coincident Right Ascension and Declination did not differ by more than a few degrees for most of the data. Specifically we found that the disagreement exceeded $5^o$ only for very few points, which if deleted made no difference to our final results. We also verified consistency using a newer 1988 compilation by Broten et al.$^{15}$ The only exception to this rule was found for Birch's data: here the disagreement with other compilations was found to be larger, but still tolerable.  All results we report are consistent, and no combination of any large set gave results significantly different from any other. The absence of information available to us on Birch's $RM$ values, plus the possibility of discrepancies in that data, led us to give results both with and without Birch's data.  In Figure 1 we show the angular distribution of data, which naturally is not isotropic due to the zone of avoidance and dominance of Northern Hemisphere measurements.  We will exhaustively show that the angular distribution is not an issue and cannot be confused with correlation.  \\medskip  \\begin{figure}[t,b] \\hbox{\\hspace{0em} \\hbox{\\psfig{file=anisa.ps,height=8cm}}} \\caption{ Aithoff-Hammer equal-area plots of the distribution of sources on the dome of the sky, in the equatorial coordinates of the data used. The distribution is somewhat non-uniform due to the zone of avoidance and dominance of Northern Hemisphere measurements.  (a) The distribution of the full data set of 332 points, excluding the 29 extra points contained in Birch's compilation. Adding Birch's data makes the set even more uniform. (b) The distribution of the same data set after the cut on rotation measure, $|RM-\\overline{RM}| >6$.  Any non-uniformity of the angular distribution is taken into account in all statistics reported. } \\end{figure}  \\medskip  The observables listed for galaxy $i$ include a major axis orientation angle $\\psi_i$, a linear polarization angle $\\chi_i$, and the angular coordinates of the galaxies on the sky. Other variables may include a resolution parameter, degree of polarization, and the Faraday rotation measure $RM$. The rotation measure is the slope of plots of measured polarization angle versus wavelength-squared.  This is known to measure intervening magnetized plasma parameters. A-priori, $RM$ has nothing to do with the variable $\\chi$, which is the polarization angle after Faraday rotation is subtracted.  However we have retained this variable, which seems to be informative. Consistent with restricting the study to uncorrelated isotropy, we integrate over the redshift, which happens to be incomplete in the data set in any event.  We let $\\beta=\\chi-\\psi$ be the angle between the plane of polarization and the symmetry axis of the source.  The variables $ \\chi$ and $\\psi$ are determined up to a multiple of $\\pi$; $\\beta$ runs from $-\\pi$ to $\\pi$.$^{10}$ To deal with the $\\pi$ ambiguity of polarization and axis measurements, one can map $\\beta \\rightarrow Y(\\Omega)$, where $\\Omega$ is a variable defined on twice the interval.  A popular map is ``{\\it Map 1}'', $\\Omega_1(\\beta) = 2 \\beta$.  The function $Y$ is represented by a Fourier series with periodicity $2\\pi$, assuring that the transformation $\\beta \\rightarrow \\beta^\\prime = \\beta \\pm \\pi$ leaves $Y(\\Omega)$ invariant.  The first Fourier components create a 2-component vector-like object $\\vec Y(\\Omega) = (\\cos(\\Omega), \\sin(\\Omega))$. When the components of $\\vec Y(\\Omega)$ are used in statistical analysis, there is naturally a Jacobian factor which represents the choice of {\\it Map}.  By no means, then, is {\\it Map 1} sacred, and other maps are discussed below. The angular positions on the dome of the sky are mapped into their 3-dimensional Cartesian vector positions $\\vec X$ on a unit sphere.  Since we do not model this distribution, but take it from the data, this standard map is adequate. When coordinate origins are changed, the components of $\\vec X$ transform by standard rules; one can go on to make nicely transforming distributions and tensor correlations. The two choices measuring $\\chi$ relative to $\\psi$ or $\\psi+\\pi/2$ correspond to $\\vec Y\\rightarrow -\\vec Y$.  This does not mix the 2 components of $\\vec Y$, which will be called ``even'' (for $\\cos(\\Omega)$) and ``odd'' (for $\\sin(\\Omega)$) following the transformation property of being even or odd, respectively, under parity (spatial inversion). As discussed in detail elsewhere,$^{10}$ functions of the offset angles have the corresponding parity if they are even or odd functions of $\\beta$, as intuitively evident from the handed ``sense of twist'' a parity-odd quantity conveys. The invariant correlations discussed below avoid any question of coordinate origin (either in polarization quantities or in angular positions on the dome of the sky) by being totally independent of the choice of angular origin.   The standard assumption of statistical independence corresponds to a distribution $g(\\Omega,\\vec X)=h(\\Omega) f(\\vec X )$.  This is a very broad class of distributions, with $h(\\Omega)$ and $f(\\vec X )$ completely unrestricted, which nevertheless has symmetries allowing it to be tested. All statistics will be compared to baselines using the actual distribution of the data $f(\\vec X)$ on the dome of the sky in Monte Carlo simulations.  Statistics based on assuming independence of polarizations and positions will be compared with a simple correlated ansatz of the form $h(\\Omega) C(\\Omega, \\vec X) f(\\vec X)$. The case $C=1$ reduces to the uncorrelated case. ",
        "conclusions": "\\noindent In presenting a study of restricted scope, our conclusions are most crisply phrased in a negative sense: {\\it  the null hypothesis of uncorrelated isotropy is not supported. On the basis of significance, it can be ruled out}.  By the nature of this study, one is constrained from concluding prematurely what the correlation found may represent. Under many separate statistical probes, the evidence against isotropy in the data is significant at $95\\%-99.99\\%$ (roughly $2-4\\sigma$) confidence levels.   This is not the first such finding, but just one more among a number of studies accumulating over the years.   While no evidence of systematic bias is found, we strongly reiterate the possibility. Yet the persistence of the effect seems to indicate physical processes outside the framework which has been used to interpret the data conventionally.  Associated with this behavior are persistent axis parameters concordant with the axis parameters found in Nodland and Ralston,$^1$ and which subsequently have been found to coincide with the CMB dipole direction.$^{20,21}$  Nevertheless this is a new field and it would be premature to fix on a physical origin now.  We therefore postpone more detailed conclusions, and recommend that physical models be used to suggest suitable directions of research.  Local effects, while traditionally held to be under control, can potentially be ruled out with redshift information.   Resources exist to generate cosmological radio data sets with many more points, and the time may be ripe for clever technological advances that could be revolutionary.  New analysis combined with new data might tell us what is causing the effect.  \\bigskip  \\nonumsection{Acknowledgements} \\noindent We thank Borge Nodland, Hume Feldman, Doug McKay and G. K. Shukla for useful comments. Supported by DOE grant number 85 ER40214, the KU General Research Fund, the NSF-K*STAR Program under the Kansas Institute for Theoretical and Computational Science and DAE grant number DAE/PHY/96152.    \\nonumsection{References} \\noindent"
    },
    "9803/astro-ph9803346_arXiv.txt": {
        "abstract": "We present high-resolution imaging of the young binary, T Tauri, in continuum emission at $\\lambda$=3~mm.  Compact dust emission with integrated flux density 50 $\\pm$ 6 mJy is resolved in an aperture synthesis map at 0\\farcs5 resolution and is centered at the position of the optically visible component, T Tau N\\null.  No emission above a 3$\\sigma$ level of 9 mJy is detected 0\\farcs7 south of T Tau N at the position of the infrared companion, T~Tau~S\\null.  We interpret the continuum detection as arising from a circumstellar disk around T~Tau~N and estimate its properties by fitting a flat-disk model to visibilities at $\\lambda$ = 1 and 3~mm and to the flux density at $\\lambda$ = 7~mm.  Given the data, probability distributions are calculated for values of the free parameters, including the temperature, density, dust opacity, and the disk outer radius.  The radial variation in temperature and density is not narrowly constrained by the data. The most likely value of the frequency dependence of the dust opacity, $\\beta$ = $0.53^{+0.27}_{-0.17}$, is consistent with that of disks around other single T Tauri stars in which grain growth is believed to have taken place. The outer radius, R = 41$^{+26}_{-14}$~AU, is smaller than the projected separation between T Tau N and S, and may indicate tidal or resonance truncation of the disk by T~Tau S\\null.  The total mass estimated for the disk, log(M$_D$/M$_\\odot$) = ${-2.4}^{+0.7}_{-0.6}$, is similar to masses observed around many single pre--main-sequence sources and, within the uncertainties, is similar to the minimum nebular mass required to form a planetary system like our own. This observation strongly suggests that the presence of a binary companion does not rule out the possibility of formation of a sizeable planetary system.   {\\it Subject headings:} circumstellar matter --- stars:pre-main sequence --- star:individual (T Tauri) ",
        "introduction": "Observations of T~Tau at $\\lambda$ = 2.8 mm were carried out from 1996 December to 1997 March with the 9-element BIMA array.  Continuum emission was measured in an 800 MHz bandwidth centered at 108 GHz. Data were taken with the array in the A, B, and C configurations, providing baseline coverage from 2.1 k$\\lambda$ to 420 k$\\lambda$ with sensitivity to emission on size scales up to $\\sim$60\\arcsec. Interleaved observations of a nearby quasar, 0431+206, were included as a check on the phase de-correlation on long baselines.  The phase calibrator was 0530+135, with an assumed flux of 3.1~Jy during A array.  Data were calibrated and mapped using the Miriad package.  A map of 0431+206 yielded an image of a point source, indicating little atmospheric degradation of the resolution in the observations of T~Tau.  An aperture synthesis image of T Tau was constructed with data from the A array alone (the longest baselines) and is displayed in Figure \\ref{map}.  The beam size is $0\\farcs59 \\times 0\\farcs39$ at a position angle of 48\\arcdeg. It is clearly evident that the compact 3~mm emission arises from circumstellar material surrounding T~Tau~N only.  Peak emission of 32 mJy beam$^{-1}$ is centered at RA(J2000) 04:21:59.424 and Dec(J2000) 19:32:06.41 with an absolute positional uncertainty of $\\sigma$ = $\\pm$0\\farcs07 as determined from observations of 0431+206.  The peak position is within 1.6$\\sigma$ of the Hipparcos coordinates for T~Tau~N, but 7.8$\\sigma$ distant from the position of T~Tau~S\\null.  The emission is resolved with an approximate deconvolved source size of 0\\farcs45 $\\times$ 0\\farcs32 (FWHM) (63 $\\times$ 45~AU) at PA 19\\arcdeg, assuming an elliptical Gaussian shape.  The integrated flux density, 50 $\\pm$ 6 mJy (where the uncertainty is dominated by the flux calibration error), is consistent with the 100 GHz value measured at lower resolution by Momose et al.\\ (1996), 48 $\\pm$ 7 mJy.  No emission is detected at the position of T~Tau~S above the 3$\\sigma$ upper limit of 9 mJy beam$^{-1}$, and no circumbinary emission is apparent.  The integrated flux density detected in the compact C array is no greater than that for A array. Since the C array observations are most sensitive to emission on larger spatial scales, up to 60$''$, the absence of detectable excess emission implies that the majority of continuum emission detected in a 1$'$ aperture at $\\lambda$ = 3 mm originates from a circumstellar region around T Tau N\\null.  The 3$\\sigma$ upper limit on circumbinary emission, accounting for thermal noise and flux calibration errors, is 17 mJy.  Our measurement of the 108 GHz flux density from T~Tau~N is plotted in Figure \\ref{sed} together with high-resolution measurements at 43 GHz ($\\lambda$ = 7~mm) (Koerner et al.\\ 1998), 267 GHz ($\\lambda$ = 1~mm) and 357 GHz ($\\lambda$ = 0.8~mm) (Hogerheidje et al.\\ 1997).  All four points can be fitted by a single power-law curve with $\\chi^2$ = 1.8 and goodness of fit 0.4.  The best-fit spectral index $\\alpha$=d(log($F_\\nu$))/d(log($\\nu$)) is 2.30$\\pm$0.1.  This value is in good agreement with those from disks around single T Tauri stars and consistent with thermal emission from large circumstellar dust grains (Beckwith \\& Sargent 1991; Mannings \\& Emerson 1994; Koerner et al.\\ 1995), suggesting that the continuum emission from T Tau N has a single origin in thermal dust emission along the entire wavelength range from $\\lambda$ = 1 to 7~mm.  We argue that the 3mm emission arises from a circumstellar disk around T Tau N\\null.  The spectral slope is consistent with radiation from dust grains (see Koerner et al.\\ 1998 for a more stringent limit on the possible contribution of free-free or gyrosynchrotron emission). The star is detected optically, even though the lower limit to the dust mass is high; if the dust were distributed in a uniform density sphere with a size given by the A array fit, the extinction to the star would be A$_V \\sim$ 1000.  \\section {Circumstellar disk models for T~Tau~N}  \\subsection{Model description}  To refine estimates of the parameters of dust around T~Tau~N, we fit the visibility amplitudes directly with a model of a circumstellar disk.  While computationally intensive, fitting in the visibility plane rather than the image plane avoids the non-linear process of deconvolution and allows the instrumental errors to be included in a consistent manner.  Our disk model is thin and circularly symmetric with a flux from an annular region at radius $r$ given by \\begin{equation} dS_{\\nu}(r) = {2 \\pi \\cos\\theta \\over D^2 } B_{\\nu} (1 - e^{-\\tau}) r \\, dr, \\label{disk_flux} \\end{equation} where $B_{\\nu}$ is the Planck function, $D$ is the distance, and $\\theta$ the inclination angle.  Flared disks have been invoked to explain the flat infrared spectral indices of some T Tauri stars (Kenyon \\& Hartmann 1987); however, the effects due to flaring are not important at these long wavelengths (Chiang \\& Goldreich 1997). The visibility amplitude $V$ at $uv$-distance $\\eta$ is calculated with a Hankel transform, \\begin{equation} V(\\eta) = 2\\pi \\int S_{\\nu}(r)J_0(2\\pi\\eta r/D)r\\,dr, \\label{hankel} \\end{equation} where $J_0$ is the Bessel function. The inner radius is set by the dust destruction temperature (2000~K); the exact value has little effect on the millimeter flux density. The temperature, surface density and dust opacity are described by power-law relations: $ T(r) = T_{\\rm 10~AU} (r/{\\rm 10~AU})^{-q},\\ \\Sigma (r) = \\Sigma_{\\rm 10~AU} (r/{\\rm 10~AU})^{-p},$ and $\\kappa_{\\nu} = \\kappa_o (\\nu/\\nu_o)^{\\beta}$.  Due to its dependence on the product of surface density and dust opacity, the optical depth scales with the corresponding power-law exponents in radius and frequency, \\begin{equation} \\tau (r,\\nu) = {\\Sigma \\kappa_{\\nu} \\over \\cos\\theta} \\equiv \\tau_{\\rm 10~AU} \\Big({r \\over {\\rm 10~AU}}\\Big)^{-p}\\Big({\\nu \\over \\nu_o}\\Big)^{\\beta}. \\end{equation} The reference value for $\\kappa_o$ is 0.1~g$^{-1}$~cm$^{2}$ at $\\nu_o = 1200$ GHz ($\\lambda$ = 250~$\\mu$m; Hildebrand 1983).  Eight parameters were varied in the model-data comparison: $T_{\\rm 10~AU}$, $ q$, $\\tau_{\\rm 10~AU},$ $p,$ $\\beta,$ $r_{out},$ $\\theta$, $\\mbox{and}\\ \\alpha$, the position angle of the disk on the sky.  In addition to the 3~mm visibilities, those measured at 1~mm by Hogerheidje et al.\\ (1997) were used together with the 7~mm flux density (Koerner et al.\\ 1998).  The IRAS 100 $\\mu$m flux of 120 Jy (Strom et al.\\ 1989) was included as an upper limit.  For each value of $\\theta$ and $\\alpha$, the $u$ and $v$ coordinates for each visibility were de-projected to those of a face-on disk, then binned in annuli of de-projected $uv$-distance for comparison to model values calculated by Eqn.\\ \\ref{hankel}. Over 5 million models were compared to the data within an 8-dimensional grid in parameter space. Logarithmic grid spacings were used for the temperature, optical depth and outer radius.  To quantitatively assess the reliability of estimates of properties of a disk around T Tau N, we calculate the probability distribution for each parameter over the entire range of models.  A detailed description of this Bayesian approach is given in Lay et al.\\ (1997).  Given the data, the probability of a model with a particular set of parameter values is proportional to $e^{-\\chi^2}$ where $\\chi^2$ is the standard squared difference between data and model, weighted by the uncertainty in the data.  A systematic error in overall flux calibration was accounted for by normal weighting of a range of model flux scalings with 1$\\sigma$ corresponding to a 10\\% difference in flux. The final probability for a given model was taken as the sum of probabilities over all flux scalings and multiplied by a factor of $\\sin \\theta$ to account for the fact that edge-on disks are more likely than face-on in a randomly oriented sample.  Finally, the relative likelihood of each parameter value was calculated by adding the probabilities of all models with that parameter value.   \\subsection{Parameter results}  The resulting probability distributions for the disk parameters are given below.  The ranges considered for $T_{\\rm 10~AU}$, $\\tau_{\\rm 10~AU}$, and $r_{out}$ were sufficiently wide to bracket all values with significant probability.  Parameters values quoted are at the median of the probability distribution with an error range encompassing 68\\% of the total probability.  However, it is important to keep in mind that the distribution is not Gaussian.  The disk outer radius is $r_{out} = 41^{+26}_{-14}$~AU (Figure \\ref{prob}).  The probability that it exceeds the projected binary separation (100 AU) is only 3\\%.  Models with large outer radii generally have higher values of $p$ and lower values of $\\tau_{\\rm 10~AU}$.  The value for the dust mass opacity index, $\\beta = 0.53^{+0.27}_{-0.17}$ (Figure \\ref{prob}), is consistent with values measured for circumstellar disks around T Tauri stars (Beckwith \\& Sargent 1991; Koerner et al.\\ 1995).  There is an 85\\% probability that $\\beta$ $\\ge 0.30$, the value calculated by assuming optically thin emission (S$_\\nu \\propto \\nu^{2 + \\beta}$) and fitting a straight line to the 7, 3, and 1~mm fluxes.  As discussed below, this implies that at least some of the emission is optically thick.  There is some correlation between models with high values of $\\beta$ and those with high $\\tau$ and low $T$.  The temperature, $T_{\\rm 10~AU} = 26^{+34}_{-13}$~K, and $\\lambda$ = 3~mm optical depth, $\\tau_{\\rm 10~AU} = 0.50^{+0.83}_{-0.36}$, are not as tightly constrained as $r_{out}$. Although many models have $\\tau > 1$ at 10~AU, most (89\\%) radiate more than half their total emission in an optically thin regime.  Note that for an optically thin disk in the Rayleigh-Jeans regime, Eqn.\\ \\ref{disk_flux} becomes degenerate in $T$ and $\\tau$.  Consequently, temperature and optical depth are anti-correlated for $\\tau <$ 0.5 at 10~AU.  The data do not narrowly constrain values for $q$, $p$, $\\theta$ or $\\alpha$.  The ranges used for $p$ and $q$ were $p$=0.5--2.0 and $q$=0.4--0.75.  Steeper density profiles are slightly favored over shallow ones; the probability that $p$ is $\\ge$ 1.5 is 65\\%.  The lack of preferred values for $p$ and $q$ is largely due to the degeneracy of $p$ and $q$ for optically thin emission.  It may be possible to constrain the disk parameters further by including data from additional wavelengths.  Mid-infrared flux densities are often used to determine $q$ and $T_o$, for example.  We chose not to include these data, however, because the mid-infrared flux traces material within a few AU of the star, while the millimeter data is sensitive mainly to material tens of AU away.  The simple power-law relations assumed in the model may not be valid over such a large range of disk radii and physical conditions.  The disk model masses, weighted by the model probabilities, were binned to derive a median value log(M$_D$/M$_\\odot$) = ${-2.4}^{+0.7}_{-0.6}$ (Figure \\ref{prob}). The wide range in mass is due largely to the range of $\\tau$ values that fit the data.  For a given 3~mm flux, disks with higher mass correspond to those with higher $\\tau$ and $\\beta$.  The median value for the disk mass, M$_D$ = $4 \\times 10^{-3}$ M$_\\odot$, is toward the low end of disk masses typically derived for classical T~Tauri stars (e.g.\\ Beckwith et al.\\ 1990).  We note, however, a large dependence of mass estimates on values assumed for the other parameters.  Beckwith et al.\\ assumed $\\beta$ = 1 and an outer radius of 100~AU\\null.  If we consider only models with $\\beta \\ge 0.75$ and an outer radius $>$ 50~AU, the median mass increases to $10^{-2}$ M$_\\odot$, similar to the minimum mass solar nebula. ",
        "conclusions": "It has been conjectured that the formation of planetary systems arises from the collapse of protostellar clouds which rotate more slowly than clouds from which binaries form (Safronov \\& Ruzmaikina 1985).  This, in turn, raises the possibility that planetary systems may fail to form in the binary environment.  In contrast, our observations and modeling demonstrate that a substantial mass of material, with size like that of the solar system, can exist around a star in a binary system with separation not less than 100 AU\\null. This material constitutes a large reservoir available to planet-forming processes; the low value of the mass opacity index, $\\beta$ =0.53, further suggests that the formation of larger grains may already be underway as a first step toward planetesimal formation (Beckwith \\& Sargent 1991; Mannings \\& Emerson 1994; Koerner et al.\\ 1995).  Our observations rule out a greater circumstellar mass of dust around T Tau S---whether in an edge-on circumstellar disk or compact spherical envelope---as a simple explanation for the origin of increased extinction along the line of sight to T Tau's infrared companion. Due largely to the low optical depth of dust continuum emission at $\\lambda$ = 3mm, however, the true source of extra extinction is not identified unambiguously.  Our estimate of the outer radius is smaller than the projected separation between T~Tau~N and S and suggests that a disk around T Tau~N does not obscure T~Tau~S\\null.  However, we did not consider models with an exponential density profile at the outer edge like that of Hogerheidje et al.\\ (1997).  We note that since $\\kappa(500\\ {\\rm nm})/\\kappa(3\\ \\rm{mm}) \\sim 10^3$--$10^4$ (e.g., Pollack et al.\\ 1994), the material that provides the extinction toward T Tau S could easily be undetectable at 3 mm.  Thus, our data are consistent either with obscuration of T Tau S by tenuous outer regions of the T Tau N disk or with truncation of the T Tau N disk by T Tau S as discussed below.  If the binary components are in a bound orbit, as suggested by their common proper motion, the distribution of circumstellar material at the outer edge of the disk will be affected by the gravitational influence of the companion.  In models of disk/companion interactions, the size of the circumstellar disk depends on the mass ratio of the stars and their separation (Papaloizou \\& Pringle 1977, Artymowicz \\& Lubow 1994).  The radius of the circumprimary disk typically ranges from 0.3 to 0.4 times the separation for mass ratios of 1 to 0.3 and circular orbits.  If the orbital plane of the system is viewed nearly face-on and the binary separation is 100--110~AU, the tidal disk radius would be 30--40~AU\\null.  The disk size estimated from modeling our observations, 27 $< R <$ 67~AU, is consistent with the range of values predicted for tidal truncation.  If tidal truncation has occurred and we are seeing the true size of the disk in our maps, then the disk around T Tau N is not obscuring T Tau S\\null. However, high-resolution observations at sub-millimeter wavelengths or in molecular transitions that better trace low-column-density material are needed to adequately solve this problem.  Finally, we point out that the disk around T Tau N is similar to those observed around some single low-mass stars, regardless of the reliability of model assumptions that lead to an estimate of the absolute value of its mass and size, since the millimeter-wave flux from T Tauri is among the brightest measured from a large sample of young low-mass stars (cf.\\ Beckwith et al.\\ 1990, Osterloh \\& Beckwith 1995), and the nominal FWHM size of the emission is similar to that of single-star disks for which sufficiently high-resolution observations have been carried out (e.g. Lay et al.\\ 1994; Mundy et al.\\ 1996). If the largest disks around T Tauri stars typically yield planetary systems like our own, it is plausible that the disk around T Tau N will too, in spite of the presence of a companion at a distance of 100 AU or greater. Since the majority of pre-main-sequence stars are in multiple systems, the possibility of planetary formation in binaries like T Tauri invites consideration of the likelihood that a non-negligible fraction of binary stars contribute substantially to the estimated fraction of stars that possess planetary systems."
    },
    "9803/astro-ph9803170_arXiv.txt": {
        "abstract": "We use recent data obtained by three (OSSE, BATSE, and COMPTEL) of four instruments on board the Compton Gamma Ray Observatory, to construct a model of Cyg X-1 which describes its emission in a broad energy range from soft X-rays to MeV $\\gamma$-rays self-consistently. The $\\gamma$-ray emission is interpreted to be the result of Comptonization, bremsstrahlung, and positron annihilation in a hot optically thin and spatially extended region surrounding the whole accretion disk. For the X-ray emission a standard corona-disk model is applied. We show that the Cyg X-1 spectrum accumulated by the CGRO instruments during a $\\sim$4 year time period between 1991 and 1995, as well as the HEAO-3 $\\gamma_1$ and $\\gamma_2$ spectra can be well represented by our model.  The derived parameters match the observational results obtained from X-ray measurements. ",
        "introduction": "One of the brightest sources in the low-energy $\\gamma$-ray sky, Cyg X-1, has been extensively studied during the last three decades since its discovery (\\cite{Bowyer65}, for a review see \\cite{Oda77,LiangNolan84}). It is a high-mass binary system (HDE~226868) with an orbital period of 5.6~days consisting of a blue supergiant and presumably a black hole (BH) with a mass in excess of $5M_\\odot$ (\\cite{Dolan92}). The separation of the two components is $\\approx4\\times10^{12}$ cm (\\cite{Beall84}). A periodicity of 294~d found in X-ray and optical light curves is thought to be related to precession of the accretion disk (\\cite{Priedhorsky83,Kemp83}).  The X-ray flux of Cyg X-1 varies on all observed timescales down to a few milliseconds (e.g., \\cite{Cui97}), but the average flux exhibits roughly a two-modal behaviour.  Most of its time Cyg X-1 spends in a so-called `low' state where the soft X-ray luminosity (2--10 keV) is low. The low-state spectrum is hard and can be described by a power-law with a photon index of $\\sim1.7$ in the 10--150 keV energy band. There are occasional periods of `high' state emission, in which the spectrum consists of a relatively stable soft blackbody component and a weak and variable hard power-law component.  Remarkable is the anticorrelation between the soft and hard X-ray components (\\cite{LiangNolan84}), which is clearly seen during the transition phases between the two states.  Cyg X-1 is believed to be powered by accretion through an accretion disk. Its X-ray spectrum indicates the existence of a hot X-ray emitting and a cold reflecting gas. The soft blackbody component is thought to consist of thermal emission from an optically thick and cool accretion disk (\\cite{ShakuraSunyaev73,Pringle81,Balucinska95}).  The hard X-ray part ($\\ga10$ keV) with a break at $\\sim150$ keV has been attributed to thermal emission of the accreting matter Comptonized by a hot corona with temperature from tens to hundred keV (\\cite{SunyaevTitarchuk80,LiangNolan84}). A broad hump peaking at $\\sim20$ keV (\\cite{Done92}), an iron K$\\alpha$ emission line at $\\sim6.2$ keV with an equivalent width $\\sim100$ eV (\\cite{Barr85,Kitamoto90}, see also \\cite{Ebisawa96} and references therein), and a strong iron K-edge (e.g., see \\cite{Inoue89,Tanaka91,Ebisawa92},1996) have been interpreted as signatures of Compton reflection of hard X-rays off cold accreting material.  In addition, there have also been sporadic reports of a hard spectral component extending into the MeV region. The most famous one was the so-called `MeV bump' observed at a $5\\sigma$ level during the HEAO-3 mission (\\cite{Ling87}). For a discussion of the pre-CGRO data and $\\gamma$-ray emission mechanisms see, e.g., a review by Owens \\& McConnell (1992). The COMPTEL spectrum accumulated over 15 weeks of real observation time during the 1991--95 time period shows significant emission out to several MeV (\\cite{McConnell97}), which, however, remained always by more than an order of magnitude below the MeV bump reported from the HEAO-3 mission.  The annihilation line search provided only tentative (1.9$\\sigma$) evidence for a weak 511 keV line with a flux of $(4.4\\pm2.4)\\times10^{-4}$ photons cm$^{-2}$ s$^{-1}$ (\\cite{LingWheaton89}). Recent OSSE observations (\\cite{Phlips96}) resulted only in upper limits with values of $\\le7\\times10^{-5}$ cm$^{-2}$ s$^{-1}$ for a narrow 511 keV line and $\\le2\\times10^{-4}$ cm$^{-2}$ s$^{-1}$ for a broad feature at 511 keV.  Although an unified view for the X-ray spectra of BH candidates and their spectral states has yet to be constructed, the qualitative picture seems to be quite clear.  Current popular models include an optically thick disk component, a hot Comptonizing region (e.g., \\cite{Haardt93,Gierlinski97}), and/or an advection-dominated accretion flow (e.g., \\cite{Abramowicz95,NarayanYi95} and references therein). The spectral changes are probably governed by the mass accretion rate (e.g., \\cite{Chen95,Esin97}).  \\begin{deluxetable}{lc}% \\tablecolumns{2} \\footnotesize \\setlength{\\tabcolsep}{0.25em} \\tablecaption{ Luminosity of Cyg X-1.  \\label{table1}} \\tablewidth{7cm} \\tablehead{ \\colhead{Energy band} & \\colhead{Luminosity, $10^{36}$erg/s} } \\startdata $\\geq0.02$ MeV  & $26$ \\nl 0.02--0.2 MeV   & $20.5$ \\nl 0.2--1 MeV      & $4.8$\\nl $\\geq1$ MeV     & $0.6$\\nl \\enddata \\end{deluxetable}%  This picture, however, provides no explanation for the observed $\\gamma$-ray emission (e.g., McConnell et al. 1997). The hard MeV tail can not be explained by standard Compton models because they predict fluxes which are too small at MeV energies, and thus another mechanism is required. The models developed so far connect the $\\gamma$-ray emission with a compact hot core ($\\sim400$ keV or more) in the innermost part of the accretion disk, which emits via bremsstrahlung, Compton scattering, and annihilation (\\cite{LiangDermer88,SkiboDermer95}), or with $\\pi^0$ production due to collisions of ions with nearly virial temperature (e.g., \\cite{KolykhalovSunyaev79,JourdainRoques94}).  Li, Kusunose \\& Liang (1996) have shown that stochastic particle acceleration via wave-particle resonant interactions in plasmas ($\\sim100$ keV) around the BH could provide a suprathermal electron population, and is able to reproduce the hard state MeV tail.  The possibility of Comptonization in the relativistic gas inflow near the BH horizon has been discussed by Titarchuk \\& Zannias (1998).  We use the recent data obtained by three of four instruments aboard CGRO to construct a model of Cyg X-1, which describes its emission in a wide energy range from soft X-rays to MeV $\\gamma$-rays (\\cite{Moskalenko97}).  Instead of a compact (pair-dominated) $\\gamma$-ray emitting region, we consider an optically thin and spatially extended one surrounding the whole accretion disk. It produces $\\gamma$-rays via Comptonization, bremsstrahlung and positron annihilation. For the X-ray emission the corona-disk model is retained.  In section 2 we discuss the combined OSSE--BATSE--COMPTEL spectrum of Cyg X-1. Our model and the inferred results are described in sections 3--4, and the implications are discussed in section 5. The applied formalism is given in the Appendix. ",
        "conclusions": "The data obtained recently by the CGRO instruments allow us to construct a model of Cyg X-1 which describes its emission from soft X-rays to MeV $\\gamma$-rays self-consistently.  This model is based on the suggestion that the $\\gamma$-ray emitting region is a hot optically thin and spatially extended proton-dominated cloud, the outer corona. The emission mechanisms are bremsstrahlung, Comptonization, and positron annihilation.  For X-rays a standard corona-disk model is applied.  The CGRO spectrum of Cyg X-1 accumulated over a $\\sim$4 years period between 1991 and 1995, as well as the HEAO-3 $\\gamma_1$, and $\\gamma_2$ spectra can be well represented by our model.  The derived parameters match also the basic results of the X-ray observations.  A fine tuning of the model would require further Monte Carlo simulations and more accurate spectral measurements. In this respect, the solution of the discrepancy between the OSSE and BATSE normalizations would be of particular importance."
    },
    "9803/astro-ph9803200_arXiv.txt": {
        "abstract": "s{Recent observations of microlensing events in the Large Magellanic Cloud suggest that a sizable fraction of the galactic halo is in the form of Massive Astrophysical Compact Halo Objects (MACHOs). Although the average MACHO mass is presently poorly known, the value $\\sim 0.1 M_{\\odot}$ looks as a realistic estimate, thereby implying that brown dwarfs are a viable and natural candidate for MACHOs. We describe a scenario in which dark clusters of MACHOs and cold molecular clouds (mainly of $H_2$) naturally form in the halo at galactocentric distances larger than 10-20 kpc. Moreover, we discuss various experimental tests of this picture.}   \\normalsize\\baselineskip=15pt ",
        "introduction": "Since 1993 several microlensing events have been detected towards the Large Magellanic Cloud by the MACHO and EROS collaborations. Everybody agrees that this means that Massive Astrophysical Compact Halo Objects (MACHOs) have been discovered. Yet, the specific nature of MACHOs is unknown, mainly because their average mass turns out to depend strongly on the assumed galactic model. For instance, the spherical isothermal model would give $\\sim 0.5 M_{\\odot}$ whereas the maximal disk model would yield $\\sim 0.1 M_{\\odot}$ for that quantity. What can be reliably concluded today is only that MACHOs should lie in the mass range $0.05 M_{\\odot} - 1~M_{\\odot}$. Remarkably enough, the MACHO team has claimed that the fraction of galactic matter in the form of MACHOs is fairly model independent and -- within the present statistics -- should be $\\sim 50 \\%$.  What is the most realistic galactic model? Regretfully, no clear-cut answer is presently available. Nevertheless, the current wishdom -- that the Galaxy ought to be best described by the spherical isothermal model -- seems less convincing than before and nowadays various arguments strongly favour a nonstandard galactic halo. Indeed, besides the observational evidence that spiral galaxies generally have flattened halos, recent determinations of both the disk scale length, and the magnitude and slope of the rotation at the solar position indicate that our galaxy is best described by the maximal disk model. This conclusion is further strengthened by the microlensing results towards the galactic centre, which imply that the bulge is more massive than previously thought. Correspondingly, the halo plays a less dominant r\\^ole than within the spherical isothermal model, thereby reducing the halo microlensing rate as well as the average MACHO mass. A similar result occurs within the King-Michie halo models, which also take into account the finite escape velocity and the anisotropies in velocity space (typically arising during the phase of halo formation). Moreover, practically the same conclusions also hold  for flattened galactic models with a substantial degree of halo rotation. So, the expected average MACHO mass should be smaller than within the spherical isothermal model and the value $\\sim 0.1~M_{\\odot}$ looks as a realistic estimate to date. This fact is of paramount importance, since it implies that brown dwarfs are a viable and natural candidate for MACHOs. Still -- even if MACHOs are indeed brown dwarfs --  the problem remains to explain their formation, as well as the nature of the remaining dark matter in galactic halos.  We have proposed a scenario in which dark clusters of MACHOs and cold molecular clouds -- mainly of $H_2$ -- naturally form in the halo at galactocentric distances larger than $10-20$ kpc (somewhat similar ideas have also been put forward by  Ashman and by Gerhard and Silk). Below, we shall review the main features of this model, along with its observational implications. ",
        "conclusions": ""
    },
    "9803/astro-ph9803085_arXiv.txt": {
        "abstract": "We compute the polarization of the Ly$\\alpha$ line photons emerging from an anisotropically expanding and optically thick medium, which is expected to operate in many Ly$\\alpha$ emitting objects including the primeval galaxy DLA~2233+131 and Lyman break galaxies. In the case of a highly optically thick medium, the escape of resonance line photons is achieved by a large number of resonant local scatterings followed by a small number of scatterings in the damping wing. We show that some polarization can develop because the wing scatterings are coupled with strong spatial diffusion which depends on the scattering geometry and kinematics. The case of a slab with a finite scattering optical depth and expansion velocity of $\\sim 100~\\kms$ is investigated and it is found that Ly$\\alpha$ photons are emergent with the linear degree of polarization up to 10 per cent when the typical scattering optical depth $\\tau {\\gtrsim} 10^5$. We subsequently investigate the polarization of Ly$\\alpha$ photons emerging from a spherical shell obscured partially by an opaque component and we obtain $\\sim$ 5 per cent of polarization. It is proposed that a positive detection of polarized Ly$\\alpha$ with P-Cygni type profile from cosmological objects can be a strong test of the expanding shell structure obscured by a disk-like component. ",
        "introduction": "Various astronomical objects in the cosmological scales show P-Cygni type profiles in the Ly$\\alpha$ emission. These objects include the most remote galaxy at $z=4.92$ gravitationally lensed by CL1358+62 \\markcite{fra97} (Franx et al. 1997), high $z$ galaxies observed with the {\\it Hubble Space Telescope} and the Keck telescopes \\markcite{ste96a, ste96b, gia96, low97} (Steidel et al. 1996, Giavalisco et al. 1996, Lowenthal et al. 1997) and the damped Lyman $\\alpha$ (hereafter DLA) candidates \\markcite{djor96, djor97} (Djorgovski et al. 1996, 1997). Similar P-Cygni Ly$\\alpha$ profiles are found in nearby starburst galaxies, which are sometimes classified as Wolf-Rayet galaxies, blue compact galaxies, or H~II galaxies \\markcite{kun96, hec97, sah97, leq95, leg97} (e.g. Kunth et al. 1996, Heckman and Leitherer 1997, Sahu and Blades 1997, Lequeux et al. 1995, Legrand et al. 1997, etc.). The column density $N_{HI}$ of neutral hydrogen in these systems is usually found to be in the range $N_{HI}\\sim 10^{19-21}~\\cm^{-2}$.  The primeval galaxies or the first star clusters expected to be found at $z>5$ epoch may possess a central super star cluster surrounded by neutral hydrogen of high column density \\markcite{hai97a, hai97b} (Haiman and Loeb 1997a,b). These surrounding layers can be accelerated by the expanding H~II region just outside the super star cluster. It is hoped that in the near future with the advent of the {{\\it Next Generation Space Telescope} (NGST), the infrared spectra of these infant galaxies will be accessible and that the observational confirmation of the ubiquity of P-Cygni type Ly$\\alpha$ profiles may test the above hypothesis.   \\markcite{AL98} Ahn and Lee (1998, hereafter AL98) investigated the Ly$\\alpha$ line formation in a thick and expanding medium. It was emphasized that the profile formation should be studied by accurately computing the contributions from photons back scattered by receding medium and wing-scattered photons \\markcite{leg97} (see also Legrand et al. 1997). It is well known that the properties of the Ly$\\alpha$ photons scattered in the damping wing are characterized by the Rayleigh phase function \\markcite{ste80} (e.g. Stenflo 1980). This is in contrast with the degree of polarization $p=0$ resulting from a resonance transition between $1S_{1/2}$ and $2P_{1/2}$ and $p=3/7$ obtained for the $1S_{1/2}$ and $2P_{3/2}$ transition \\markcite{lee94b} (Lee et al. 1994).  In a moderately thick and static medium a negligibly polarized flux is expected because the photons are locally scattered many times and get isotropized before they escape to the observer \\markcite{lee94b} (e.g. Lee 1994). However, in a very thick medium, the escape is achieved by a large number of local resonant scatterings followed by a small number of scatterings in the damping wing. Hence, in the wing regime the spatial diffusion becomes important and the radiation field may get anisotropic depending on the scattering geometry. Therefore, the emergent line photons may get polarized and also anisotropic kinematics introduced in the medium can enhance the polarization.  In this {\\it Letter}, we compute the polarized flux of the emergent Ly$\\alpha$ from an optically thick and expanding slab. This result is applied to a hemi-spherical shell that is expected in various systems including primeval galaxies exhibiting P-Cygni profiles. ",
        "conclusions": "It seems a general consensus that the P-Cygni type Ly$\\alpha$ emissions are originated from expanding envelopes of H~II regions, which are indicative of the massive star formation. These are often obscured by dust lanes or thick molecular disks \\markcite{ich94, sco98} (Ichikawa et al. 1994, Scoville et al. 1998).  In this {\\it Letter} we computed the polarization of the Ly$\\alpha$ photons that are transferred through an optically thick and expanding neutral hydrogen layer. Anisotropic expansion and high column density are coupled to enhance scatterings in the damping wing into the direction corresponding to the largest velocity gradient, which results in highly polarized emergent flux. In particular, in a spherical shell with column density $\\sim 10^{20}~\\cm^{-2}$ and expansion velocity $\\sim 100~\\kms$ we find that the averaged degree of polarization of the emergent Ly$\\alpha$ line photons reach as high as 0.05 when $\\mu = 0.5$.  There are three interesting classes of primeval objects showing P-Cygni type profiles; the DLA candidate galaxies including DLA~2233+131 \\markcite{djor96, lu96, lu97} (Djorgovski et al. 1996, Lu et al. 1996, 1997) and DLA~2247-021 \\markcite{djor97} (Djorgovski 1997), the Lyman break galaxies at $3<z<4$ \\markcite{ste96, ste97, low97} (Steidel et al. 1996, 1997, Lowenthal et al. 1997), and the remote galaxies observed in the gravitational lens surveys \\markcite{fra97, fry97, tra97} (Franx et al. 1997, Frye et al. 1997, Trager et al. 1997).  Firstly, several DLA galaxies with $3<z< 4$ are listed including DLA~2233+131 with $z=3.15$ by \\markcite{djor96, djor97} Djorgovski et al. (1996,1997) who concluded that these objects are progenitors of normal disk galaxies today.  In view of the point that the galactic rotation would erase the P-Cygni structure, the Ly$\\alpha$ emitters or the giant H~II regions might be concentrated in a compact region. It is noted that they also exhibit the P-Cygni type Ly$\\alpha$ profiles, which is one of prominent characteristics of nearby dwarf starbursts \\markcite{meu95, kun96, leg97, lei97} (e.g. Meurer et al. 1995, Kunth et al. 1996, Legrand et al. 1997, Leitherer 1997). According to \\markcite{meu95} Meurer et al. (1995), there is evidence for existence of a dust lane or a thick H I disk component partly obscuring the UV sources or the super star clusters.  Secondly, interesting astronomical objects showing P-Cygni type profiles are the Lyman break galaxies at $z \\sim 3$ discovered by \\markcite{ste96} Steidel et al. (1996) \\markcite{low97} (see also Lowenthal et al. 1997).  From $\\sim 75\\%$ of the Lyman break galaxies obtained from their recent survey, \\markcite{ste97} Steidel et al. (1997) detected Ly$\\alpha$ emission lines which are weaker than expected from the UV continuum luminosities. They ascribed this to the resonant scattering \\markcite{kun98} (see Kunth et al. 1998).  The third group includes the remote galaxies which can be observed using the gravitational lens effect of the clusters of galaxies \\markcite{fra97, fry97, tra97} (Franx et al. 1997, Frye et al. 1997, Trager et al. 1997). Majority of the obtained profiles show the prominent P-Cygni type.  One other important application may be found in the first generation star clusters investigated by \\markcite{hai97a, hai97b} Haiman and Loeb (1997a,b). These primeval objects are expected to be observed at the redshift of $5<z<10$ and therefore the Ly$\\alpha$ emission will be located most possibly in the near IR band where the NGST can be used. The neutral hydrogen layers enveloping the super star clusters may trace the Hubble-type expansion, and the mechanism of the Ly$\\alpha$ line formation may be studied in a similar way introduced in this work \\markcite{AL98} (Ahn \\& Lee 1998).  We conclude that the Ly$\\alpha$ lines emerging from an expanding and optically thick medium are measurably polarized by up to 5 per cent and therefore a good constraint on theoretical models is provided by band polarimetry."
    },
    "9803/astro-ph9803198_arXiv.txt": {
        "abstract": "The {\\small CELESTE} experiment will be an Atmospheric Cherenkov detector designed to bridge the gap in energy sensitivity between current satellite and ground-based gamma-ray telescopes, 20 to 300 GeV. We present test results made at the former solar power plant, Themis, in the French Pyrenees. The tests confirm the viability of using a central tower heliostat array for Cherenkov wavefront sampling. ",
        "introduction": "The {\\small CELESTE} experiment uses the Electricit\\a'e de France central receiver solar power plant at Themis (N. 42.50$^\\circ$, E. 1.97$^\\circ$, 1650 m. a.s.l.) as a gamma-ray telescope \\cite{proposal}. The project is fully funded and will begin observations in early 1998 with 18 heliostats, growing to 40 in the following year. Figure \\ref{principle} sketches the principal of the {\\small CELESTE} approach. This paper describes test results accumulated during the design and construction phase. Emphasis is placed on Cherenkov measurements made using six heliostats between October 1996 and February 1997.  \\subsection{Science} Active Galactic Nuclei (AGNs), pulsars, and supernova remnants are complex ``cosmic accelerators''. High energy radiation dominates the power output of certain classes of these objects \\cite{vonM,Egpuls}. Study of their high energy spectra provides insight into the origin of cosmic rays, especially at the highest energies, and the nature of AGNs.  Photons from distant galaxies also probe the extragalactic medium: gamma-rays incident on near infrared photons are above threshold for $e^+e^-$ pair production and are thus absorbed, with approximately one attenuation length for GeV to TeV photons traversing cosmological distances. In this way gamma-ray spectra provide information on galaxy formation and, indirectly, on the nature of dark matter \\cite{ir}.  Around 1990 two breakthroughs revolutionized the field. First, ground-based atmospheric Cherenkov detectors became sensitive, reliable instruments above a few hundred GeV, led by the Whipple imager \\cite{whipple} and followed by the Themistocle and {\\small ASGAT} wavefront samplers at Themis \\cite{them,asgat}. Recent measurements by {\\small CAT} and other imagers of flaring in Mrk 501 underscore the rapid improvement of ground-based detectors \\cite{cat}. Second, the {\\small EGRET} instrument on the Compton satellite measured the spectra of over 150 point sources and mapped the galactic diffuse gamma-ray emission, in the energy range $0.1 < E_\\gamma < 10$ GeV \\cite{Egcat}. The energy range currently inaccessible either by satellite or ground-based detectors, $20 < E_{\\gamma} < 200$ GeV, is particularly rich with information on the acceleration and absorption processes. Bridging this energy gap is a key scientific goal.   \\subsection{Solar Arrays as Atmospheric Cherenkov Detectors} The Atmospheric Cherenkov Technique has been described extensively elsewhere: see for example \\cite{weekes,whipple,them,asgat,cat}. A few points bear repeating.  The minimum energy threshold of a Cherenkov telescope is limited by accidental trigger coincidences induced by photons from the night sky. For a constant diffuse night sky light $\\phi$ (photons per unit time, area, and solid angle), and a coincidence time gate $\\tau$ the threshold scales as \\begin{equation} \\label{thresh_form} E_{threshold} \\propto \\sqrt{ {\\Omega \\tau \\phi} \\over {A \\epsilon}} \\end{equation} where $\\Omega$ is the solid angle seen by a phototube and $\\epsilon$ is the photon collection efficiency. Current telescopes have pushed $\\tau$ and $\\Omega$ to their practical limits. Until technological progress improves quantum efficiencies and hence $\\epsilon$, lower energy thresholds require larger mirror areas, $A$.  Solar heliostat arrays offer the advantage of a large mirror area available without additional construction costs. The optical configuration is similar to that of multiple mirror wavefront samplers such as Themistocle. The main difference is the common focus of all heliostats is situated at the top of the tower. Secondary optics disentangle the light coming from  each of the heliostats, as sketched in figure \\ref{principle}. The {\\small STACEE} collaboration is building a similar detector in the United States \\cite{stacee}.  Imagers are superior to samplers above 200 GeV, since the fine sampling of their cameras allows excellent hadron rejection. But below 100 GeV, hadron backgrounds are naturally suppressed because the Cherenkov yield of the hadron showers decreases faster than that of the gamma showers. Hence, the advantage of the imagers in this regard is less.   At low energies, the background due to cosmic ray electrons is greater than above 200 GeV due to an energy spectrum steeper than for hadrons. Since electron induced showers are practically indistinguishable from gamma-ray showers, angular resolution is the key to suppressing the electron background. But angular resolution, especially at lower energies, tends to be limited by shower development rather than by instrument response. Scattering in the first few generations of the shower and deflection of low energy secondaries in the geomagnetic field are the factors limiting angular precision. Hence, a high-granularity imager does not necessarily outperform a wavefront sampler. Both should achieve angular resolutions around $0.1^\\circ$.  Sections 2 and 3 describe the apparatus and the main results of the prototype tests. Based on these results we   compare  sampling  arrays  and imagers for the low energy domain in section 4. ",
        "conclusions": "We have used a solar array to detect Cherenkov light from air showers. We accurately measure the shower light emitted in the small volume at the intersection of the fields-of-view of two groups of three heliostats each.  Triggering on widely spaced heliostat groups has been shown to give good hadron and muon rejection.  The minimum trigger threshold at which the Cherenkov signal dominated the night sky background was 5 photoelectrons per heliostat. In our analysis the threshold was 7 photoelectrons, which corresponds to a hadron shower energy slightly above 300 GeV, or to a gamma-ray energy of 80 GeV. At this threshold the trigger rate was 6 Hz. We expect a threshold of 2 -- 3 photoelectrons and a rate of 20 Hz with the optics now being installed, corresponding to a gamma-ray energy threshold below 30 GeV.  Our test results, and in particular the low trigger rate, confirm that wavefront sampling is an alternative to imaging as a basis for a high-performance gamma-ray detector in the 20 to 200 GeV energy range. We are presently commissioning a full-scale device which we expect to provide the first ground-based detection of cosmic gamma-rays in this energy range. The ongoing observations could establish that, even independently of existing solar arrays, wavefront sampling might be the best way to access an  unexplored  spectral  window  known  to  be  particularly  rich   in astrophysical information. \\\\  {\\bf ACKNOWLEDGEMENTS}\\\\ We gratefully acknowledge the contributions of the technical staffs of our laboratories and the IMP (CNRS) at Odeillo. Drs. Philippe Roy and Louis Behr play key roles in the continued progress of the project. We thank Electricit\\a'e de France for allowing use of the Themis site. Funding was provided by the Institut National de Physique des Particules et de Physique Nucl\\a'eaire of the Centre National de Recherche Scientifique; the Grant Agency of the Czech Republic; by the University of Bordeaux; by the Ecole Polytechnique; and by the Regional Councils of Aquitaine and of Languedoc-Roussillon."
    },
    "9803/astro-ph9803017_arXiv.txt": {
        "abstract": "Multi-wavelength imaging and spectroscopy of the $z=0.708$ radio galaxy 3C441 and a red aligned optical/infrared component are used to show that the most striking aspect of the radio-optical ``alignment effect'' in this object is due to the interaction of the radio jet with a companion galaxy in the same group or cluster. The stellar population of the red aligned continuum component is predominately old, but with a small post-starburst population superposed, and it is surrounded by a low surface-brightness halo, possibly a face-on spiral disc. The [O{\\sc iii}]500.7/[O{\\sc ii}]372.7 emission line ratio changes dramatically from one side of the component to the other, with the low-ionisation material apparently having passed through the bow shock of the radio source and been compressed. A simple model for the interaction is used to explain the velocity shifts in the emission line gas, and to predict that the ISM of the interacting galaxy is likely to escape once the radio source bow shock has passed though. We also discuss another, much fainter, aligned component, and the sub-arcsecond scale alignment of the radio source host galaxy. Finally we comment on the implications of our explanation of 3C441 for theories of the alignment effect. ",
        "introduction": "The consequences of radio jets impacting on density inhomogeneities have been invoked to explain many properties of high redshift radio sources, such as bending and asymmetries in arm length. Evidence for this is seen in the form of correlations between emission-line gas and the side of the radio galaxy with the shorter arm-length, or higher depolarisation (McCarthy, van Breugel \\& Kapahi 1991; Liu \\& Pooley 1991). The extent to which this is related to the so-called ``alignment effect'' whereby the axis of extended optical continuum and line emission is found to be co-aligned with the radio axis, is unclear, although many attempts to explain radio-optical alignments involve a dense external medium. Such a medium is needed, for example, to act as the source of a scattering surface for quasar light from the AGN (Bremer, Fabian \\& Crawford 1997), or as a medium for jet-induced star formation (e.g.\\ Best, Longair \\& R\\\"{o}ttgering 1996 \\& references therein).  3C441 is a $z=0.708$ radio galaxy with an asymmetric radio structure, which appears to be a rare example of a radio source with a red aligned component outside the radio lobes. Such a component is clearly hard to obtain in either jet-induced star formation or scattered quasar models. This red component is seen just beyond the end of the shorter, north-west radio lobe [component `c' of Eisenhardt \\& Chokshi (1990); see Fig.\\ 1]. This lobe appears to possess a radio jet which bends round to the west at the tip of the lobe. Just to the south of the red component, mostly within the radio lobe, is a arc-shaped clump of emission line gas seen in the [O{\\sc ii}]372.7 emission line image of McCarthy, van Breugel \\& Spinrad (1994). Spectroscopy of the [O{\\sc ii}] emission line by McCarthy, Baum \\& Spinrad (1996) shows that this emission line material is redshifted by $\\approx 800\\;$km s$^{-1}$ with respect to the radio galaxy, and weak emission extends to just beyond the red component.  In this paper, we use our own spectroscopy, infrared and radio imaging, and archive data from the {\\em Hubble Space Telescope} {\\em (HST)} (also presented in Best, Longair \\& R\\\"{o}ttgering 1997c) to attempt to explain the properties of 3C441 in terms of models for the interaction of radio jets with their environments. In particular, we address the problems of the origin of the red continuum light from the aligned component and the properties of the extended emission line region.  We assume an Einstein -- de Sitter cosmology with a Hubble constant $H_0=50\\;$km s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\setlength{\\unitlength}{1mm} \\begin{figure} \\begin{picture}(75,75) \\put(0,-20){\\special{psfile=3c441csed.ps hscale=40 vscale=40}} \\end{picture} \\caption{\\small{The SED of the compact red knot `c'. The continuous line is from our own spectroscopy, the near infrared points from Eisenhardt \\& Chokshi (1990) (open squares) and our own near-infrared and optical photometry (solid triangles; note the $J$ and $K$ points have no associated error bars as the data were non-photometric). The dotted line is the model spectral energy distribution of a 6-Gyr old galaxy in which star formation proceeded at a constant rate in a 1-Gyr burst before ceasing (Bruzual \\& Charlot 1993).}} \\end{figure}  \\subsection{The continuum knot}   Component `c' in Fig.\\ 1a and in Eisenhardt \\& Chokshi (1990), the red aligned knot beyond the northwest radio hotspot, is slightly resolved on the {\\em HST} images. Its spectral energy distribution (SED) seems to be consistent with that of an old stellar population (Fig.\\ 3), although the amplitude of the 4000\\AA$\\;$break seems lower than that of the model 6-Gyr old stellar population (corresponding approximately to the age of the Universe in our assumed cosmology) which is otherwise a good fit to the SED.  A stellar origin for the light is indeed confirmed by a close inspection of the spectrum of `c' (Fig.\\ 4) in which stellar absorption features, in particular the ``G-band'' from CH absorption in cool stellar envelopes can be seen at $z=0.714$, redshifted by $1000\\, {\\rm km\\, s^{-1}}$ with respect to the radio galaxy ($z=0.708$). Also visible are H$\\delta$ and the calcium H and K lines, the latter possibly blended with H$\\epsilon$. The weakness of the 4000\\AA$\\;$break, despite the apparent dominance of the SED by an old stellar population may then be explained by the existence of a small post-starburst population of A-stars, and indeed the detection of H$\\delta$ absorption in our spectrum is consistent with this. Such populations have been detected in other AGN host galaxies and companions (Tadhunter, Dickson \\& Shaw 1996; Canalizo \\& Stockton 1997), and may trace mergers or interactions in small clusters or groups (Zabludoff et al.\\ 1996).  Whether there is a link between a starburst $\\sim 10^8$yr ago in a close companion or the host itself and the triggering of the AGN is unclear -- certainly if so there must be a delay between the starburst and the development of a radio source if typical radio source lifetimes are $\\sim 10^7$yr (Alexander \\& Leahy 1987). In the case of 3C441, component `c' is $\\approx 125$ kpc from the host galaxy, so it is conceivable that if its relative velocity with respect to the host in the plane of the sky is of the same order as its relative velocity along the line of sight, it could have been close enough to interact $10^8$yr ago. There is, however, no sign of a corresponding starburst in the host galaxy, and E+A galaxies are relatively common in high-$z$ clusters (Dressler \\& Gunn 1990; Caldwell \\& Rose 1997). Furthermore, the high velocity encounter implied by this may not have have produced sufficient disruption to initiate the starburst.    \\setlength{\\unitlength}{1mm} \\begin{figure} \\begin{picture}(75,75) \\put(0,-20){\\special{psfile=spectrum.ps hscale=40 vscale=40}} \\end{picture} \\caption{\\small{The red-arm spectrum of knot `c' of 3C441, showing the stellar absorption features mentioned in the text. The raw spectrum has been smoothed by a 9-pixel (2.4 nm) box-car filter and corrected for atmospheric absorption.}} \\end{figure}   \\subsection{The emission line gas in `c'}  Contour plots of the 2-D spectra around the [O{\\sc ii}]372.7 and [O{\\sc iii}]500.7 emission lines are shown in Fig.\\ 5. There is a striking change in the ionisation and luminosity across the position of the continuum knot: the side nearest to the radio galaxy has a high luminosity, and a low ionisation as measured by the ratio of [O{\\sc iii}]500.7/[O{\\sc ii}]372.7 emission lines ($0.35 \\pm 0.1$). In contrast, beyond the knot the [O{\\sc iii}]500.7/[O{\\sc ii}]372.7 ratio is much higher ($>10\\pm 3$) and the total luminosity lower by a factor of $\\approx 6$. There is a velocity gradient of $\\approx 300$ km s$^{-1}$ in the line emission across the continuum knot, in the sense that the side nearest the radio galaxy is blueshifted, and that on the far side is, within errors of $\\approx 100\\, {\\rm km\\, s^{-1}}$, at rest with respect to the starlight in `c'. There is also a faint tail of yet more highly blueshifted emission (up to $\\approx 600\\, {\\rm km\\, s^{-1}}$) on the side nearest the radio galaxy, pointing towards the radio galaxy.  Note that there is no sign of emission lines in Fig.\\ 4, which used a narrow extraction about the continuum peak. In contrast, in Fig.\\ 5 the extended emission-line flux from `c' is clearly visible and is quite strong when integrated over the entire emission region (Table 1).  \\setlength{\\unitlength}{1mm} \\begin{figure*} \\begin{picture}(150,60) \\put(-25,-200){\\special{psfile=linefig.ps}} \\end{picture} \\caption{\\small{2-D spectrum of 3C441 with the slit aligned along the axis joining the radio galaxy and `c'. Left: the [O{\\sc ii}] emission line with wavelength on the horizontal axis and distance along the slit vertically. The radio galaxy is to the top and component `c' below it. Right: the [O{\\sc iii}]5007 emission line. The figures are $11.2\\, {\\rm nm} \\times 30\\, {\\rm arcsec}$ in size (11.2 nm $\\approx$ 5280 km s$^{-1}$ close to [O{\\sc ii}] and $\\approx$ 3930 km s$^{-1}$ close to [O{\\sc iii}]). The contour levels are spaced at intervals of $2\\times 10^{-21} {\\rm Wm^{2}nm^{-1}}$ per pixel, starting from $2\\times 10^{-21} {\\rm Wm^{2}nm^{-1}}$, each pixel was $0.28 {\\rm nm} \\times 0.33 {\\rm arcsec}$ in size.}} \\end{figure*}   \\subsection{The radio structure}  The asymmetric radio structure of 3C441 is interesting in the context of the models to explain the aligned emission. The asymmetry in jet brightness either side of the nucleus is very pronounced, and can be interpreted either as an asymmetry produced by Doppler boosting, or in terms of 3C441 having a radio structure transitional between FRI and FRII, with one side FRII-like and the other more FRI-like. The lack of a radio central component, commonly seen in radio galaxies with Doppler boosted one-sided jets (e.g. 3C22, Rawlings et al.\\ 1994), argues strongly in favour of the latter explanation, and indeed the flaring of the jet just to the NW of the host galaxy is reminiscent of the ``Mach disk'' structure seen in the M87 jet (Owen, Hardee \\& Cornwell 1989).  \\subsection{The radio galaxy}  The spectrum of `a', the host galaxy of the radio source is shown in Fig.\\ 2. It is a typical moderately-high ionisation narrow-line radio galaxy spectrum, with strong emission lines superposed on a stellar spectrum dominated by an old stellar population with a strong 4000\\AA$\\;$break.  Close inspection of the images shows, however, that even in this case, where the integrated light is dominated by old stellar populations there is evidence for morphological peculiarity. In particular there is a distinct blue aligned component to the south east of the host galaxy peak (Fig.\\ 6). Whether this represents a spiral arm, a merger remnant or some form of radio source-induced aligned component is unclear. As its separation from the radio galaxy is only about 0.5-arcsec, it is hard to tell from the spectra whether it is line or continuum dominated, but the more diffuse material extending $\\approx 2$-arcsec to the south of the radio galaxy is definitely continuum dominated. Although much weaker relative to the smooth underlying host galaxy in the F785LP image, the aligned component is nevertheless visible, along with some more diffuse aligned emission on the other side of the peak, to the northwest.  In Fig.\\ 7, the position angle of the host galaxy (derived from the second moments of the flux distribution) is plotted for various isophotes and apertures. This shows two interesting aspects of the alignment. First, the alignment persists into the $K$-band, suggesting the aligned light is fairly red. Second, in the {\\em HST} images, there is evidence for isophotal twisting as the aperture size is increased to include the inner aligned components highlighted in Fig.\\ 6 (PA $\\sim 150$ deg), and later the low surface brightness emission round the host at PA $\\sim 0$, seen best in Fig.\\ 1. The low-surface brightness material to the northwest may be mostly line emission though, as the emission-line image of McCarthy et al.\\ (1995) shows that the [O{\\sc ii}] emission is also roughly aligned along PA 0.  \\setlength{\\unitlength}{1pt} \\begin{figure} \\begin{picture}(200,100) \\put(-240,-350) {\\special{psfile=3c441rg_555.ps}} \\put(-120,-350){\\special{psfile=3c441rg_785.ps}} \\put(0,80){(a)} \\put(120,82){(b)} \\end{picture} \\caption{Close-up of the host galaxy (`a') greyscaled so as to show the aligned components near the nucleus: (a) F555W image; (b) F785LP image. Both images are 10 arcsec square and have been smoothed with a $\\sigma = 0.05$ arcsec gaussian.} \\end{figure}   \\setlength{\\unitlength}{1pt} \\begin{figure*} \\begin{picture}(500,200) \\put(-165,-250){\\special{psfile=kbandten_bm.ps vscale=80 hscale=80}} \\put(15,-250){\\special{psfile=ibandten_bm.ps vscale=80 hscale=80}} \\put(195,-250){\\special{psfile=vbandten_bm.ps vscale=80 hscale=80}} \\end{picture} \\caption{Position angle as a function of aperture size and isophotal cutoff for the host galaxy of 3C441. The isophotal cutoff level is in units of the sky noise, and the aperture radii (indicated by circles of increasing size) are 0.6, 0.9, 1.1, 1.7 and  2.4 arcsec in (a), and 0.3, 0.4, 0.6, 0.8, 1.2, 1.7, and 2.4 arcsec in (b) and (c). The radio source PA (defined as the PA of the line joining the hotspots) is indicated by a dotted line.} \\end{figure*}   \\subsection{The relationship of the radio and optical components}  The relative astrometry was performed using the results of the APM scans of the Palomar Sky Survey plates. The positions of four stars were used to align the corners of the radio map with the optical images. The uncertainty in the overlay from the scatter in the fit was $\\approx 0.3$ arcsec. This astrometry places the host galaxy (`a' in Fig.\\ 1a) just to the SE of the first appearance of the northern radio jet. Component `d' is apparently aligned along the jet direction, although positioned to the side of it. The point of deflection of the jet is approximately coincident with the peak of the [O{\\sc ii}]372.7 emission, about 3-arcsec SE of the peak of the continuum knot `c'. In the F555W image (Fig.\\ 1a), there is an arc of emission just to the north of the north-west radio hotspot and apparently centered on `c'. Fig.\\ 1b shows the F785LP image; here there is a halo of diffuse emission around `c'.  By subtracting a spectrum centered on the continuum knot in an aperture 1.7-arcsec wide from the total spectrum of the aligned component `c' (in a 6.8-arcsec wide aperture), we have been able to estimate the line contribution to the continuum magnitudes measured for the nebular region surrounding `c'. In both filters this is $\\approx 10$ per cent overall, but in the regions of brightest line emission in the interaction region to the south of `c' our spectrum suggests that the line contribution of [O{\\sc ii}] to the F555W flux rises to dominate the overall flux in the filter, consistent with the arc of emission seen in the F555W region just above the radio hotspot consisting entirely of line emission. The continuum emission present in this extraction is bluer than the emission from the knot.  The linear object `d' (Fig.\\ 8) is reminiscent of the aligned component in 3C34 (Best, Longair \\& R\\\"{o}ttgering 1997a). Like the object in 3C34, it has no emission lines visible either in our spectra or in the narrow-band image of McCarthy et al.\\ (1994), but is well aligned with the radio structure and lies within the radio lobe. Its optical--near infrared colours are bluer than `c' ($K=20.8; J>22.7; m_{785}=23.3; m_{555}=25.1$ in 3-arcsec diameter apertures, where $m_{785}$ and $m_{555}$ are AB magnitudes in the two {\\em HST} images), but show a break between the 785LP and 555W filters, consistent with a 4000$\\;$\\AA$\\;$break at the redshift of the radio source.  \\setlength{\\unitlength}{1pt} \\begin{figure} \\begin{picture}(200,100) \\put(-548,-740){\\special{psfile=3c441d_555.ps hscale=200 vscale=200}} \\put(-430,-740){\\special{psfile=3c441d_785.ps hscale=200 vscale=200}} \\put(0,80){(a)} \\put(120,80){(b)} \\end{picture} \\caption{Close-up of component `d': (a) F555W image; (b) F785LP image. Both images are 5 arcsec square and have been smoothed with a $\\sigma = 0.1$ arcsec gaussian.} \\end{figure}   \\subsection{Other companion objects}  There are several objects around the radio galaxy which appear to be members of a group or cluster around it. 3C441 is in a sample of radio galaxies whose clustering properties we are currently evaluating, and a formal estimate of the clustering amplitude, $B_{\\rm gq}$, for this radio source will be presented in Wold et al.\\ (in preparation). For the purposes of this paper, we simply compared the counts of galaxies with magnitudes $m_1$ to $m_1+3$ (where $m_1$ is the magnitude of the radio source host galaxy) in the object frame of the F785LP image (the WF3 CCD) with the average of those in the two side frames (WF2 and WF4). This revealed an excess of $11.5 \\pm 7.1$ galaxies, indicating the possible presence of a cluster, but not at a high confidence level. Clearly though this is likely to be an underestimate of the cluster richness as many cluster members may be outside the restricted field of the CCD, and present on the other frames, increasing the estimate of the background count."
    },
    "9803/astro-ph9803221_arXiv.txt": {
        "abstract": "We modified the Press-Schechter (PS) formalism and then analytically derived a constrained mass distribution function $n(M|\\varphi)$ for the regions having some specified value of the primordial gravitational potential, $\\varphi$. The resulting modified PS theory predicts that gravitationally bound clumps with masses corresponding to rich clusters are significantly biased toward the regions of negative primordial potential - the troughs of the potential. The prediction is quantitative, depending on the mass and the depth of the troughs, which can be tested in large N-body simulations. As an illustration of the magnitude of the effect we calculate the constrained mass function for the CDM model with $\\Gamma = \\Omega h = 0.25$ normalized to $\\sigma_{8} = 1$. In particular, we show that the probability of finding a clump of mass $10^{14} - 10^{15}h^{-1}M_{\\odot}$ in the region of negative initial potential is $1.3 - 3$ times greater (depending on the mass) than that in the region of positive initial potential. The scale of the potential fluctuations $R_{\\varphi}=\\sqrt{3} \\sigma_{\\varphi}/\\sigma_{\\varphi'}$ is shown to be $\\approx 120 h^{-1}{\\rm Mpc}$ for the spectrum in question. The rms mass density contrast on this scale is only about $\\sigma _{\\delta}(R_{\\varphi}) \\approx 0.03$. Assuming that the modified PS theory is statistically correct, we conclude that clusters are significantly biased ($b \\ge 10$, $b$ is a bias factor defined by $\\Delta n_{cl}/ n_{cl} =b \\Delta \\rho_m/\\rho_m$) toward the regions having negative initial potential. ",
        "introduction": "Assuming the standard hierarchical model of the structure formation from Gaussian fluctuations due to gravitational instability, we study the effect of primordial gravitational potential fluctuations on massive objects such as galaxy clusters and perhaps superclusters, i.e., clusters of clusters (\\cite{bah-son84}). We employ the PS formalism as a tool and modify it for this study.  Some effect of the primordial gravitational potential upon the structure formation has been already noted. \\cite{kof-sha88} have noticed that the adhesion approximation predicts that the formation of voids is associated with positive peaks of the primordial gravitational potential. Sahni, Sathyaprakash, \\& Shandarin (1994) studied the effect  and measured a significant correlation between the sizes of voids and  the value of primordial gravitational potential in numerical simulations of the adhesion model. By investigating the evolution of correlation between the potential and the density perturbations, Buryak, Demianski, $\\&$ Doroshkevich (1992) showed that the formation of super large scale structures is mainly determined by the spatial distribution of the gravitational potential. Recently, Madsen et al. (1997) have demonstrated by N-body simulations that the under dense and the over dense regions are closely linked to the regions with the positive and the negative gravitational potential respectively. Thus, given all these results showing the important role of the primordial gravitational potential in the structure formation, it would be interesting to calculate the effect of the primordial potential upon the mass distribution function.  The mass distribution function $n(M)$ is defined such that $n(M)dM$ is the comoving number density of gravitationally bound objects in the mass range $(M,M + dM)$. The standard Press-Schechter (hereafter, PS) formalism provides an effective tool to evaluate $n(M)$ in spite of various criticism on it (see \\cite{mon98}), and is widely used in cosmology (e.g. \\cite{gro-etal97}; \\cite{kit-sut97}; \\cite{bah-fan98}; \\cite{rob-gaw-sil98}; \\cite{wan-ste98}). Also, \\cite{lee-sh98} have shown by applying the dynamics based on the Zel'dovich approximation to the PS formalism that it is very robust with respect to the underlying dynamics.  The following two equations represent the essence of the PS formalism (\\cite{pre-sch74}): \\begin{equation} n(M) = \\frac{\\bar{\\rho}}{M}\\bigg{|} \\frac{dF}{dM}\\bigg{|} , \\end{equation} \\begin{equation} F(M) = \\int^{\\infty}_{\\delta_{c}}\\! p(\\delta) d\\delta. \\end{equation} Here $p(\\delta)$ is the probability density distribution of the linearly extrapolated density contrast $\\delta$ smoothed on a comoving filtering scale $R$ which is related to the mass by $M=M(R)=\\alpha{\\bar\\rho}R^3$. The proportionality constant $\\alpha$ is either determined by the shape of the smoothing window function or sometimes is used as a free parameter in order to provide a better agreement with numerical results. In the case of a sharp k-space filter which is actually consistent with the PS formalism (see \\cite{pea-hea90}), the filtering scale $k_c = 2\\pi/R$ in k-space and mass are related as $M =  6\\pi^{2}\\bar{\\rho}k_c^{-3}$. The density threshold value $\\delta_c$ for collapse was originally given as $\\delta_{c}\\approx 1.69$ according to the Top Hat spherical model. However, it has been shown that the lowered value of $\\delta_c$ in the range from $1.3-1.6$ gives a better fit in N-body simulations, which depends on the the initial spectrum and the type of the filter (e.g.,  \\cite{gro-etal97}).  In this Letter we investigate and show how much the primordial gravitational potential $\\varphi$ affects the mass distribution function of galaxy cluster. Modifying the PS formalism, we derive  a constrained mass distribution function $n(M|\\varphi$) defined as the comoving number densities of clumps of mass $M$ in the regions where the primordial gravitational potential fluctuation satisfies some specified conditions. The Cold Dark Matter model (CDM) with $\\Gamma = \\Omega  h = 0.25$ and $\\sigma_8=1$ is used to demonstrate the significance of the effect. ",
        "conclusions": "The PS formalism has been proved to be a simple but very effective tool widely used for constraining cosmological models. We have modified it by considering the dependence of mass function on the initial perturbation of gravitational potential. The resulting modified PS theory predicts that the clumps with masses greater than roughly $10^{14}h^{-1}M_{\\odot}$ have a noticeable tendency to form in the troughs of the primordial gravitational potential (the regions where the primordial potential fluctuations were negative). This quantitative prediction can be tested in large N-body simulations. Regardless of the outcome it will shed light on the PS formalism; if our prediction is confirmed, it will show a new potency of the PS technique. Otherwise a new limitation to the formalism will be established.  Assuming that the prediction is correct at least qualitatively, \\footnote{N-body simulations (e.g., Madsen et al. 1997) and the adhesion model (Sahni et al. 1994) have already visually demonstrated this bias effect of the gravitational potential.} we would like to discuss some of its obvious consequences. The scale of the initial potential \\begin{equation} R_{\\varphi} = \\sqrt{3} \\sigma_{\\varphi}/ \\sigma_{\\varphi'} =\\sqrt{3\\frac{\\int^{\\infty}_{k_{l}}\\! dk k^{-2}P(k)} {\\int^{\\infty}_{0}\\! dk P(k)}} \\approx 120 h^{-1} {\\rm Mpc} \\end{equation} does not depend on any ad hoc scale; the dependence on $k_l$ is exremely weak ($\\propto \\sqrt{ln{(1/k_l)}}$ for the Harrison-Zel'dovich spectra assumed here). It is, perhaps, worth mentioning that the scale of the potential is also practically independent of the smoothing scale unless it exceeds the value of a few tens of $h^{-1}{\\rm Mpc}$. The density scale $R_{\\delta_{k_c}}$ is determined by the scale of the smoothing window function $k_c$ that has only one ``natural'' scale corresponding nonlinearity $k_c=k_{nl}$. For the model in question the scale of the primordial potential is found to be $R_{\\varphi} \\approx 120 h^{-1} {\\rm Mpc}$. The scale of the density contrast field reaches this value $R_{\\delta} = \\sqrt{3} \\sigma_{\\delta}/\\sigma_{\\delta'} \\approx 120 h^{-1} {\\rm Mpc}$ only after it is smoothed  on $k_c \\approx 0.017 h {\\rm Mpc^{-1}}$. The corresponding density variance on this scale is $\\sigma_{\\delta}(0.017 h {\\rm Mpc^{-1}}) \\approx 0.03$. On the other hand, the number of clumps with masses $10^{14} - 10^{15} h^{-1} M_{\\odot}$ can easily be  30\\% greater in the troughs of the potential than the mean density $n(>M) = 0.5[n(>M|\\varphi<0)+ n(>M|\\varphi>0)]$ (see Fig. 1). Thus, the bias factor $b$ (defined by the relation $\\Delta n_{cl}/ n_{cl} = b \\Delta \\rho_m/\\rho_m$) reaches at least $10$ on the scale about $120 h^{-1} {\\rm Mpc}$.  Qualitatively the bias phenomenon  can be explained as follows. The initial density contrast is proportional to the Laplacian of the initial potential ($\\delta \\propto  \\nabla^2 \\varphi$). Therefore the two fields are cross-correlated: the positive peaks of $\\delta$ are more likely to be found in the troughs of the potential where it is negative. The correlation is not very strong (for $k_c = 0.25 h {\\rm Mpc^{-1}}$ corresponding to $\\sigma_{\\delta} = 1$ the crosscorrelation coefficient $\\kappa = \\sigma_v^2(0.25 h {\\rm Mpc^{-1}})/ \\sigma_{\\varphi} \\sigma_{\\delta}(0.25 h {\\rm Mpc^{-1}}) \\approx 0.12$).  But the clusters are extreme objects corresponding to the tail of the mass function, and thus very sensitive to the environment. That is why the clusters put one of the strongest constraints on cosmological models (\\cite{kly-rhe94}, \\cite{bo-my96}, \\cite{fan-etal97}, \\cite{bah-fan98}).  Incorporating the motion of mass into dynamics can only increase the bias effect due to the nonlinear effects although they are quite small on the scale in question. But, the point is not in the magnitude of the nonlinear effects but rather in their sign. On the scale of the potential the mass moves from the peaks of the potential to the troughs. Using the Zel'dovich approximation one can easily estimate the rms displacement of the mass  on the scale of the potential (\\cite{sh93}): \\begin{equation} d_{rms} = \\sqrt{{{\\int_0^{0.017h}P(k) dk} \\over {\\int_0^{0.25h}P(k)k^2 dk}}} \\approx 3 h^{-1} {\\rm Mpc}. \\end{equation} It is relatively small compared to the scale of the potential but coherent on the scale of the potential field, and therefore it can only enhance the bias effect. Another nonlinear effect is related to the rate of growth of perturbations. For the perturbations on the scale of a few Mpc the potential troughs/peaks may be viewed as patches with slower/faster expansion rate that corresponds to the increase/decrease of the rate of growth of small-scale perturbations. Similarly, the bias is enhanced in the redshift space because the velocity field is directed toward the troughs and away from peaks of the potential. Both effects can increase the bias by about 5\\% depending on the initial spectrum.  Another  way of calculating the constrained mass function would be using the peak-background split technique suggested by \\cite{kai84} to explain the enhanced correlation function of reach clusters. Obviously, the initial potential resembles the smoothed initial density field if the filter has a sufficiently large scale, but the former is never identical to the latter. The potential itself can be viewed as a smoothed density field with a very soft scale-free filter $W(k) \\propto k^{-2}$. Typically the density field is filtered with much harder filters (e.g. top-hat, Gaussian, or sharp $k$-space filters), that impose the scale which is an ad hoc parameter. The magnitude of the bias in our approach is determined by the crosscorrelation of the density contrast smoothed at the scale ($k_c$) of nonlinearity ($\\sigma_{\\delta_{k_c}}=1$) with the initial potential that does not have any ad hoc parameters. Probably, the value of the crosscorrelation coefficient determines the bias in the peak-split approach as well. The crosscorrelation of the density field $\\delta_{k_c}$ smoothed on the scale of nonlinearity ($k_c = 0.25 h {\\rm Mpc^{-1}}$) with the field $\\delta_{k_{\\varphi}}$ smoothed on the scale of the potential ($k_{\\varphi} = 0.017 h {\\rm Mpc^{-1}}$) is about $4$ times weaker than the correlation of $\\delta_{k_c}$ with the initial potential $\\varphi$. Thus, we expect that the bias of galaxy clusters on such large scales as the scale of the initial potential ($\\approx 120 h^{-1} {\\rm Mpc}$) is stronger toward the troughs of the potential fluctuations rather than to the peaks of the density fluctuations $\\delta_{k_c}$ smoothed with the corresponding filter. We have not applied the split peak-background approach because it is not clear how to avoid arbitrarines in choosing the scale that splits the density into small-scale peaks and large-scale background field. This question requires a separate study.  Applying this effect to observations one has to take into account the following issues. The gravitational potential does not evolve much on large scales especially in the Einstein-de Sitter universe (Kofman \\& Shandarin 1988; \\cite{pau-mel95}; \\cite{mel-etal96}). Therefore, the potential at present is very similar to the primordial one on scales much greater than the scale of nonlinearity. A simple explanation to this in the frame of the standard scenario of the structure formation is due to the fact that the mass has been displaced by the distance about $10 h^{-1}{\\rm Mpc}$ (\\cite{sh93}). Therefore, the potential on  scales greater than, say, $30 h^{-1}{\\rm Mpc}$ has been changed very little.  Clusters can be used as {\\it statistical} tracers of the potential. In addressing this question it is worth noting that the shot noise is an important factor since clusters are rare objects. Using the observational mass function (Bahcall \\& Cen 1993) one can estimate that an average spherical patch of the radius  $\\approx 60 h^{-1} {\\rm Mpc}$ contains about $30$ clusters with the masses greater than $10^{14} h^{-1} M_{\\odot}$. Thus, the shot noise is about 18\\% on this scale which is comparable with the bias itself (see Fig. 1, the bottom panel). However, Fig. 2 suggests that the most massive clusters [$M>10^{15} h^{-1} M_{\\odot}$] are very likely to reside in the regions of negative potential ($P > 75\\%$) and very unlikely in the regions of high potential ($P < 5\\%$ if $\\varphi > \\sigma_{\\varphi}$). More detailed analysis will be present elsewhere.  Probably, the best candidates for the markers of the troughs in the field of the primordial potential fluctuations are superclusters (defined as clusters of clusters) (\\cite{bah-son84}) especially with highest density enhancements (Shapley supercluster) and the giant geometrical patterns in the cluster distribution (\\cite{tul-etal92})."
    },
    "9803/hep-ph9803378_arXiv.txt": {
        "abstract": "The scattering of solar neutrinos on electrons is sensitive to the neutrino magnetic moments through an interference of electromagnetic and weak amplitudes in the cross section.  We show that future low-energy solar neutrino experiments with good angular resolution can be sensitive to the resulting azimuthal asymmetries in event number and should provide useful information on non-standard neutrino properties such as magnetic moments. We compare asymmetries expected at Hellaz (mainly pp neutrinos) with those at the Kamiokande and Super-Kamiokande experiments (Boron neutrinos), both for the case of Dirac and Majorana neutrinos and discuss the advantages of {\\sl low energy experiments}. Potentially interesting information on the solar magnetic fields may be accessible. ",
        "introduction": "Most non-standard properties of neutrinos arise from non-zero masses \\cite{fae,revnu}. Among these electro-magnetic dipole moments play an important role \\cite{MoPal}.  Here we are concerned with a particular effect in neutrino-electron scattering for neutrinos from the Sun which possess a Dirac magnetic moment \\cite{VVO} or transition magnetic moments \\cite{BFD} in the case of Majorana neutrinos. The latter is especially interesting first of all because it is more fundamental theoretically, and because Majorana neutrinos are the ones which arise in most extensions of the Standard Model. Moreover, the effects of Majorana transition moments can be resonantly enhanced when neutrinos propagate in media \\cite{RFSP} such as the Sun, providing one of the attractive solutions to the solar neutrino problem \\cite{akhmedov97}. Another practical advantage in favour of Majorana transition moments is that, in contrast to Dirac-type magnetic moments, these are substantially less stringently constrained by astrophysics \\cite{Raffelt}.  For {\\sl pure left-handed neutrinos} the weak interaction and the electro-magnetic interaction amplitudes on electrons do not interfere, since the weak interaction preserves neutrino helicity while the electro-magnetic does not. As a result the cross section depends quadratically on $\\mu_\\nu$.  However, if there exists a process capable of converting part of the initially fully polarized $\\nu_e$'s, then an {\\sl interference term} arises, proportional to $\\mu_\\nu$, as pointed out e.g. in ref. \\cite{Barbieri}.  This term depends on the angle between the component of the neutrino spin transverse to its momentum and the momentum of the outgoing recoil electron.  Therefore the event count rates expected in an experiment would exhibit an {\\sl asymmetry} with respect to the above defined angle.  Such asymmetry would not show up in earth-bound laboratory experiments even with stronger magnetic fields, since the helicity-flip could be caused only by the presence of a neutrino mass and is therefore negligible \\cite{grimus}. However, in the solar convective zone one may find a magnetic field extended over a tenth or so of the solar radius and, most importantly, the neutrino depolarization could be resonant in the Sun. Even if the Sun possesses only a relatively modest large-scale magnetic field $B_{\\perp} \\sim 10^4$ G in the convective region ($L\\sim L_{conv} \\simeq 3\\times 10^{10}$ cm), and for a neutrino magnetic moment of the order $10^{-11} \\mu_B$ such a spin-flip process may take place with sizeable rates, since in such a case one has $\\mu_{\\nu} B_{\\perp}L \\sim 1 $.  Barbieri and Fiorentini considered \\cite{Barbieri} the conversions $\\nu_{eL} \\to \\nu_{eR}$ in the Sun as a result of the spin-flip by a toroidal magnetic field in the convective zone. They showed that the azimuthal asymmetry could be observable in a real time solar $^8$B-neutrino experiment and as large as 20\\% for an electron kinetic energy threshold of $W_e=5$ MeV.  They chose a fixed $\\nu_e$ survival probability $P_e=1/3$ (as suggested at that time by the Homestake experiment) and the maximal Dirac magnetic moment allowed by laboratory experiments, $\\mu_\\nu\\simeq 10^{-10}\\mu_B$.  On the other hand, Vogel and Engel \\cite{Vogel} emphasized that if an asymmetry in the scattering of solar neutrinos exists, recoil electrons will be emitted copiously along the direction of the neutrino polarization in the plane orthogonal to the neutrino momentum. They calculated the asymmetry expected for solar $^8$B neutrinos with $\\mu_\\nu=10^{-10} \\mu_B$ and concluded that it would be difficult to detect because of the poor angular resolution of the experiments. Moreover, as we will see later, both \\cite{Barbieri} and \\cite{Vogel} {\\sl overestimated} the asymmetry. Thus their calculations are not accurate.  In this paper we correct results for the asymmetry in the case presented by \\cite{Barbieri} and \\cite{Vogel} for high energy $^8$B neutrinos.  In addition we compare them with the expected asymmetry in the case of a Majorana transition magnetic moment of the same magnitude. More importantly, we show the sensitivity of planned solar neutrino experiments in the {\\sl low energy region} ($\\omega \\lsim 1$ MeV) to the azimuthal asymmetries that are expected in the recoil electron event rates, arising from the above electro-weak interference term.  We calculate the asymmetry for the low-energy $pp$-neutrinos fixing the survival probability at $P_e=0.5$. This gives the maximum expected asymmetry and seems phenomenologically reasonable in order to convert the initial solar $\\nu_{eL}$'s via the Resonant Spin-Flavour Precession (RSFP) scenario.  In particular we calculate the asymmetry that could be observed in the azimuthal distribution of events in an experiment like the proposed Hellaz \\cite{Hellaz}, sensitive to the fundamental $pp$ neutrinos from the Sun. The Multi-Wire-Chamber in Hellaz should measure both the recoil electron energy $T$ and the recoil electron scattering angle $\\theta$ with good precision. Moreover Hellaz should be sensitive to the azimuthal angle $\\phi$, measuring the number of events in $\\phi$--bins.  We discuss the sensitivity of the Hellaz experiment for probing $\\mu_\\nu$ and compare it with planned accelerator experiments.  In particular there are very interesting new projects, such as the future ITEP-Minnesota experiment, where they plan to search $\\mu_{\\nu}/\\mu_B$ down to $3\\cdot 10^{-11}$ with reactor anti-neutrinos \\cite{Voloshin}, and the LAMA experiment, which will use a powerful isotope neutrino source \\cite{Bernabei}.  Finally, we also refine our calculations of the azimuthal asymmetry expected for $pp$ neutrinos at Hellaz using a realistic energy-dependent conversion probability $P_e$ based on a simple model for resonant spin flip conversions in the Sun. ",
        "conclusions": "Measuring azimuthal asymmetries in future low-energy solar neutrino-electron scattering experiments with good angular resolution should be a feasible and illuminating task.  Such asymmetries should provide useful information on non-standard neutrino properties such as magnetic moments, as well as on solar magnetic fields. The effect follows from an interference of electro-magnetic and weak amplitudes in the cross section.  We have seen that low-energy experiments such as Hellaz (sensitive mainly to pp neutrinos) should provide a much better means for the study of azimuthal asymmetries than accessible at the Kamiokande or Super-Kamiokande experiments (sensitive to Boron neutrinos).  For equal values of the magnetic moments, the expected asymmetries are larger for Dirac neutrinos than for Majorana neutrino transition moments. However, the Dirac neutrino case is probably less likely, as there is no resonant conversion in the Sun. One exception would be the case of Dirac neutrinos in the presence of twisting magnetic fields \\cite{akhpetsmi}. However, although in this case resonant conversions in matter can take place one expects (as mentioned in section 3) a washing out of the asymmetry effect due to the changing magnetic field direction. Therefore the RSFP scenario remains as the most promising possibility. It is also the most interesting one theoretically, since Majorana neutrinos are more fundamental and arise in most models of particle physics beyond the Standard Model.  Note that the discussion given above we have assumed $\\nu_e$ magnetic moments of the order $10^{-11}\\mu_B$ which is consistent with present laboratory experiments. Apart from possible effects in red giants, a $\\nu_e$ transition moment of $10^{-11}\\mu_B$ is compatible with astrophysical limits, given the present uncertainties in these considerations."
    },
    "9803/hep-ph9803244_arXiv.txt": {
        "abstract": "{ \\setlength{\\baselineskip}{16pt} \\noindent  Using the results of a high precision calculation of the solar neutrino survival probability for Earth crossing neutrinos in the case of MSW $\\nu_e \\rightarrow \\nu_s$\\ transition solution of the solar neutrino problem, we derive predictions for the one-year averaged day-night (D-N) asymmetries in the deformed recoil-e$^-$ spectrum and in the energy-integrated event rate due to the solar neutrinos, to be measured with the Super - Kamiokande detector. The asymmetries are calculated for three event samples, produced by solar $\\nu_e$\\ crossing the Earth mantle only, the core, and the mantle only + the core (the full night sample), for a large set of representative values of the MSW transition parameters $\\Delta m^2$ and $\\sin^2 2\\theta_V$ from the ``conservative'' $\\nu_e \\rightarrow \\nu_s$ solution regions, obtained by taking into account the possible uncertainties in the predictions for the $^{8}$B and $^{7}$Be neutrino fluxes. The effects of the uncertainties in the value of the bulk matter density and in the chemical composition of the Earth core on the predictions for the D-N asymmetries are investigated. The dependence of the D - N effect related observables on the threshold recoil - e$^-$\\ kinetic energy, $T_{e,th}$, is studied. It is shown, in particular, that for $\\sin^2 2\\theta_V \\leq 0.030$\\ the one year average D-N - asymmetry in the sample of events due to the core-crossing neutrinos is larger than the asymmetry in the full night  sample by a factor which, depending on the solution value of \\dms, can be $\\sim (3 - 4)$ ($\\dms ~< 5\\times 10^{-6}~{\\rm eV^2}$) or $\\sim (1.5 - 2.5)$ ($5\\times 10^{-6}~{\\rm eV^2}~ \\ltap ~\\dms ~\\ltap 8\\times 10^{-6}~{\\rm eV^2}$). We find, however, that at small mixing angles $\\sin^2 2 \\theta_V ~\\ltap~ 0.014$, the D-N asymmetry in the case of solar $\\nu_e \\rightarrow \\nu_s$\\ transitions is considerably smaller than if the transitions were into an active neutrino, $\\nu_e \\rightarrow \\nu_{\\mu(\\tau)}$. In particular, a precision better than 1\\% in the measurement of any of the three one year averaged D-N asymmetries considered by us would be required to test the small mixing angle nonadiabatic $\\nu_e \\rightarrow \\nu_s$ solution at $\\sin^2 2\\theta_V ~\\ltap ~0.01$. For $0.0075~ \\ltap~ \\sin^2 2\\theta_V ~\\leq~ 0.03$ the magnitude of the D-N asymmetry in the sample of events due to the core-crossing neutrinos is very sensitive to the value of the electron number fraction in the Earth core, $Y_e(core)$: a change of $Y_e(core)$ from the standard value of 0.467 to the conservative upper limit of 0.50 can lead to an increase of the indicated asymmetry by a factor of $\\sim (3 - 4)$. Iso - (D-N) asymmetry contours in the $\\Delta m^2 - \\sin^2 2 \\theta_V$\\ plane for the Super-Kamiokande  detector are derived in the region $\\sin^2 2\\theta_V \\geq 10^{-4}$ for the three event samples studied for  $T_{e,th} = 5~{\\rm MeV ~and~7.5~MeV}$, and in the case of the samples due to the core crossing and (only mantle crossing + core crossing) neutrinos - for $Y_e(core) = 0.467~{\\rm and~} 0.50$. The possibility to discriminate between the $\\nu_e \\rightarrow \\nu_s$ and $\\nu_e \\rightarrow \\nu_{\\mu(\\tau)}$ solutions of the solar neutrino problem by performing high precision D-N asymmetry measurements is also discussed. } \\eec  \\newpage ",
        "introduction": "\\indent In the present article we continue the systematic study of the \\daynight\\ effect for the \\SK\\ detector began in refs. \\cite{ArticleI}\\ and \\cite{ArticleII}. Assuming that the solar neutrinos undergo two - neutrino MSW $\\nue \\rightarrow \\numt$\\ transitions in the Sun, and that these transitions are at the origin of the solar neutrino deficit, we have performed in \\cite{ArticleII}\\ a high - precision calculation of the one - year averaged solar \\nue\\ survival probability for Earth crossing neutrinos, $\\PTot(\\nue\\rightarrow\\nue)$, reaching the \\SK\\ detector. The probability $\\PTot(\\nue\\rightarrow\\nue)$\\ was calculated by using, in particular, the elliptical orbit approximation (EOA) to describe the movement of the Earth around the Sun. Results for $\\PTot(\\nue\\rightarrow\\nue)$\\ as a function of $\\Enu/\\dms$, $\\Enu$\\ and \\dms\\ being the neutrino energy and the neutrino mass squared difference, have been obtained for neutrinos crossing the Earth mantle only, the core, the inner 2/3 of the core and the mantle + core (full night) for a large representative set of values of \\SdTvS, where $\\theta_V$\\ is the neutrino mixing angle in vacuum, from the ''conservative'' MSW solution region in the \\dms\\ - \\SdTvS\\ plane, derived by taking into account the possible uncertainties in the fluxes of \\BHt\\ and \\BeSv\\ neutrinos (see, e.g., ref. \\cite{SPnu96,KPUNPUB96}; for earlier studies see ref. \\cite{KS94}).  We have found in \\cite{ArticleII}, in particular, that for $\\SdTvS \\leq 0.013$\\ the one - year averaged \\daynight\\ asymmetry \\footnote{A rather complete list of references on the \\daynight\\ effect is given in ref. \\cite{ArticleI}; for relatively recent discussions of the effect see, e.g., \\cite{ Hata:Langacker:1994Earth, Baltz:Weneser:1994, Gelb:Kwong:Rosen:1996, Lisi:Montanino:1997, Bahcall:Krastev:1997, Hata:1997} } in the probability $\\PTot(\\nue\\rightarrow\\nue)$\\ for neutrinos crossing the Earth core can be larger than the asymmetry in the probability for (only mantle crossing + core crossing) neutrinos by a factor of up to six. The enhancement is even larger for neutrinos crossing the inner 2/3 of the core. We have also pointed out to certain subtleties in the calculation of the time averaged \\nue\\ survival probability $\\PTot(\\nue\\rightarrow\\nue)$\\ for neutrinos crossing the Earth, which become especially important when $\\PTot(\\nue\\rightarrow\\nue)$\\ is computed, for instance, for the core crossing neutrinos only~ \\footnote{For further details concerning the technical aspects of the calculations see ref. \\cite{ArticleII} as well as ref. \\cite{NU4DN}.}.  The results obtained in \\cite{ArticleII}\\ were used in \\cite{ArticleI}\\ to investigate the \\daynight\\ asymmetries in the spectrum of the recoil electrons from the reaction $\\nu + e^- \\rightarrow \\nu + e^-$\\ caused by the \\BHt\\ neutrinos and in the energy-integrated event rate, to be measured by the \\SK\\ experiment. We have computed in \\cite{SKDNII:spectrum}\\ the \\daynight\\ asymmetry in the recoil-$e^-$\\ spectrum for the same large set of representative values of \\dms\\ and \\SdTvS\\ from the ``conservative'' MSW solution region for which results in \\cite{ArticleII}\\ have been presented. The \\daynight\\ asymmetry in the $e^-$-spectrum was found for neutrinos crossing the Earth mantle only, the core  and the mantle + core. In \\cite{ArticleI} we have included only 12 representative plots showing the magnitude of the \\daynight\\ asymmetry in the recoil-e$^{-}$ spectrum to be expected in the case of the two - neutrino MSW $\\nue \\rightarrow \\numt$\\ transition solution of the solar neutrino problem. The spectrum asymmetry for the sample of events due to core crossing neutrinos only was found to be strongly enhanced for $\\SdTvS~\\lsim~0.013$\\ with respect to the analogous asymmetries for the mantle and for the (only mantle crossing + core crossing) neutrinos. We presented in \\cite{ArticleI} also detailed results for the one-year averaged \\daynight\\ asymmetry in the \\SK\\ signal for the indicated three samples of events. We have found that indeed for $\\sin^22\\theta_{V}\\leq 0.013$\\ the asymmetry in the sample corresponding to core crossing neutrinos can be larger than the asymmetry in the sample produced by only mantle crossing or by (only mantle crossing + core crossing ) neutrinos by a factor of up to six. We have investigate in \\cite{ArticleI} the dependence of the D-N asymmetries in the three event samples defined above on the threshold e$^{-}$ kinetic energy being used for the event selection. The effect of the uncertainties in the Earth matter density and electron number density distributions on the predicted values of the D-N asymmetries were studied as well. We derived also iso - (\\daynight) asymmetry contours in the region of $\\SdTvS~\\gsim~10^{-4}$\\ in the \\dms - \\SdTvS\\ plane for the signals in the \\SK\\ detector, produced by neutrinos crossing the mantle only, the core and the mantle + core (full night sample). The iso-asymmetry contours for the sample of events due to core-crossing neutrinos were obtained for two values of the fraction of electrons, $Y_e$, in the core: for $Y_e (core) = 0.467$ and 0.500 \\footnote{Results for other D-N effect related observables, as the D-N asymmetry in the zenith angle distribution of the events and in the mean recoil-e$^{-}$ energy, have been obtained in refs. \\cite{Lisi:Montanino:1997, Bahcall:Krastev:1997}, where iso - (D-N) asymmetry contour plots for the full night (i.e., mantle + core) signal for the \\SK\\ detector for one value of $Y_e(core) = 0.0467$ have been presented as well.}. The results derived in \\cite{ArticleI} confirmed the conclusion drawn in \\cite{ArticleII}\\ that the \\SK\\ detector will be able to probe not only the large mixing angle adiabatic solution region, but also an important and substantial part of the $\\SdTvS~\\lsim~0.014$\\ nonadiabatic region of the MSW $\\nue \\rightarrow \\numt$\\ transition solution of the solar neutrino problem. We have found, in particular, that in a large sub-region of the ``conservative'' nonadiabatic solution region located at $\\SdTvS~\\lsim~0.0045$, the D-N asymmetry in the sample of events due to the core-crossing neutrinos only is negative and has a value in the interval (-1\\%) - (-3\\%).  In the present article we realize the same program of studies for the Super-Kamiokande detector for the alternative possibility of solar neutrinos undergoing two-neutrino matter-enhanced transitions in the Sun and in the Earth into a sterile neutrino, \\nus. As is well-known, the solar $\\nu_e$ matter-enhanced transitions into a sterile neutrino, $\\nu_e \\rightarrow \\nu_s$, provide one of the possible neutrino physics solutions of the solar neutrino problem (see, e.g., \\cite{KPNPB95,Hata:Langacker:1994Earth,KPL96,SPnu96}). The reference solution region in the $\\Delta m^2 - \\sin^22\\theta$ plane, i.e., the region obtained (at 95\\% C.L.) by using the predictions of the reference solar model of Bahcall and Pinsonneault from 1995 with heavy element diffusion \\cite{BP95} ~(BP95) for the different components of the solar neutrino flux (pp, pep, $^{7}$Be, $^{8}$B and CNO), lies within the bounds: \\vspace{-1.2cm} \\bec\\beq 2.8\\times 10^{-6}~{\\rm eV}^2 ~\\ltap ~\\Delta m^2 ~\\ltap ~7.0\\times 10^{-6}~{\\rm eV}^2,~~ \\eeq\\eec \\vspace{-1.0cm} \\bec\\beq 4.8\\times 10^{-3}~ \\ltap ~\\sin^22\\theta ~\\ltap ~1.4\\times 10^{-2}~. \\eeq\\eec \\noindent The reference solution is of the small mixing angle nonadiabatic type. A reference large mixing angle (adiabatic) solution (present in the case of $\\nu_{e} \\rightarrow \\nu_{\\mu (\\tau)}$ transitions) is practically ruled out by the solar neutrino data \\cite{KPNPB95,KPL96,SPnu96}. If we allow for possible uncertainties in the predictions for the fluxes of the $^{8}$B and $^{7}$Be neutrinos \\cite{KPL96}, the solar neutrino data is described in terms of the hypothesis of $\\nu_e \\rightarrow \\nu_s$ transitions for larger ranges of values of the parameters $\\Delta m^2$ and $\\sin^22\\theta$, belonging to the intervals: \\vspace{-1.2cm} \\bec\\beq 2.8\\times 10^{-6}~{\\rm eV}^2 ~\\ltap ~\\Delta m^2 ~\\ltap ~8.0\\times 10^{-6}~{\\rm eV}^2,~~ \\eeq\\eec \\vspace{-1.0cm} \\bec\\beq 8.0\\times 10^{-4} ~\\ltap ~\\sin^22\\theta ~\\ltap ~3.0\\times 10^{-2}~. \\eeq\\eec \\vspace{-0.4cm} \\noindent and \\vspace{-1.2cm} \\bec\\beq 5.3\\times 10^{-6}~{\\rm eV}^2 ~\\ltap ~\\Delta m^2 ~\\ltap ~1.2\\times 10^{-5}~{\\rm eV}^2,~~ \\eeq\\eec \\vspace{-1.0cm} \\bec\\beq 0.13 ~\\ltap ~\\sin^22\\theta ~\\ltap ~0.55~. \\eeq\\eec \\noindent The ``conservative'' solution regions lying within the bounds determined by eqs. (3) - (4), and eqs. (5) - (6) have been obtained by treating the $^{8}$B neutrino flux as a free parameter in the relevant analysis of the solar neutrino data, while the the $^{7}$Be neutrino flux was assumed to have a value in the interval $\\Phi_{Be} = (0.7 - 1.3)~\\Phi^{BP}_{Be}$, where $\\Phi^{BP}_{Be}$ is the flux in the reference solar model \\cite{BP95}. For values  of $\\Delta m^2$ and $\\sin^22\\theta$ from the solution regions (5) and (6) the $\\nu_e \\rightarrow \\nu_s$ transitions of $^{8}$B neutrinos having energy $E_{\\nu} \\geq 5~{\\rm MeV}$ are adiabatic. However, this adiabatic solution is possible only for large values of the $^{8}$B neutrino flux \\cite{KPL96}, $\\Phi_{B} \\cong (2.5 - 5.0)~\\Phi^{BP}_{B}$, $\\Phi^{BP}_{B}$ being the reference model flux \\cite{BP95}. Such values of $\\Phi_{B}$ seem totally unrealistic from the point of view of the contemporary solar models and we consider the indicated adiabatic solution here for completeness. It should be added that in deriving the conservative solution regions represented by eqs. (3) - (4) and eqs. (5) - (6) the limit on the D-N effect derived in \\cite{Hata:Langacker:1994Earth} on the basis of the data obtained in the Kamiokande II and III experiments \\cite{KamDN} was utilized.  Let us note that the preliminary result on the D-N effect from the Super-Kamiokande experiment after approximately one year (374.2 days) of data taking reads \\cite{SKSB97}: \\vspace{-0.6cm} \\bec\\beq \\bar{A}^{SK}_{D-N} \\equiv \\, \\frac{\\bar{R}^{D} - \\bar{R}^{N}} {\\bar{R}^{D} + \\bar{R}^{N}} =  - 0.031 \\pm 0.024 \\pm 0.014, \\eeq\\eec \\vspace{0.2cm} \\noindent where $\\bar{A}^{SK}_{D-N}$ is the average energy integrated D-N asymmetry and $\\bar{R}^{D}$ and $\\bar{R}^{N}$ are the observed average event rates caused by the solar neutrinos during the day and during the night in the Super-Kamiokande detector in the period of data taking. The first error in eq. (7) is statistical and the second error is systematic. The data were obtained with a recoil$-e^{-}$ threshold energy $E_{e,th} = 6.5~{\\rm MeV}$.  In addition of performing i) detailed high precision calculations of the D-N asymmetries in the recoil-e$^{-}$ spectrum and ii) of the energy-integrated D-N asymmetries for the three samples of events (due to core crossing, only mantle crossing and only mantle + core crossing neutrinos), and of studying iii) the effects of the recoil-e$^{-}$ energy threshold variation and iv) of the uncertainties in the chemical composition and matter density of the Earth's core on the calculated D-N effect related observables, we also analyze qualitatively the possibility to distinguish between the two solutions of the solar neutrino problem involving matter-enhance transitions of the solar $\\nu_e$ respectively into active neutrinos and into sterile neutrinos, $\\nu_e \\rightarrow \\nu_{\\mu (\\tau)}$ and $\\nu_e \\rightarrow \\nu_s$, by performing high-precision D-N asymmetry measurements.  \\vspace{-0.3cm} ",
        "conclusions": "\\indent  In the present article we have performed a rather detailed quantitative study of the D-N effect for the Super-Kamiokande detector for the solution of the solar neutrino problem involving two-neutrino matter-enhanced transitions of the solar neutrinos into a sterile neutrino, $\\nu_e \\rightarrow \\nu_s$. The one year average D-N asymmetry, \\AsymRs, has been calculated (using the high precision methods developed in refs. \\cite{ArticleI,ArticleII}) for three samples of events, \\mantle\\ (M) , \\core\\ (C) and \\night\\ (N), produced respectively by the solar neutrinos crossing the Earth mantle only, the Earth core, and by the only mantle crossing + the core crossing neutrinos (the full night sample). The asymmetry calculations require the knowledge of the one year averaged spectrum of the recoil electrons and energy-integrated even rate, produced by the solar neutrinos during the day (the \\DAY\\ sample). Results for the D-N asymmetry in the recoil-e$^{-}$ spectrum for the same three samples of events, \\AsymSs (\\Te), s=N,C,M, have also been obtained. The asymmetries have been calculated for a large representative set of values of the neutrino transition parameters \\dms\\ and $\\sin^22\\theta$ from the ``conservative'' $\\nu_e \\rightarrow \\nu_s$ transition solution regions (eqs. (3) - (6)), derived by taking into account the possible uncertainties in the predictions for the $^{8}$B and $^{7}$Be neutrino fluxes. We have investigated the dependence of the three D-N asymmetries studied, on the recoil-e$^{-}$ kinetic energy threshold $T_{e,th}$, which can be varied in the Super-Kamiokande experiment, by performing calculations of all the indicated D-N asymmetries for $T_{e,th} = 5.0~$ MeV and $T_{e,th} = 7.5~$ MeV. The effect of the estimated uncertainties in the knowledge of the bulk matter density and the chemical composition of the Earth core \\cite{Stacey:1977,PREM81,CORE} on the predictions for the D-N asymmetries, has been studied as well by deriving results for \\AsymSs (\\Te)\\ and \\AsymRs\\ both for the standard value of the electron number fraction in the core $\\Ye (core) = 0.467$ and for the estimated conservative upper limit on $\\Ye (core)$, $\\Ye (core) = 0.50$ (see \\cite{CORE} and \\cite{ArticleI}). Iso-(D-N) asymmetry contour plots for the \\night, \\core\\ and \\mantle\\ samples of events in the region $10^{-7}~{\\rm eV^2} ~\\leq ~\\dms~\\leq 10^{-4}~{\\rm eV^2}$, $10^{-4}~\\leq~ \\sin^22\\theta_V ~\\leq~ 1$, have been obtained for $T_{e,th} = 5.0~$ MeV and $T_{e,th} = 7.5~$ MeV, and for the \\night\\ and \\core\\ samples - for $\\Ye (core) = 0.467$ and $\\Ye (core) = 0.50$. The main results of this study are collected in Tables I - VII and are shown graphically in Figs. 4 - 7.  We have found that, as like in the case of the $\\nue \\rightarrow \\numt$\\ solution, the division of the data collected at night into a \\core\\ and \\mantle\\ samples is a rather effective method of enhancing the D-N asymmetry at small mixing angles, $0.001~\\ltap~ \\sin^22\\theta_V ~\\ltap~ 0.03$: the asymmetry in the \\core\\ sample $|\\AsymRC|$ is larger than the asymmetry in the \\night\\ sample $|\\AsymRN|$ typically by a factor of (3 - 4) if $\\dms ~< 5\\times 10^{-6}~{\\rm eV^2}$, and by a factor of $\\sim (1.5 - 2.5)$ for $5\\times 10^{-6}~{\\rm eV^2}~ \\ltap ~\\dms ~\\ltap 8\\times 10^{-6}~{\\rm eV^2}$ (Table II). However, the enhancement is not as strong as in the case of the $\\nue \\rightarrow \\numt$\\ transition solution \\cite{ArticleI}. Moreover, in the interesting region $0.005~\\ltap~ \\sin^22\\theta_V ~\\ltap~ 0.014$ the D-N asymmetries in the \\core\\ and \\night\\ samples found for the $\\nue \\rightarrow \\nus$\\ solution, $|\\AsymRC (sterile)|$ and $|\\AsymRN (sterile)|$, are substantially smaller - at least by a factor of 4 and typically by a factor of 5 to 10, than the asymmetries corresponding to the $\\nue \\rightarrow \\numt$\\ solution, \\AsymRC (active) and \\AsymRN (active). Similar conclusion is valid for the \\mantle\\ sample asymmetries. This remarkable difference in the magnitudes of the asymmetries $|\\AsymRs (sterile)|$ and $|\\AsymRs (active)|$ in the corresponding small mixing angle solution regions is a consequence of the different roles the neutron number density distribution in the Earth $n_{n}(r)$ plays in the solar neutrino transitions in the two cases: the $\\nue \\rightarrow \\numt$\\ transitions, as is well-known, depend only on the electron number density distribution, $n_{e}(r)$, while the  $\\nue \\rightarrow \\nus$\\ transitions depend on the difference ($n_{e}(r) - 0.5~n_{n}(r)$). In the Sun one has \\cite{BP95} $0.5~n_{n}(r) \\ll n_{e}(r)$ and $n_{n}(r)$ influences little the $\\nue \\rightarrow \\nus$\\ transitions. In contrast, due to the neutrality and approximate isotopic symmetry of the Earth matter, one has in the Earth: $n_{e}(r) - 0.5~n_{n}(r) \\cong 0.5~n_{e}(r)$. This difference between  the number density distributions $n_{e}(r)$ and ($n_{e}(r) - 0.5~n_{n}(r)$) in the Earth is at the origin of the dramatic difference between the magnitudes of the D-N asymmetries corresponding to the small mixing angle $\\nue \\rightarrow \\nus$\\ and $\\nue \\rightarrow \\numt$\\ transition solutions discussed above. Correspondingly, it leads to a shift towards smaller (by a factor of $\\sim 2$) values of \\dms\\ and larger values of \\SdTvS\\ of the iso - \\daynight\\ asymmetry contours in the $\\dms - \\sin^22\\theta_V$ plane corresponding to the $\\nue \\rightarrow \\nus$\\ solution with respect to the analogous contours for the $\\nue \\rightarrow \\numt$\\ solution (compare Figs. 3a - 3c in \\cite{ArticleI} with Figs. 5a, 6a and 7a).  At small mixing angles even the \\core\\ asymmetry corresponding to the $\\nue \\rightarrow \\nus$\\ solution is rather small (Table II, Fig. 6a): for $0.0012~\\ltap~ \\sin^22\\theta_V ~\\ltap~ 0.008$ and $2.8\\times 10^{-6}~{\\rm eV^2}~ \\ltap ~\\dms ~\\ltap 4\\times 10^{-6}~{\\rm eV^2}$ we find $(-2\\%)~\\ltap ~\\AsymRC (sterile)~\\ltap ~(-1\\%)$. For other values of \\dms\\ from the small mixing angle ``conservative'' solution region $0.001~ \\ltap~\\sin^22\\theta_V ~\\ltap ~0.009$ one obtains $|\\AsymRC (sterile)| \\leq 1\\%$. We have $\\AsymRC (sterile)~\\gtap ~1\\%$ in the solution region $\\sin^22\\theta_V ~\\gtap~ 0.009$ and $3.0\\times 10^{-6}~{\\rm eV^2}~ \\ltap ~\\dms ~\\ltap ~4.4\\times 10^{-6}~{\\rm eV^2}$. In addition, \\AsymRC (sterile)\\ has a minimum in the interval $0.008~\\ltap~\\sin^22\\theta_V ~\\ltap ~0.03$ at $\\dms \\cong 6.0\\times 10^{-6}~{\\rm eV^2}$ and for this value of \\dms\\ one has $\\AsymRC \\geq 1\\%$ only when $\\sin^22\\theta_V \\geq 0.012$. The \\night\\ and \\mantle\\ asymmetries are larger than 1\\% in absolute value only if $\\sin^22\\theta > 0.010$ (Table II, Figs. 5a and 7a).  Replacing the threshold energy $T_{e,th} = 5~$MeV with 7.5 MeV can, depending on the SMA solution value of \\dms, increase $|\\AsymRC|$ (by a factor $\\sim (1.2 - 1.5)$), decrease it somewhat or leave the asymmetry practically the same; it changes little the magnitudes of \\AsymRN\\ and \\AsymRM (Tables III and IV, Figs. 5b, 6c and 7b).  The asymmetries \\AsymRC\\ and \\AsymRN, however, are rather sensitive to the value of \\Ye (core) (Tables II - IV and Figs. 5a, 5b and 6a - 6d). The dependence of \\AsymRC\\ and \\AsymRN\\ on \\Ye (core) is particularly strong in the ``conservative'' solution interval $0.0075 ~\\ltap ~\\sin^22\\theta_V ~ < ~0.030$, where a change of the value of \\Ye (core)\\ from 0.467 to 0.50 leads to an increase of $|\\AsymRC|$ and $|\\AsymRN|$ by factors of $\\sim ( 2 - 4)$.  The predicted D-N asymmetries in the recoil-e$^{-}$ spectrum for the three samples of events are small in the SMA solution region (Figs. 4.1 - 4.16). The spectrum asymmetry for the \\night\\ sample, for instance, at $\\SdTvS < 0.014$ satisfies $|\\AsymSN(\\Te)|~\\ltap~ 1\\%$ for $5~{\\rm MeV} \\leq T_e \\leq 14~{\\rm MeV}$, and is hardly observable with the Super-Kamiokande detector. This conclusion is valid both for $Y_e(core) = 0.467$ and $Y_e(core) = 0.50$. Analogous results are valid for the \\core\\ sample spectrum asymmetry \\AsymSC (\\Te): one has $|\\AsymSC(\\Te)| \\geq 4\\%$ only if $\\SdTvS \\geq 0.01$; at $\\SdTvS \\cong 0.014$ the asymmetry \\AsymSC(\\Te) reaches 16\\%.  The upper limit on the D-N asymmetry \\AsymRN\\ following from the Super-Kamiokande data (eq. (7)) rules out (at 95\\% C.L.) the ``conservative'' large mixing angle (adiabatic) solution possible in the case of solar $\\nue \\rightarrow \\nus$\\ transitions for unrealistically large values of the $^{8}$B neutrino flux \\cite{KPL96}.  A qualitative analysis performed by us indicates that the measurement of the \\core\\ and \\mantle\\ sample asymmetries, which are independent observables, can help to discriminate between the $\\nue \\rightarrow \\numt$\\ and the $\\nue \\rightarrow \\nus$\\ transition solutions of the solar neutrino problem.  The results obtained in the present study suggest that it will be difficult to probe the small mixing angle nonadiabatic $\\nue \\rightarrow \\nus$\\ transition solution of the solar neutrino problem at $\\SdTvS ~\\ltap~ 0.01$ by performing high precision measurements of the event rate and the recoil-e$^{-}$ spectrum D-N asymmetries with the Super-Kamiokande detector. The precision required to test the indicated solution region exceeds, for most values of the parameters \\dms\\ and $\\SdTvS$ from the region, the precision in the D-N asymmetry measurements which is planned to be achieved in the Super-Kamiokande experiment."
    },
    "9803/astro-ph9803309_arXiv.txt": {
        "abstract": "We discuss the emergent spectra from accreting black holes, considering in particular the case where the accretion is characterized by relativistic bulk motion. We suggest that such accretion is likely to occur in a wide variety of black hole environments, where the strong gravitational field is expected to dominate the pressure forces, and that this likely to lead to a characteristic high-energy spectroscopic signature; an extended power-law tail.  It is in the high (soft) state that matter impinging upon the event horizon can be viewed directly, and the intrinsic power-law is seen. Certain types of Active Galactic Nuclei (AGN) may represent the extragalactic analog of the high-soft state accretion, which would further support our ideas, demonstrating the stability of the ($\\alpha\\sim1.8$) power-law. This stability is due to the asymptotic independence of the spectral index on the mass accretion rate and its weak dependence on plasma temperatures. We have computed the expected spectral energy distribution for an accreting black hole binary in terms of our three model parameters: the disk color temperature, a geometric factor related to the illumination of the black hole site by the disk and a spectral index related to the efficiency of the bulk motion upscattering. We emphasize that this is a fully self-consistent approach, and is not to be confused with the more common phenomenological methods employing additive power law and black-body or multi-color disk. A test of the model is presented using observational data from the Compton Gamma Ray Observatory and the Rossi X-Ray Timing Explorer, covering $\\simeq 2-200$ keV for two recent galactic black hole X--ray nova outbursts. The resulting model fits are encouraging and, along with some observational trends cited from the literature, they support our bulk-motion hypothesis. ",
        "introduction": "Do black holes interact with an accretion flow in such a way that a unique observational signature can be identified -- that is one  which is entirely distinct from those associated with other compact objects, based solely on the radiation observed at infinity? This is a crucial question confronting both theoretical and observational astrophysicists today; for recent reviews of the astrophysics of black holes [see e.g., \\cite{liang97}, \\cite{zhang97a}]. Certainly a large body of evidence has been accumulated which supports the {\\it existence} of black holes, the most convincing arguments being those invoking dynamical mass determinations [e.g. \\cite{orosz97}]. Other arguments have recently been advanced suggesting that X-ray nova flux histories demonstrate the existence of  black-hole horizons [\\cite{narayan96}], and in AGN, asymmetrical line features have identified and interpreted as originating in massive black hole environments [\\cite{tanaka95}, \\cite{fabian95}]. Also,  quasi-periodic oscillations (QPOs) have been attributed to dynamical time scales associated with the innermost stable orbits in black hole binaries [\\cite{mrg97}].  A distinct feature of black hole spacetime geometries, as opposed to those associated with  other compact objects, is the presence of the event horizon. Near the horizon the strong gravitational field is expected  to dominate the pressure forces and thus drive the accreting material into free fall. In contrast, for other compact objects the pressure  forces are dominant near the surface and the free fall state is absent. Recently, Titarchuk and Zannias (1998) (hereafter TZ98) have developed the relativistic radiative  transfer theory demonstrating that  high-energy photons are produced by upscattering  from the converging inflow within a few Schwarzschild radii. Only some  fraction of the radiation emitted by the accretion disk illuminates the converging inflow site. It can be such a situation that the radiation density (or pressure)  determined by the injected energy of those soft disk photons and by the weak amplification they experience [\\cite{tmk97}; hereafter TMK97] is much smaller than the Eddington value. {\\it We argue that then this difference is crucial and it results in a unique observational signature for accreting black holes}. \\par \\noindent As explained above, this signature originates from upscattering of low energy photons by fast moving electrons  with velocities, $v$, approaching the speed of light, $c$. A soft photon of energy $E$, in the process of multiple scattering off  the electrons, gets substantially blue-shifted to energy \\begin{equation} E^{\\prime}=E{{1-(v/c)\\cos\\theta}\\over{1-(v/c)\\cos\\theta^{\\prime}}} \\end{equation} \\noindent due to  Doppler effect provided at least one photon is scattered in the direction of electron motion (i.e. when $\\cos\\theta^{\\prime}\\approx 1$). For example, in the first scattering event we assume the direction of incident photon,  $\\theta_1$, is nearly normal to the electron velocity, and the direction of the scattered photon is nearly aligned with the electron velocity. In the process the its outward propagation through the converging-inflow medium, the angle between the photon and electron velocity increases. Thus, in the second event the cosine angle, $\\cos\\theta_2$, tends to approach zero. The angle of outgoing photon, $\\theta_2^\\prime$, has to be large enough,  in order for the Doppler boosted photon to reach an observer.  Any system having a disk structure around  a compact object is expected to have a  source of low-energy photons. {\\it The boosted photon component is seen as the extended power-law at energies much higher than the characteristic energy of the soft photons. And it is entirely independent of the initial spectral and spatial distributions of the low-energy photons.} The spectral index of the boosted photon  distribution is determined only by the mass accretion rate and the plasma temperature of the bulk flow. {\\it The presence of this high-energy power-law  component is a generic feature of the model.} A key ingredient in support of our claim comes from the exact relativistic transfer calculations describing the Compton scattering of the low-energy radiation field of the Maxwellian distribution of fast moving electrons (TZ98). It was proven mathematically that the power law is always present as a part of the black hole spectrum over a wide energy range, extending up to 500 keV. A turnover in the spectrum at about this energy, i.e. at E$\\ltorder m_ec^2$, is a prediction of our model. Other extended power-law components, which may be related to the relativistic electron motion, e.g. in a jet, are not uniquely constrained to this energy band because they are not tied with the electron rest mass $m_ec ^2$. The observations with CGRO/OSSE could in principle confirm or refute this prediction. In practice, the data  thus far obtained are signal-to-noise limited and cannot address this issue in a definitive manner. \\par \\noindent In this letter we extract the most important points regarding the radiative transfer and  present observational evidence which supports the model.  \\section { Bulk-Motion Spectral Models} It has been shown elsewhere (TMK97, TZ98, \\cite{ct95}; hereafter CT95) that two effects, the bulk motion upscattering and the Compton (recoil) downscattering (herein BMC), compete forming the hard tail of the spectrum as an extended power law. The soft part of the spectrum comes from the disk photons seen directly and a subset of those photons which escape from the BMC atmosphere after undergoing a few scatterings but without any significant  energy change. It has also been shown that without taking  into account special and general  relativistic effects, one is able to reproduce the main features of the full relativistic formalism: the overall spectral energy distribution and the dependence of the high-energy power on mass-accretion rate (TMK97; TZ98); also, refer to  recent calculations by Laurent \\& Titarchuk (1998) (hereafter LT98).  In the relativistic treatment, the Compton downscattering becomes less efficient at high energies, due to Klein-Nishina effects.  Also, there is a possibility that  the electron distribution in the converging inflow can deviate from a Maxwellian,  flattening at high velocities, since there is insufficient time for it to thermalize. The hard photon power-law thus extends to higher energies. At the same time however, the spectrum is steepened as a  result of gravitational redshift effects.  We have assumed that there is an external illumination of the converging flow by the low-energy black body radiation of an accretion disk having a characteristic temperature $T_{c}$. Furthermore we have assumed that this illuminating radiation impinges on the BMC atmosphere with a certain geometry, which we have paramaterized in terms of a ''fraction\" $ f$. This fraction is really the first expansion coefficient of the spatial source photon distribution over the set of the eigenfunctions of the BMC formulation (TMK97, Eq. 30). As we mentioned above (\\S 1) the spectral index  is independent of the illumination fraction, $f$. It is clearly demonstrated in TMK97 (Figs 4).  We remind the reader that all reasonable theoretical spectra must exhibit a smooth transition from blackbody-like spectrum to a pure power-law, typically, at energies in the 5--12  keV range for stellar black holes.  In the soft state when the accretion rate is higher, the soft photons from the the Keplerian disk cool the hot region (Compton cloud) due to thermal Comptonization and free-free emission (CT95). The cooler converging inflow, as it rushes towards the black hole, scatter the soft-photons within the a radius, $r\\sim \\dot m r_s$ -- some of the photons  then undergo outward radiative diffusion. Here $\\dot m=\\dot M/\\dot M_E$, $r_s$ is Scwarzschild radius, $\\dot M$ is the net accretion rate (including accretion from the disk plus any halo or other non-keplerian component), $\\dot M_E \\equiv L_E/c^2=4\\pi GMm_p/ \\sigma_Tc~$ is the Eddington accretion rate, $M$ is the mass of the central object, $m_p$ is the proton mass and $G$ is the gravitational constant. It transfers its momentum to the soft-photons  to produce the power-law component extending to energies comparable to the kinetic energy of electrons in the converging inflow, i.e. of order $m_ec^2$. On the other hand {\\it in the hard state, the hot emission cloud covering the BMC zone prevents us from seeing the photons that are upscattered to subrelativistic energies within a few Scwarzschild radii}.  The luminosity of the upscatterd component, the hard power law, has to be very small compared to the Eddington luminosity in order for the BMC model to be valid. The relative normalization of the soft component to the hard power-law is less important provided the inferred luminosity of the hard power-law remains consistent with the assumption of negligible radiation pressure near the black-hole horizon.  The BMC spectral model can be described as the sum of a thermal (disk) component and the convolution of some fraction of this component $g(E_0)$ with the upscattering Green's function $I(E,E_0)$ (TMK97, Eq. 30). The Green's function has the form of a broken power-law with spectral indices $\\alpha$ and $\\alpha+\\zeta$ for high  $E\\geq E_0$ and low $E\\leq E_0$ energy parts respectively, \\begin{equation} F_{\\nu} (E)=\\int_0^{\\infty}I(E,E_0)g(E_0)dE_0. \\end{equation} \\par \\noindent The above convolution is insensitive to the value of the Green's function spectral index $\\alpha+\\zeta$, which is always much greater than one. TZ98 presented rigorous proof that the hard power-law tail is a signature of a Schwarzschild black hole. Furthermore, the same statement is valid for the case of a  rotating (Kerr) black hole, although a higher mass accretion rate is required to provide the same efficiency for the soft photon upscattering.  We note that the processes of absorption and emission (as free-free or synchrotron radiation) can be neglected provided the plasma temperature of the bulk flow is of order 1 keV or greater for characteristic number densities of order $10^{18}$ cm$^{-3}$ and for magnetic field strengths in the proximity of the black hole of order $10^4-10^5$ gauss or less (CT95).  \\section {Application to  Recent High Energy Observations} As a test of the model we have collected data resulting from high-energy observations covering recent activity periods in two galactic X--ray novae: GRO J1655--40 [\\cite{zhang97b}] and GRS 1915+105 [\\cite{chaty96}]. X--ray novae comprise perhaps the best test case of the methodology described here, since they are in low-mass binary systems -- avoiding  the added complications which may arise from the OB star winds in high-mass binary BHCs such as Cygnus X--1 -- and because they become exceptionally bright in outburst exhibiting  frequent and pronounced high-energy spectral state transitions [e.g. \\cite{csl97}, \\cite{ebisawa94} \\cite{esin97}]. Furthermore, as a group they comprise the most convincing galactic black--hole candidates.  We constructed composite high-energy spectra for GRO J1655--40 during an outburst in the spring of 1996 (\\cite{hynes98}), covering the ~2-200 keV spectral region and fit these data by the BMC model. This was accomplished using summed, standard mode (128 channel) data from the RXTE/PCA and the 16-channel BATSE/LAD earth-occultation data bracketting the pointed observations. In addition, there was substantial outburst activity in GRS 1915+105 during the latter part of 1996 [e.g. \\cite{bandy98}]. We utilized some of the available data from the same instruments for this event as well.  The BMC model described in section 2 was imported into the ''XSPEC\" software package which was used to perform all of the model fitting described here. Our resulting fits are shown in Figure 1.  For GRO J1655-40 we obtained a blackbody color temperature  of $kT_{c}=1.1\\pm0.1$~keV for the soft photon source, a energy spectral index of $\\alpha=1.60\\pm0.03$, and a geometric factor $f$, parameterizing the fraction of the total soft photon flux illuminating the BMC inflow atmosphere, of $f=0.32\\pm0.02$. The observed 2-100~keV flux was $5.7\\times 10^{-8}$ ergs/cm$^2$/s. We note that the corresponding luminosity in the hard power-law component is about $1.5\\%$ of $L_{E}$ (with an assumption of the distance to the source, $3.2$ kpc and the mass of the central object,  7 solar masses), which is consistent with our assumption of negligible radiation pressure near the event horizon.  Similar results; $kT_{c}=0.9\\pm0.1$, $\\alpha=1.68\\pm0.03$  and $f=0.72\\pm0.02$, were obtained for GRS 1915+105.  From the inferred 2-200~keV luminosity for GRO~J1655-40, $\\sim 5\\%$ of $L_{E}$, we derive a mass accretion rate (in Eddington units) of order 1, bearing in mind the efficiency of gravitational to radiative energy conversion is of order 5\\% or less, e.g. Shakura \\& Sunyaev 1973 (hereafter SS73). This value of $\\dot M$ is consistent with expected values within the BMC framework. This suggests that the line-of-sight column density of the BMC atmosphere is of order $10^{24}$ cm$^{-2}$ (see, TMK97, Eq. 2). However, because the best-fitted color temperature is about 1 keV, (and this is a lower bound on the BMC plasma temperature) we conclude that the detected X-ray spectrum is not significantly modified by absorption (see also \\S 2).  The temperatures we infer for these two sources are somewhat lower than values previously reported [e.g. Zhang et al. (1997b)], however this is to be expected. As noted, our procedure represents a fully self-consistent model deconvolution, whereas most previous approaches are phenomenological, i.e. power law plus black body or multicolor disk. Mathematically, one expects the power law component contribute significant soft-energy flux with the net effect of skewing the thermal residual to higher apparent temperatures. This will not occur with approach, as the hard power law turns over towards low energies (see Fig 3, TMK97). The inferred spectral indices also agree extremely well with our model predictions. In TZ98, calculations of the $\\dot m - \\alpha$ relationship were presented. For the low temperature limit, an asymptotic lower limit of $\\alpha\\simeq1.8$ was calculated; for the higher BMC plasma temperatures (of order 10 keV)  this limit is significantly lower, $\\sim1$, and for mass accretion rates of  $\\dot m\\sim1$ (see below), $\\alpha$ is precisely in the $1.5-1.8$ range we find (LT98, Titarchuk 1998). Thus, we feel our observational test provides extremely encouraging support of our methodology.  Using our inferred  color temperature $T_c$ and spectral index $\\alpha$, along with the  measured flux normalization, we can estimate the mass accretion rate, black hole mass and source distance within the framework of standard accretion disk theory (e.g. SS73). This additionally requires certain assumptions  regarding a ''hardening factor\"  -- the ratio of color temperature to the effective plasma temperature (\\cite{shimura95}, hereafter ShT95). In fact, the spectral index (TMK97, TZ98, LT98) depends on the mass accretion rate and the plasma temperature; the disk color temperature  is  $\\propto (\\dot m/m)^{1/4}$ (SS73), and the normalization is $\\propto \\dot m m/d^2$ (where $d$ is the distance to the source).  An ideal test case is GRO~J1655-40, since its mass and distance are known to a high degree of accuracy (relative to other BHCs). The distance we infer, $3.8\\pm1.4$ kpc, is consistent (at the $1-\\sigma$ level) with previous determinations (\\cite{Hjellming95}). Also, the black hole mass we calculate can be reconciled with determinations from dynamical studies (\\cite{orosz97}) provided we used the hardening factor 1.9 (ShT95). This assumed a mass accretion rate $\\dot m=3$, which was the value obtained from Monte Carlo simulations performed to calculate the spectral index dependence on the mass accretion rate and plasma temperature (LT98). Again, this is an encouraging result suggesting that with further refinement, one has a method of mass and distance determination independent of the conventional quiescent spectroscopic and photometric studies, which are not always plausible. ",
        "conclusions": "The successful application of our basic model to observational data and the inferred physical parameters of those systems in comparison to independent determinations is  encouraging.  We postulate that   (i){\\it The soft state detected in GRO J1655-40  and GRS 1915+105 represents a generic feature of accreting Galactic black holes. An extragalactic analog may now be evident in the Narrow Line Seyfert 1 (NLS1) galaxy population.} (ii) {\\it The BMC spectra represent a characteristic signature of black hole horizons. The disk flux, of order 5\\% $L_{edd}$ tends to cool the ambient environment and the generic hard power-law components is seen by the observer.} (iii) {\\it Because this spectral feature is formed very close to the horizon, ($2-3R_s$), the variability timescales of high-energy line and continuum radiation should be associated with the  crossing time scale $t_{cross}\\sim10^{-5}M/M_{\\sun}$} s. (iv) {\\it The  variability seen in the soft component is not expected to be correlated with the hard component. It is related to the illumination geometry of a small area of the black hole horizon site, whereas  the soft radiation seen by the observer directly emanates from a major fraction of the entire disk which comprises a much larger area.}  (v) {\\it QPOs emanating from the inner edge of the accretion disk should lead to a pronounced hard- X-ray variability signature, because the seed photons for the converging flow upscattering come from the same inner disk region.} (vi) {\\it The appearance  of  an  additional bump in the energy range 10-20 keV can be explained in terms of downscattering (reflection) effects (ST80)  from the inner edge of the accretion disk.}  Our conclusions are consistent with various observations of Galactic and extragalactic black hole systems. For example,  in NLS1s the X-ray power-law is significantly steeper and its normalization is more variable,  with time scale of order  $10^4$ s, than in broad-line Seyfert 1 galaxies. This suggests the NLS1s may represent the extragalactic analog of the high-soft state (\\cite{pounds95}, \\cite{brandt97}, \\cite{comastri98}). Several groups have independently reached similar conclusions (CT95, \\cite{pounds95}). We further note that the equivalent widths of Fe features detected  tend to be large, in some cases $\\sim500$ eV (\\cite{comastri98}, hereafter C98).  It is worth noting that the  detection of the strong hydrogen-like iron line is expected if the source of hard energy photons ($>7$ keV) is located inside the the converging inflow region(and alternatively, line radiation could form in the cooler, Compton cloud ambient to the converging inflow region). It is easy to show that the ionization parameter $\\xi=L/(r^2n)\\approx 10^5$ ergs~cm~$s^{-1}$ is a typical value for the converging inflow and it is almost independent of the central object mass. In this case, only the hydrogen-like iron would be expected (Kallman and McCray 1982), which has been confirmed recently by BeppoSAX observations (C98).  Another supportive example is the black hole X-ray binary LMC X-3 which  appears to always be in the high state. Its hard-tail component varies independently of the soft component [e.g., \\cite{ebisawa93}].  Recent RXTE observations of Cyg X-1 during a state transition [\\cite{cui97}] revealed a striking decrease in the soft-to-hard photon lag times as the source passes from the hard to soft state. This is very strong empirical evidence that the soft seed  photons, which comprise a of fraction the disk thermal component, and the hard-X-ray power law emanate from a common compact region, again consistent with our model.  The detection of 67 Hz QPOs  from GRS~1915+105 by RXTE was recently reported by \\cite{mrg97}.  It was  clearly demonstrated that this feature is  associated with  the high energy  component visible in the PCA. This can be explained in terms of a QPO in the inner edge of the disk or by $g-$mode disk oscillations occurring within the characteristic radius of 4 $r_s$ [\\cite{tlm98}]. This should lead to variations in the hard spectral component since significant changes in the illumination geometry  of the converging inflow site, can occur (TMK97, TZ98).  Similar intrepretation can be applied to the  300 Hz QPOs detected by RXTE in GROJ1655--40 [\\cite{rem97}].  In conclusion,  we wish to emphasize once again that the observations presented here, along with some observational trends presented by others in the literature, and the relativistic theory prompt us to claim that {\\it we have identified a generic spectral signature black hole accretion.}  \\vspace{0.2in}  \\centerline{\\bf{ ACKNOWLEDGMENTS}} We wish to acknowledge the anonymous referee for who made a number of useful comments on the initial draft, as well as Jean Swank and Menas Kafatos for discussion and useful suggestions. Portions of this work were supported by the Rossi X--Ray Timing Explorer and Compton Gamma Ray Observatory Guest Observer Programs. L.T. also would like to acknowledge support from NASA grant NAG5-4965.  \\newpage"
    },
    "9803/astro-ph9803286_arXiv.txt": {
        "abstract": "\\rightskip=\\leftskip The Fornax cluster galaxy FCC 35 shows an unusual multiply-peaked integrated \\ion{H}{1} profile (Bureau, Mould \\& Staveley-Smith 1996).  We have now observed FCC 35 with the Australia Telescope Compact Array (ATCA) and have found a compact \\ion{H}{1} source with $M_{HI}$ = 2.2 $\\times$ $10^{8}$ $M_{\\odot}$, and a spatially overlapping complex of \\ion{H}{1} gas with the same mass. By combining optical observations with the \\ion{H}{1} data, we are able to identify FCC 35 as a young compact source of star formation with a nearby intergalactic \\ion{H}{1} cloud which is devoid of stars.  We classify FCC 35 as a blue compact dwarf (BCD) or \\ion{H}{2} galaxy, having large amounts of neutral hydrogen, very blue colors ($(U-V)$ = 0.1), and a low metallicity spectrum with strong narrow emission lines.   Together with the presence of the \\ion{H}{1} cloud, this suggests that FCC 35 is the result of a recent interaction within the Fornax cluster. ",
        "introduction": "Our interest in FCC 35 began in 1994 when this galaxy was observed as part of an investigation of the Tully-Fisher relation in Fornax (Bureau, Mould, \\& Staveley-Smith 1996; hereafter BMS).  Optically, FCC 35 was identified as a member of the Fornax cluster by Ferguson (1989) and classified as a possible BCD/Sm IV.  BMS photometry revealed a high surface brightness and an offset nucleus, despite a regular light profile.  The 21 cm single-dish observations of FCC 35 showed three distinct \\ion{H}{1} peaks within a 700 km s$^{-1}$ velocity range (see Fig.~\\ref{fig:f1}).   There were no other known \\ion{H}{1} sources within the Parkes' beam, so the explanation for the three-peaked profile was unknown.  This anomaly served to motivate further studies of FCC 35.  Generally, BCDs like FCC 35 are high surface brightness dwarf galaxies which appear to be undergoing an intense period of star formation (Thuan \\& Martin 1981).  The cause of these star formation bursts is not well understood, but in many cases interaction is considered a likely mechanism.  Interactions can induce rapid star formation (Bushouse 1987) and the resulting internal motions can continue to induce bursts for up to 10$^8$ yrs after the interaction (Noguchi 1991).  There are several types of interaction which can trigger the BCD phenomenon.  The first involves an encounter between two spiral galaxies and the formation of \\ion{H}{1}-rich tidal tails. BCDs have been observed and modeled to form at the end of these tails which can extend for hundreds of kpc (e.g. Duc et al. 1997; Barnes \\& Hernquist 1992, 1996; Elmegreen, Kaufman \\& Thomasson 1993; Mirabel, Lutz \\& Maza 1991; Schweizer 1978).   The surrounding environment is often left in a disordered state for some time after this type of interaction.  Another type of starburst-inducing interaction involves an intergalactic \\ion{H}{1} cloud of similar mass to the progenitor galaxy (Taylor, Brinks \\& Skillman 1993, Taylor 1997).  Taylor et al. (1995, 1996) completed a survey of relatively isolated BCDs (also called H~{\\footnotesize II} galaxies) and found $\\approx$57\\% of these to have an \\ion{H}{1} companion which could be triggering the star formation burst.  Finally, close encounters between two galaxies of different mass often induce star formation in the smaller component, possibly creating a BCD (Ostlin \\& Bergvall 1993; Lacey \\& Silk 1991).  Putting all of these possibilities together strongly suggests that the star formation bursts which produce BCDs are indeed due to interactions.  In a cluster environment these interactions are much more likely to occur than in the field.  The denser the environment, the higher the potential for an interaction among member galaxies.  Fornax is a well-studied nearby cluster (d = 18.2 Mpc; Madore et al. 1996) which has one of the highest galaxy volume densities in the Local Supercluster (Held \\& Mould 1994) and more than twice the central surface density of Virgo (Ferguson \\& Sandage 1988).  The relatively low velocity dispersion of Fornax ($\\approx$ 400 \\kms) further favors interaction between cluster members. FCC 35 is therefore in an ideal environment to be affected by the mechanisms described above.  The ATCA observations of FCC 35, presented here, reveal two distinct \\ion{H}{1} sources; one compact and regular associated with the optical FCC 35, and one extended and irregular with no optical counterpart.  The sources overlap spatially but are separated by 140 km s$^{-1}$ in velocity. These observations indicate that FCC 35 has an \\ion{H}{1} companion of comparable mass.  This brings us to the various interaction scenarios.  The starburst which is now FCC 35 may be due to a previous interaction with this \\ion{H}{1} source or both components could be the result of a spiral-spiral interaction.  It is also conceivable that the \\ion{H}{1} companion itself resulted from a gas outflow associated with the star formation burst. Investigating these possibilities is one of the central topics of this paper.  In this paper we present a combination of data which helps to reveal the origin and evolution of FCC 35.  We discuss the ATCA observations and data reduction in $\\S$2.1 and the optical imaging and spectroscopy in $\\S$2.2.  In $\\S$3.0 we present the results of these observations, including the \\ion{H}{1} distribution ($\\S$3.1), \\ion{H}{1} kinematics ($\\S$3.2), stellar distribution ($\\S$3.3), and physical conditions of the gas ($\\S$3.4).  Finally, in $\\S$4.0, we discuss these results and their implications for the formation and evolution of dwarf galaxies, in particular with respect to various interaction scenarios.  The 21 cm and optical observations together provide a unique source of information on the nature of BCDs and their companions in clusters. ",
        "conclusions": "\\subsection{The \\ion{H}{1} Cloud}  A position-velocity cut through the center of both FCC 35 and the \\ion{H}{1} cloud (Fig.~\\ref{fig:f13}) shows that despite the spatial overlap , the two components are not connected in velocity space. The cloud may either be a unique source within the Fornax cluster or a foreground or background \\ion{H}{1} object.  The former seems to be the most likely considering the velocity of the cloud (V$_{r,\\odot}$ = 1658 km s$^{-1}$) and the mean heliocentric velocity of the Fornax cluster ($\\langle{v}\\rangle = 1450 \\pm$ 34 km s$^{-1}$, $\\sigma_v$ = 350 km s$^{-1}$; Held \\& Mould 1994).  The Fornax cluster has a central number density of 500 galaxies Mpc$^{-3}$ (Ferguson 1989), and tidal debris from interactions should be expected (e.g. Theuns \\& Warren 1997).  Some of the intergalactic material may be primordial, but in a dense cluster it is likely to have arisen from galactic harassment (Moore et al. 1996). This is especially true when considering the proximity of our \\ion{H}{1} cloud to the center of the Fornax cluster ($03^{h}35^{m}, -35.7^{\\circ}$; Ferguson 1989). The interaction which formed the cloud may not have involved the galaxy FCC 35, and the cloud may simply be ``passing by'' at this stage.  Its structure could be affected by FCC 35's presence (see Figs.~\\ref{fig:f2} \\&~\\ref{fig:f4}), but this is difficult to confirm due to the irregular kinematics of the cloud and the sensitivity of its inferred structure on the weighting used in the reduction.  It is also possible that the \\ion{H}{1} cloud is a remnant of an interaction in which FCC 35 was involved.  We note that NGC 1316C is located only 6$^{\\prime}$ away (in projection) from FCC 35 ($\\Delta$V = 150 km s$^{-1}$) and the two could have interacted in the past.  It is conceivable that FCC 35 and the cloud were a single object which was ripped apart tidally into two parts of comparable mass.  This seems unlikely, however, considering the regular and compact structure of FCC 35.  Yet another possibility is that the \\ion{H}{1} cloud {\\em and} FCC 35 formed through an interaction between two spirals.  Dwarf galaxies and massive \\ion{H}{1} clouds have been predicted to form (Barnes \\& Hernquist 1992; Elmegreen et al. 1993) and observed forming (Schweizer 1978; Mirabel et al 1991) in the tidal tails which arise from these interactions. The \\ion{H}{1} cloud could then be the remaining gas from a tidal tail.  This possibility will be discussed further in the next section.  \\subsection{FCC 35}  The amount of neutral hydrogen in FCC 35 ($M_{HI}/M_{Tot}$ = 0.5), together with the optical data, presents a picture of a blue compact dwarf (BCD) or \\ion{H}{2} galaxy. Exponential surface brightness profiles are typical of BCDs (see Fig.~\\ref{fig:f10}), as are offset nuclei (Fig.~\\ref{fig:f9}a; Drinkwater \\& Hardy 1991). The spectrum of FCC 35 is also similar to that of the general BCD population, with relatively low metallicity and strong narrow emission lines (Fig.~\\ref{fig:f12}; Izotov et al. 1997; Masegisa et al. 1994; Walsh \\& Roy 1993; Thuan \\& Martin 1981). The strong H${\\alpha}$ and [OIII] emission lines are a signature of the star formation activity in FCC 35, as are the blue colors towards the galaxy's nucleus (see Table~\\ref{tab:t7}).  The formation of BCDs is still not understood. Interaction may be responsible for a significant fraction, but there are also explanations based upon an evolutionary sequence among dwarf galaxies.   Davies \\& Phillipps (1988) propose a sequence, dI$\\leftrightarrow$BCD$\\leftrightarrow$dE, which involves repeatedly induced star formation bursts and explains the similarities between different types of dwarf galaxies.  The trigger of the BCD phenomenon is described by Gordon \\& Gottesman (1981).  They find the majority of dwarf irregulars to have an extended \\ion{H}{1} halo and suggest that the infall of this halo fuels the star formation bursts.  FCC 35 does have an extended \\ion{H}{1} halo, but the presence of the \\ion{H}{1} cloud suggests that this star formation burst is interaction related.  One type of interaction which has been found to produce BCDs and intergalactic \\ion{H}{1} clouds is the interaction between two spiral galaxies.  The tidal tails formed as a result of these interactions can extend for hundreds of kpc (Hibbard \\& Van Gorkom 1996) and create objects of up to 10$^9$ M$_{\\odot}$ (Elmegreen et al. 1993).  The tidal features often have low mass-to-light ratios (Hibbard \\& Van Gorkom 1996), and the models of Barnes and Hernquist (1992) predict that these objects would have very little dark matter. Elmegreen et al. (1993) also predict that dwarf galaxies formed as a result of this type of interaction should contain old stars from the original disks plus new stars from the interaction-induced star formation bursts. FCC 35's upper limit on the ionized gas abundance (Z $\\leq$ 0.25 Z$_{\\odot}$) is consistent with tidal formation from the outer disk of a spiral galaxy.  Indeed, FCC 35 fulfills many of the criteria related to the spiral-spiral interaction scenario (see Table~\\ref{tab:t3}). It has a relatively low $M_{tot}$/$L_V$ and a (corrected) rotation curve (Fig.~\\ref{fig:f8}) which indicates a truncated mass distribution and (presumably) small amounts of dark matter. However, if this is the origin of FCC 35, we would perhaps expect to observe more remnant \\ion{H}{1} in the surrounding region. This was not apparent in the channel maps obtained within the 43$^{\\prime}$ primary beam of the ATCA.  FCC 35's star formation burst could also have been induced through an interaction with its closest neighbor, NGC 1316C, which has not been detected in \\ion{H}{1}.  However, this scenario appears unlikely when the age of FCC 35's star burst is taken into account.  Its color,  (U-V)$_{26\\arcsec}$ = 0.1, corresponds to an age of about 10$^{7}$ years (Larson \\& Tinsley 1978).  If FCC 35's star formation burst had been directly triggered by an interaction with NGC 1316C, the relative speed of NGC 1316C would need to be at least 3000 km s$^{-1}$. This is excessive given the low velocity dispersion of the Fornax cluster.  However, we recall that internal motions resulting from interactions can induce later star formation bursts (Noguchi 1991), so our arguments do not conclusively exclude this possibility or the spiral-spiral interaction scenario.  The most plausible interaction-related cause for the star formation burst in FCC 35 is that the \\ion{H}{1} cloud, whatever its origin, has triggered it. Its relative mass and its proximity in space and velocity make it likely that the cloud is interacting with FCC 35.  This situation is not uncommon: Taylor et al. (1994) find that galaxies with \\ion{H}{1} companions tend to have a very low mass-to-light ratios and the mass of the companion is often only an order of magnitude smaller than the mass of the galaxy (see also Walter et al. 1997). The relative velocity, projected separation, and masses of the \\ion{H}{1} cloud and FCC 35 clearly show that they are unbound, so it seems likely that they will drift apart as FCC 35 fades into a low surface brightness or irregular dwarf galaxy."
    },
    "9803/astro-ph9803235_arXiv.txt": {
        "abstract": "We present an unbiased method for evaluating the ranges of ages and metallicities which are allowed by the photometric properties of the stellar populations that dominate the light of early-type galaxies in clusters. The method is based on the analysis of morphologically-classified early-type galaxies in $17$ clusters at redshifts $0.3\\simlt z\\simlt0.9$ and in the nearby Coma cluster using recent stellar population synthesis models that span a wide range of metallicities. We confirm that metallicity effects must play a role in the origin of the slope of the color-magnitude relation for cluster early-type galaxies. We show, however, that the small scatter of the color-magnitude relation out to redshifts $z\\sim1$ does not formally imply a common epoch of major star formation for all early-type galaxies. Instead, it requires that galaxies assembling more recently be on average more metal-rich than older galaxies of similar luminosity. Regardless of the true ages and metallicities of early-type galaxies within the allowed range, their photometric properties and the implied strengths of several commonly used spectral indices are found to be consistent with {\\it apparently} passive evolution of the stellar populations. Also, the implied dependence of the mass-to-light ratio on galaxy luminosity is consistent with the observed trend. The results of our unbiased analysis define the boundaries in age and metallicity that must be satisfied by theoretical studies aimed at explaining the formation and evolution of early-type galaxies in clusters. ",
        "introduction": "Early-type galaxies in clusters exhibit a linear color-magnitude (CM) relation indicating that bright galaxies are systematically redder than their faint cluster companions (Visvanathan \\& Sandage 1977\\markcite{vs77}). This remarkable relation shows very small scatter ($\\pm 0.05$ magnitude) in high precision photometry of local clusters such as Coma and Virgo (Bower, Lucey \\& Ellis 1992a\\markcite{ble92a}, 1992b, hereafter BLE92) \\markcite{ble92b} and can be extended to clusters at medium-to-high redshift ($0\\simlt z\\simlt 1$) (Ellis et al 1997, Stanford, Eisenhardt \\& Dickinson 1998)\\markcite{el97}\\markcite{sed98}. A first attempt at explaining the universality of the CM relation involves using the age of each galaxy as the main determinant of its color. Ageing stellar populations redden progressively as stars with decreasing initial mass evolve off the main sequence. Therefore, if the colors of cluster galaxies are purely controlled by age, the small scatter about the CM relation implies a nearly synchronous star formation process for all galaxies of a given mass, while the slope of the CM relation implies systematically older ages for more massive galaxies. As shown most recently by Kodama \\& Arimoto (1997\\markcite{ko97}), such a picture is highly unlikely because it does not preserve the slope nor the magnitude range of the CM relation in time.  Another important factor that affects the colors of stellar populations is metallicity. At fixed age, a more metal-rich stellar population will appear redder and fainter than a more metal-poor one (e.g., Worthey 1994 \\markcite{wo94}). Hence, increasing metallicity at fixed age has a similar effect on colors as increasing age at fixed metallicity. This is usually referred to as the {\\it age-metallicity degeneracy} (Worthey 1994\\markcite{wo94}). Several studies have shown that CM relation of cluster elliptical galaxies could be primarily driven by metallicity effects (Larson 1974\\markcite{lar74}; Matteuci \\& Tornamb\\'e 1987\\markcite{mator87}; Arimoto \\& Yoshii 1987 \\markcite{ari87}; Bressan, Chiosi \\& Tantalo 1996\\markcite{br96}; Kodama \\& Arimoto 1997\\markcite{ko97}). The physical mechanism usually involved is that of a galactic wind: supernovae-driven winds are expected to be more efficient in ejecting enriched gas, and hence in preventing more metal-rich stars from forming, in low-mass galaxies than in massive galaxies with deeper potential wells. Although age is generally assumed to be the same for all galaxies in these studies, this has not been proven to be an essential requirement. In fact, scenarios in which E/S0 galaxies progressively form by the merging of disk galaxies (Schweizer \\& Seitzer 1992\\markcite{ss92}) in a universe where structure is built via hierarchical clustering also predict that the CM relation is driven primarily by metallicity effects (Kauffmann \\& Charlot 1998\\markcite{kauf98}). Moreover, age effects could be important if, for example, there is sufficiently strong feedback from early galaxy formation to bias the luminous mass distribution of subsequent generations of galaxies by the heating of intergalactic gas.  In this paper we present a new, more model-independent approach for evaluating the full range of ages and metallicities allowed by the spectro-photometric properties of early-type galaxies in clusters. The method is based on the construction of age-metallicity diagrams constrained by the colors of early-type galaxies in the nearby Coma cluster and in 17 clusters observed with the {\\it Hubble Space Telescope} ({\\it HST}) at redshifts up to $z\\approx0.9$ (Stanford et al. 1998\\markcite{sed98}). Such an analysis has hitherto been hindered because of the lack of both accurate stellar libraries for different metallicities and reliable morphological information on cluster galaxies at medium-to-high redshifts. Our results can subsequently be reframed into specific theories of galaxy formation, since they will be indispensable for any model that seeks to produce galaxies resembling those actually observed. In \\S2 we present the spectral evolution models used in this paper. The cluster sample is described in \\S3. In \\S4 we construct the age-metallicity diagrams allowed by the observations, and in \\S5 we compute the corresponding ranges in mass-to-light ratio and in several commonly used spectral indices. We discuss our main conclusions in \\S6. ",
        "conclusions": "We have shown that the tight photometric constraints on early-type galaxies in clusters allow relatively wide ranges of ages and metallicities for the dominant stellar populations. In particular, the small scatter of the CM relation out to redshifts $z\\sim1$ does not necessarily imply a common epoch of star formation for all early-type galaxies. It requires, however, that galaxies assembling more recently be on average more metal-rich than older galaxies of similar luminosity. In this context it is interesting to mention that, based on the spectral indices of nearby E/S0 galaxies, Worthey, Trager \\& Faber (1996\\markcite{wo96}) favor younger ages for more metal-rich galaxies than for metal-poor ones at fixed velocity dispersion. The results of our unbiased analysis therefore define the boundaries in age and metallicity that must be satisfied by theoretical studies aimed at explaining the formation and evolution of early-type galaxies in clusters.  The constraints obtained here on the age and metallicity ranges of E/S0 galaxies are consistent with conventional models in which the galaxies all form monolithically in a single giant burst of star formation at high redshift (e.g., Kodama et al. 1998\\markcite{ko98}, and references therein). In fact, this implies that regardless of the true ages and metallicities of early-type galaxies within the allowed range, their photometric properties will always be consistent with {\\it apparently} passive evolution of the stellar populations. As Figure~6 shows, this consistency even extends to  spectral index strengths.  Our results are also consistent with scenarios in which E/S0 galaxies are formed by the merging of disk galaxies (Schweizer \\& Seitzer 1992\\markcite{ss92}) in a universe where structure is built through hierarchical clustering (Kauffmann 1996\\markcite{kauf96}; Baugh, Cole \\& Frenk 1996\\markcite{ba96}; Kauffmann \\& Charlot 1998\\markcite{kauf98}). For such scenarios, Figure~3 constrains the metallicity and epoch of the last major event of star formation in E/S0 galaxies and their progenitors (see \\S2 and \\S4). The ages and metallicities of cluster E/S0 galaxies predicted by hierarchical models are found to be consistent with these constraints (Kauffmann \\& Charlot 1998\\markcite{kauf98}).  To better assess the origin of E/S0 galaxies in clusters one therefore needs to appeal to observational constraints other than their spectro-photometric properties. For example, conventional models of E/S0 galaxy formation are being challenged by the paucity of red galaxies found at high redshifts in deep surveys (Kauffmann, Charlot, \\& White 1997\\markcite{kcw97}; Zepf 1997\\markcite{zepf97}). Also, morphological distinction between E and S0 galaxies and the evolution of the morphology-density relation out to moderate redshifts appear to point to different formation epochs for E and S0 galaxies (Dressler et al. 1997\\markcite{dres97}).  The tightness of the CM relation is proof of a stable process in the assembly of cluster early-type galaxies. However, as we move towards greater redshifts, a drastic change is expected at lookback times that approach the formation of the first E/S0 galaxies. This change can arise as a systematic blueing, an increased scatter or a slope flattening in the CM relation (e.g., Arag\\'on-Salamanca et al. 1993; Charlot \\& Silk 1994\\markcite{cs94}; Kauffmann \\& Charlot 1998\\markcite{kauf98}). An interesting question is raised by the presence of morphologically-selected early-type galaxies with very blue colors in clusters at moderate redshifts (\\S3 and \\S4). If these objects are true cluster members, our analysis shows that they could be young metal-poor galaxies that will later join the CM relation. Hence, we need to probe deeper down the galaxy luminosity function in distant clusters in order to assess whether these objects can have any fundamental bearing on the origin of early-type galaxies."
    },
    "9803/astro-ph9803003_arXiv.txt": {
        "abstract": "We evaluate the effect of screening by bound electron in ${^7}$Be(p,$\\gamma$)$^8$B reaction, where $^7$Be target contains bound electron, in the framework of the adiabatic representation of the three particle problem. A comparison with two other approximations (united atom and folding) is presented. A good agreement between the ``united atom'' approximation and the exact solution is found. We also discuss the screening corrections induced by two K-shell electrons on a $^7$Be target. The bound electron screening effect consequences for $^7$Be and $^8$B solar neutrino fluxes are discussed. ",
        "introduction": "In recent years, an increasing attention has been devoted to an accurate estimation of electron screening effect for nuclear fusion reactions in stellar plasma and for the interactions of low-energy ion beams with atomic or molecular targets in laboratory experiments (see refs.~\\cite{Gruzinov,Shoppa,Salpeter,Mitler,Carraro,Brown,Langanke,Shaviv} and references therein).  In this Letter we present the first quantum mechanical calculation of screening effect by bound electron in \\begin{equation} \\label{creation} {}^7\\mathrm{Be}+\\mathrm{p} \\longrightarrow {}^8\\mathrm{B}+\\gamma \\end{equation} nuclear fusion from the pp-cycle in the Sun. Contribution of this reaction into the the total luminosity of the Sun is negligible small, but it is directly related to the long-standing ``Solar Neutrino Problem'', --  one of the most intriguing issue in the present-day neutrino astrophysics. Standard physics cannot explain an $^{37}$Ar production rate in the Chlorine experiment smaller than that expected from the solar $^8$B neutrino flux measured by both Kamiokande and (with better statistic) Super Kamiokande. GALLEX and SAGE experiments also indicate beryllium neutrino deficit (see the discussion in ref.~\\cite{Innocenti}). One of the most elegant solutions to the solar neutrino anomaly is resonant neutrino flavor conversion in the sun, that is the so-called MSW effect~\\cite{MSW}. It requires an extension of the minimal standard electroweak theory: neutrino masses and neutrino mixing. These neutrino oscillation parameters are determined in a way that can bridge between the predictions of the standard solar models and the solar neutrino observations.  Thus, even in the framework of the standard solar model within a hypothesis of neutrino oscillation (and MSW effect), it is apparently needed more precise calculations of nuclear fusion rates in the sun, because they can significantly affect the neutrino oscillation parameters determination.  A careful study of electron screening effect on nuclear fusion rates becomes particularly actual in view of expected high accuracy neutrino flux measurements by a number of new large detectors (Super Kamiokande, SNO). The interpretation of forthcoming data  requires relevant precise calculations of  solar neutrino fluxes and neutrino energy spectrum.  Usually, the effects of surrounding plasma on the nuclear fusion are treated in electrostatic screening approximation. This approximation, being classical or quantum, correctly reflects the major properties of a process only for high relative velocity of the colliding nuclei, when electron density in the vicinity of the fusing nuclei remains almost unchanged during the collision. In the case, when relative velocity of the nuclei is much smaller or comparable with the electron one (and this is the case at solar conditions), the electron density changes following any relative configuration of the nuclei, and the electrons have an impact on a kinetic energy shift of the nuclei. It is therefore natural to consider the phenomenon within the framework of the adiabatic approximation, which comes from the well known Born-Oppenheimer (BO) approximation.  The BO approximation allows one to treat nuclear motion independently from the electron coordinate, within a new effective potential which depends on the internuclear distance. Since, the considered nuclear velocities are smaller than the boun electron one, corrections to the BO approximation are expected to be negligible. Obviously, the fusing nuclei are from the continuum energy spectrum. An accurate treatment within the adiabatic approximation of the screening effect by electrons from continuum spectrum requires an additional research but the case of bound electrons presents no special problem (see, for example ref.~\\cite{Melezhik}).  As it was argued by A.~Dar, G.~Shaviv, and N.~Shaviv~\\cite{Shaviv,Dar}, the commonly accepted Debye-H\\\"uckel theory is not quite adequate for evaluating the screening effect in not-very-dense stars, like the Sun. There is, also, an experimental evidence that this theory does not provide correct answer for the screening~\\cite{Shoppa}. Actually, this fact is of no importance when the screening due to the plasma electrons is by itself rather small. But it is not the case for the low-lying bound electrons which do screen the electric charge of nuclei much effectively, and moreover, the screening effect drastically increases when energy of the fusing ions decreases. The electron screening can have dramatic effects in very dense stellar cores.  At low and moderate energies, the fusion cross section of ``bare'' charged nuclei colliding with the relative momenta $p$ in the center-of-mass frame is expressed as (see ref.~\\cite{Lang}): \\begin{equation} \\label{sigma} \\sigma_b(E) = \\frac{S(E)}{E}e^{-2\\pi\\eta}, \\end{equation} where $S(E)$ is the so-called astrophysical factor which incorporates all nuclear features of the process, $E$ is the collision energy of the nuclei, $\\eta = MZ_1Z_2/(m_e a_0 p)$ is the usual Coulomb parameter, $m_e$ and $M$ are the electron and reduced nuclear masses, respectively, and $a_0$ is the hydrogen Bohr radius. The exponential factor originates from the Coulomb wave function of the internuclear motion $\\psi_E^{\\mathrm C}(R)$ at $R = 0$.  The screened cross section $\\sigma_s(E)$ differs by the enhancement factor \\begin{equation} \\label{gamma} \\gamma(E)\\equiv \\frac{\\sigma_s}{\\sigma_b} = \\frac{|\\psi_E(0)|^2}{|\\psi_E^{\\mathrm C}(0)|^2}, \\end{equation} where $\\psi_E(R)$ is the wave function of the internuclear motion which accounts for the bound electron.  We evaluate the effect of electron screening of $^7$Be nucleus by one bound electron in reaction (\\ref{creation}) in the framework of the adiabatic approximation for three particle problem. This calculation is compared with two relevant approximations, ``united atom'' (UA) and folding approximations which, as we will demonstrate below, give respectively upper and lower estimates for the screening effect. In the framework of the UA approximation we estimate also the screening effect for $^7$Be nucleus with two K-shell electrons. ",
        "conclusions": "The enhancement factor (\\ref{gamma}) is plotted in fig.~\\ref{enhan} for all three approximations. As it was expected, the UA approximation always overestimate the exact solution, while the folding approximation underestimates it. Nevertheless, it is easy to see that simple UA prescription gives very close values to the exact solution at kinetic energies above $2$ keV. Therefore, the latter can be used  not only as a qualitative, but as a good quantitative approximation to the electron screening by bound electrons.  The electron screening is dramatic at very low kinetic energies of the nuclei. However, in a plasma, most of the nuclear fusions come at the Gamow peak energy, that is defined by the strong dependence of the nuclear cross section on energy (\\ref{sigma}) and the fast decrease of the exponential particle distribution. This energy is given by: \\[ E_0 = 1.22(Z_1Z_2T_6)^{2/3}(M/M_p)^{1/3}\\;{\\rm keV}, \\] where $T_6 = T/10^6$ K, $M_p$ is the proton mass. In the solar interior at $r_{\\mathrm eff}/R_\\odot = 0.06$, where the $^7$Be and $^8$B neutrino production reaches its maximum~\\cite{BahPin}, the plasma parameters are $T_6 \\approx 14.7$, the electron density $n_e \\approx 7.7/a_{0}^3$, and the Gamow peak energy in reaction (\\ref{creation}) is about $18$ keV. Then $\\gamma(E_0) = 1.1$, that is, there is 10\\% of an enhancement by one bound electron in the bohron production rate. Simple computations within the UA approximation give the screening effect as: \\begin{equation} \\label{screen} \\gamma(E_0) = e^{\\Delta E/kT}. \\end{equation}  Then, 10\\% of an enhancement by one bound electron could be easily reproduced just inserting numbers into the formula (\\ref{screen}). One can apply this formula also for a $^7$Be nucleus with two bound electrons. Then, \\[ \\Delta E = \\left(\\chi_1^{\\mathrm{B }}+\\chi_2^{\\mathrm{B }}\\right) - \\left(\\chi_1^{\\mathrm{Be}}+\\chi_2^{\\mathrm{Be}}\\right) = 227.98 \\: \\mathrm{eV}. \\] Here, $\\chi_1^{\\mathrm{B}} = 340.2$~eV and $\\chi_2^{\\mathrm{B}} = 259.4$~eV are, respectively, the fifth and forth ionization potential of the $^8$B atom and $\\chi_1^{\\mathrm{Be}} = 217.72$~eV and $\\chi_2^{\\mathrm{Be}} = 159.9$~eV are, respectively, the forth and third ionization potential of the $^7$Be atom. Thus, two bound electrons enhancement factor is given by $\\gamma(E_0) = 1.196$, i.e. roughly 20\\%.  Using the Saha equation Iben, Kalata and Schwartz~\\cite{Iben} calculated the probabilities $f_1$ and $f_2$ that one or two K-shell electrons are associated with any given $^7$Be nucleus. The calculations were perfomed under the assumption of pure Coulomb electron-ion forces, neglecting all excited states and screening. The probabilites found are \\begin{eqnarray*} f_1 &=& \\lambda\\left[1+\\lambda+0.25\\lambda^2 \\exp{\\left(-\\frac{\\Delta_\\chi}{kT}\\right)}\\right]^{-1}, \\\\ f_2 &=& 0.25\\lambda\\exp{\\left(-\\frac{\\Delta_\\chi}{kT}\\right)}f_1, \\end{eqnarray*} where \\[ \\lambda = n_e\\left(\\frac{h^2}{2\\pi m_e kT}\\right)^{3/2} \\exp{\\left(\\frac{\\chi_1^{\\mathrm Be}}{kT}\\right)}. \\] Here $k$ is the  Boltzmann's constant, and $\\Delta_{\\chi}=\\chi_1^{\\mathrm{Be}}-\\chi_2^{\\mathrm{Be}}=63.8$~eV. Inserting numbers one can obtain: $f_1 = 30\\%$, and $f_2 = 3\\%$.  Using the calculated abundances of $^7$Be ions, one can estimate the thermal averaged screening effect induced by both one and two bound electrons on a $^7$Be nucleus: \\[ \\langle\\gamma\\rangle-1 = 0.30 \\times 0.1 + 0.03 \\times 0.2 \\approx 0.04. \\] In the Standard Solar  Model (SSM) the electron capture rate by $^7$Be nucleus is taken to be about 1000 times faster than the proton capture rate~\\cite{BahPin}. Therefore, a small change in $^8$B production rate does not affect significantly the $^7$Be neutrino flux, although it makes a proportional change in $^8$B neutrino flux. Thus, the electron screening by bound electrons has the prompt consequences on bohron neutrino flux.  The electron screening by plasma electrons from the continuum spectrum is expected to contribute significantly to the total enhancement factor, since it is proportional to $n_e$, and thus it has to be taken into account in the exact prediction of neutrino flux change.  In summary, in the solar interior K-shell bound electrons enhance $^7$Be(p,$\\gamma$)$^8$B rate and increase $^8$B neutrino production rate by of about 4\\%. Therefore, bound electron screening has an effect on the solar $^8$B neutrinos, and acts with the opposite effect to the berrylium neutrinos. The main essence of the electron screening in nuclear fusions is the change in electron density on a nucleus during the collision of the nuclei. This effect can be treated only in a dynamical calculation like the present three particle calculation or the discussed UA approximation.  We thank  J.~N.~Bahcall, A.~V.~Gruzinov and V.~A.~Naumov for usefull discussions."
    },
    "9803/astro-ph9803145_arXiv.txt": {
        "abstract": "We report on the results of a multi-wavelength campaign to observe the soft X-ray transient (SXT) and superluminal jet source \\novasco\\ in outburst using \\HST, \\RXTE\\ and \\CGRO\\ together with ground based facilities.  This outburst was qualitatively quite different to other SXT outbursts and to previous outbursts of this source.  The onset of hard X-ray activity occurred very slowly, over several months and was delayed relative to the soft X-ray rise.  During this period, the optical fluxes {\\em declined} steadily.  This apparent anti-correlation is not consistent with the standard disc instability model of SXT outbursts, nor is it expected if the optical output is dominated by reprocessed X-rays, as in persistent low mass X-ray binaries.  Based on the strength of the 2175\\,\\AA\\ interstellar absorption feature we constrain the reddening to be $\\EBV=1.2\\pm0.1$, a result which is consistent with the known properties of the source and with the strength of interstellar absorption lines.  Using this result we find that our dereddened spectra are dominated by a component peaking in the optical with the expected $\\nu^{1/3}$ disc spectrum seen only in the UV.  We consider possible interpretations of this spectrum in terms of thermal emission from the outer accretion disc and/or secondary star, both with and without X-ray irradiation, and also as non-thermal optical synchrotron emission from a compact self-absorbed central source.  In addition to the prominent \\HeII\\ 4686\\,\\AA\\ line, we see Bowen fluorescence lines of \\NIII\\ and \\OIII, and possible P~Cygni profiles in the UV resonance lines, which can be interpreted in terms of an accretion disc wind.  The X-ray spectra broadly resemble the high-soft state commonly seen in black hole candidates, but evolve through two substates.  Taken as a whole, the outburst dataset cannot readily be interpreted by any standard model for SXT outbursts.  We suggest that many of the characteristics could be interpreted in the context of a model combining X-ray irradiation with the limit cycle disc instability, but with the added ingredient of a very large disc in this long period system. ",
        "introduction": "Soft X-ray transients (SXTs), also referred to as X-ray novae, \\cite{TS96} are a class of low-mass X-ray binaries (LMXBs) in which long periods of quiescence, typically decades, are punctuated by very dramatic X-ray and optical outbursts, often accompanied by radio activity as well.  The most promising models for explaining the outbursts invoke the thermal-viscous limit cycle instability previously developed for cataclysmic variables \\cite{C93}.  These have met with some success in explaining the properties of the outbursts \\cite{CCL95} but there remain difficulties (e.g.\\ Lasota, Narayan \\& Yi 1996).  Compared to cataclysmic variables, an important effect that must be included in models of SXTs is X-ray irradiation of the disc and/or the secondary star.  Irradiation of the disc will change its temperature structure \\cite{TMW90} and may induce delayed reflares (Chen, Livio \\& Gehrels 1993, Mineshige 1994).  The SXT \\novasco\\ was discovered in 1994 July when \\GRO\\ Burst and Transient Source Experiment (BATSE) observed it in outburst at a level of 1.1\\,Crab in the 20--200\\,keV energy band \\cite{Ha95}.  Since then it has undergone repeated outbursts to a similar level and shown itself to be a very atypical SXT.  The outburst history from 1994--5 has been summarised by Tavani et al.\\ \\shortcite{T96}, who draw attention to the contrast between the 1994 outbursts which were radio-loud with apparent superluminal jets observed (Tingay et al.\\ 1995, Hjellming \\& Rupen 1995) and the 1995 outbursts at a similar X-ray flux as in 1994, but radio-quiet.  The optical flux from \\novasco\\ is not as well documented, but Orosz, Schaefer \\& Barnes \\shortcite{OSB95} note that optical brightening does not always accompany X-ray outbursts.  After a period of apparent quiescence from late 1995 to early 1996, \\novasco\\ went into outburst again in late 1996 April \\cite{R96}. Orosz et al.\\ \\shortcite{O97} observed an optical rise leading the X-ray rise detected by \\XTE\\ by about 6~days. They suggested that this initial behaviour was consistent with the limit-cycle instability. The subsequent X-ray behaviour, however, was not as expected.  The soft X-ray flux (2--10\\,keV), as followed by the \\XTE\\ All Sky Monitor (ASM) remained at an approximately constant level for more than 4~months, though with considerable short term variability while the hard X-ray flux (20--200\\,keV) as monitored by \\GRO\\ BATSE was observed to rise very slowly, not reaching its peak until 4~months after the initial dramatic increase in the soft flux.  During this period we carried out a series of simultaneous \\HST\\ and \\XTE\\ visits, backed up by ground based observations and \\GRO\\ BATSE data.  We present here our spectral analysis.  A subsequent paper will address timing issues.  First we summarise the current state of knowledge on the properties of \\novasco: the context in which we interpret our results. \\subsection{System parameters} \\label{ParameterSection} The system parameters of \\novasco\\ are the best known of any SXT. They are summarised in Table~\\ref{ParameterTable}.  Hjellming \\& Rupen \\shortcite{HR95} estimate the distance from a kinematic model of the jets to be $3.2\\pm0.2$\\,kpc.  We also have a lower limit from observations of the 1420\\,MHz interstellar absorption \\cite{T95} of 3.0\\,kpc and an upper limit of 3.5\\,kpc obtained by the method of Mirabel and Rodr\\'{\\i}guez \\shortcite{MR94}.  The latter assumes that we can correctly identify the proper motions of the two jets relative to the central source and then only involves the requirement that these proper motions are produced by material moving at no more than the speed of light.  These two constraints support the distance estimate of Hjellming \\& Rupen.  Other parameters are taken from Orosz \\& Bailyn \\shortcite{OB97} who model the quiescent light curve at a time when the disc is estimated to contribute less than 10 per cent of the V band light.  Their deduced mass of $7.02\\pm0.24$\\,M$_{\\sun}$ makes it clear that the compact object in this system is a black hole. We note that an independent parameter determination by van der Hooft et al.\\ \\shortcite{vdH97} yields values consistent with those of Orosz \\& Bailyn \\shortcite{OB97}, although with larger uncertainties. \\begin{table} \\caption{Adopted parameters for \\novasco.} \\label{ParameterTable} \\begin{center} \\begin{tabular}{lc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Distance              & $3.2\\pm0.2$\\,kpc                    \\\\ Period                & $2.62157\\pm0.00015$\\,d              \\\\ Mass function         & $3.24\\pm0.09$                       \\\\ Inclination           & $69\\fdg50\\pm0\\fdg08$      \\\\ Mass ratio            & $2.99\\pm0.08$                       \\\\ Primary mass          & $7.02\\pm0.22$\\,M$_{\\sun}$          \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We have obtained a series of co-ordinated optical, UV and X-ray spectra spanning several months of the outburst of an SXT.  Although the optical light curve shows the expected decline, it has a spectrum different to that expected on theoretical grounds.  Conversely, although the X-ray spectra showed the familiar high state form, the X-ray fluxes continued to rise through the optical decline, contrary to expectations.  We have considered various interpretations of the observations and suggest the following possible scenario:  The outburst was triggered by a heating wave in the disc, causing an optical rise as the outer disc enters the hot state and a soft X-ray rise subsequently when the material starts to reach the inner disc. Inflow to the inner regions continues to rise for some time as the disc tries to find a steady state.  About a month after the initial rise the accretion mode near the black hole changes.  This results in a rise in the hard X-ray activity and the X-ray variability as the extended hard power-law component becomes prominent.  This change is accompanied by a brief radio flare.  While the X-ray activity is still rising, the disc itself is changing in thickness and/or geometry so that irradiation becomes less efficient allowing the cooling wave to move inwards producing the drop in optical flux at a nearly fixed temperature.  There is also some irradiation of the secondary star, producing an orbital variation in the continuum fluxes and \\HeII\\ 4686\\,\\AA\\ emission.  There clearly remain important unanswered questions not only about \\novasco, but about the outbursts of SXTs in general.  There are many theoretical avenues to be explored in seeking an explanation of these observations, especially in the modelling of long period, large disc systems and the exploration of non-thermal models for the optical emission.  This work also has many useful lessons for the observer.  We have demonstrated the value of co-ordinated, multi-wavelength campaigns in ruling out interpretations which might be suggested by a part of the dataset, but are inconsistent with the whole.  We suggest the following priorities for future observations of SXT outbursts:  \\begin{enumerate} \\item UV observations are {\\em crucial} to such a campaign for the following reasons.  a) It is only in this region that we may be seeing the expected $\\nu^{1/3}$ characteristic accretion disc spectrum in this dataset; identification of this is an important indicator of the disc temperature distribution.  b) It is in the UV resonance lines that we see the signature of an accretion disc wind. Higher resolution, higher signal to noise observations (possible only for a less extremely reddened source) will test models of accretion disc winds and allow an estimate of the mass loss rate. c) The 2175\\,\\AA\\ interstellar absorption feature is our best tool in estimating the reddening of these typically highly reddened objects; other measures such as Na~D-lines are not always reliable in these cases.  Without a good estimate of the interstellar reddening we cannot determine and hence interpret the intrinsic spectrum. \\item The campaign should include spectra, or at least multi-colour photometry, spanning and adequately sampling several consecutive orbits.  This will allow us to separate orbital spectral modulations from random variability and distinguish between emission from an irradiated secondary star and the accretion disc.  In the event of the discovery of an unambiguously eclipsing SXT these observations would be of central importance, allowing eclipse mapping of the emission regions. \\item Comprehensive X-ray spectra spanning as wide an energy range as possible should be an integral part of the campaign.  We observe the high energy side of a thermal component, but lower energy coverage is required to accurately characterise this component and distinguish between a single temperature blackbody and a multi-colour disc. \\item The campaign should include good red and near infrared coverage to obtain improved characterisation of the long wavelength turnover in the spectrum.  This will help in discriminating between thermal and non-thermal emission, which show different long-wavelength limits, and in the thermal case will provide more comprehensive information on the cooler parts of the system. \\end{enumerate}"
    },
    "9803/astro-ph9803102_arXiv.txt": {
        "abstract": "We present near-infrared (observed frame) spectra of the high-redshift quasar S4{\\ts}0636+68 at $z=3.2$ which was previously thought to be one of a group of ``strong'' \\ion{Fe}{2} emitters (i.e., $F(\\mbox{\\ion{Fe}{2}}{\\ts} \\lambda\\lambda\\mbox{4434--4684})/F({\\rm H}\\beta) > 1$). Our {\\it K}-band spectrum clearly shows emission lines of H$\\beta$ and [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007 as well as optical \\ion{Fe}{2} emission.  Our computed value of $F(\\mbox{\\ion{Fe}{2} } \\lambda\\lambda\\mbox{4434--4684})/F({\\rm H}\\beta) \\simeq 0.8$ for S4{\\ts}0636+68 is less than previously thought, and in fact is comparable to values found for radio-loud, flat-spectrum, low-$z$ quasars. Therefore S4{\\ts}0636+68 appears not to be a strong optical \\ion{Fe}{2} emitter.  Although more than half (5/8) of the high-$z$ quasars observed to date are still classified as strong optical \\ion{Fe}{2} emitters, their \\ion{Fe}{2}/H$\\beta$ ratios, for the most part, follow the same trend as that of low-$z$ quasars, i.e., an anticorrelation in $EW$(\\ion{Fe}{2})/$EW$(H$\\beta$) versus $EW$([\\ion{O}{3}])/$EW$(H$\\beta$), with radio-loud quasars having a mean value of $EW$(\\ion{Fe}{2})/$EW$(H$\\beta$) approximately half that of radio-quiet quasars at comparable values of $EW$([\\ion{O}{3}])/$EW$(H$\\beta$). ",
        "introduction": "Since optical \\ion{Fe}{2} emission\\footnote{% The \\ion{Fe}{2} emission feature actually extends from the near-UV into the red optical region of the spectrum. However, following previous convention, we use the term ``optical \\ion{Fe}{2} emission'' in this paper to mean the \\ion{Fe}{2} emission near H$\\beta$ (i.e., \\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684).} is often one of the prominent features in the spectra of Type 1 active galactic nuclei (AGN), it is perhaps not surprising that several  observational and theoretical studies have been made to explain the strength of this feature in quasars (e.g. Phillips \\markcite{Phillips77}1977, \\markcite{Phillips78}1978; \\markcite{Kwan81}Kwan \\& Krolik 1981; \\markcite{Netzer83}Netzer \\& Wills 1983; \\markcite{Wills85}Wills {\\it et al.}\\ 1985; \\markcite{Collin-Souffrin88}Collin-Souffrin {\\it et al.}\\ 1988; \\markcite{Zheng90}Zheng \\& O'Brien 1990; \\markcite{Joly91}Joly 1991; \\markcite{Boroson92}Boroson \\& Green 1992; \\markcite{Lipari93}L\\'{\\i}pari {\\it et al.}\\  1993; \\markcite{Wang96a}Wang {\\it et al.}\\ 1996a). Although it is known that the strength of the optical \\ion{Fe}{2}  emission shows an anticorrelation with the strength of [\\ion{O}{3}] emission (Boroson \\& Green 1992), the physical properties of the \\ion{Fe}{2}  emitting region are not yet fully understood.  Recent near-infrared (NIR) spectroscopic studies by \\markcite{Hill93}Hill {\\it et al.}\\ (1993) and \\markcite{Elston94}Elston {\\it et al.}\\ (1994; hereafter ETH) suggest that unusually strong optical \\ion{Fe}{2} emitters may be common in the high-$z$ universe ($2 < z < 3.4$). Though it is known that some low-$z$ far-infrared (FIR) selected  AGN ($L_{\\rm FIR} \\gtrsim 10^{11}$ $L_{\\sun}$) show strong \\ion{Fe}{2} emission in their optical spectra \\markcite{Lipari93}(cf.\\ L\\'{\\i}pari {\\it et al.}\\ 1993), such extreme \\ion{Fe}{2} emitters appear to be rare in the low-$z$ universe. Recently, we obtained NIR spectra of two radio-loud, flat-spectrum, high-$z$ quasars (B 1422+231 at $z=3.6$ and PKS 1937$-$101 at $z=3.8$) and found that their flux ratios of $F(\\mbox{\\ion{Fe}{2} }{\\ts}\\lambda\\lambda\\mbox{4434--4684})/F({\\rm H}\\beta)$ are much less than those of the other high-$z$ quasars (\\markcite{Kawara96}Kawara {\\it et al}.\\ 1996; Taniguchi {\\it et al.}\\ \\markcite{Taniguchi96}1996, \\markcite{Taniguchi97}1997), and in fact are similar to those of radio-loud, flat-spectrum, low-$z$ quasars with normal optical \\ion{Fe}{2} emission.  These new observations suggest that high-$z$ quasars may exhibit a range of values of $F(\\mbox{\\ion{Fe}{2} }{\\ts}\\lambda\\lambda\\mbox{4434--4684})/F({\\rm H}\\beta)$ similar to what has been observed for low-$z$ quasars.   If the strong \\ion{Fe}{2} emission could be attributed to the overabundance of iron, host galaxies of the high-$z$ quasars with strong \\ion{Fe}{2} emission would form at $z \\sim${\\ts}10 because it is usually considered to be the case that the bulk of the iron arises from Type Ia supernovae which occur $\\sim${\\ts}1--2{\\ts}Gyr after the first major epoch of star formation (e.g., \\markcite{Hamann93}Hamann \\& Ferland 1993, \\markcite{Yoshii96}Yoshii {\\it et al}.\\ 1996). It is therefore important to investigate the chemical properties of high-$z$ quasars systematically.  In this paper we present new NIR spectroscopy of S4{\\ts}0636+68, a flat-spectrum radio-loud quasar at $z=3.2$, which is reported in ETH as a very strong iron emitter. Based on our new measurements, we discuss whether the fraction of high-$z$ quasars with strong optical \\ion{Fe}{2} emission is substantially higher than that of low-$z$ quasars. ",
        "conclusions": "\\subsection{The Rest-Frame UV and Optical Spectra of S4{\\ts}0636+68} Figure \\ref{fig-1} shows the spectra of S4{\\ts}0636+68 (solid line) in the $I\\!H\\!K$ bands together with the Large Bright Quasar Survey (LBQS) composite spectrum (dashed line; \\markcite{Francis91} Francis {\\it et al}.\\ 1991) shifted to $z=3.2$. The atmospheric transmission of Mauna Kea is shown in the upper panel. The spectrum clearly shows H$\\beta$ at 2.04 \\micron{} and a broad ``bump'' of \\ion{Fe}{2} emission between 2.15{\\ts}\\micron{} and 2.23{\\ts}\\micron{}. The spike feature in our spectrum marked by `X' is caused by residual atmospheric absorption. Although ETH did not find evidence for [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007 emission lines, our $K$-band spectrum shows their presence (which will be discussed later). No prominent emission lines were found in the $I$-, $J$-, and $H$-band regions of the spectrum. \\ion{Mg}{2}{\\ts}$\\lambda$2798 would be expected to appear at $\\sim$ 1.18 \\micron{} but the $J$-band spectrum (not shown in Figure \\ref{fig-1}) was too noisy to be able to study \\ion{Mg}{2}. Since the efficiency of KSPEC in the $J$-band is not high, we do not use the $J$-band data in this paper.  \\placefigure{fig-1}  In Figure 2, we compare our results with previous optical-NIR spectroscopic studies of S4{\\ts}0636+68 (\\markcite{Sargent89}Sargent {\\it et al}.\\ 1989; \\markcite{Bechtold94}Bechtold {\\it et al}.\\ 1994; ETH). Their basic data are summarized in Table \\ref{tbl-1}. Our $K$-band spectrum is twice as bright as that of ETH. On the other hand, our $I$-band spectrum is 40{\\ts}\\% fainter than that of \\markcite{Bechtold94}Bechtold {\\it et al}.\\ (1994). These flux discrepancies may be due to possible time variation inherent in the object or to calibration errors in the absolute photometry. However, since the two optical spectra taken at different observing dates over two years are quite consistent with each other (\\markcite{Sargent89}Sargent {\\it et al}.\\ 1989; \\markcite{Bechtold94}Bechtold {\\it et al}.\\ 1994), it seems unlikely that this quasar is highly variable. Therefore, we consider the possibility that the discrepancy may be mostly due to calibration errors. First, we note that our $I\\!H\\!K$ spectra were taken simultaneously and thus there is no internal calibration error in our spectra. Second, the optical spectra of \\markcite{Sargent89}Sargent {\\it et al}.\\ (1989) and \\markcite{Bechtold94}Bechtold {\\it et al}.\\ (1994) show good agreement with each other and thus their photometric calibration seems reliable. Further, the optical power-law slope, $\\alpha=-0.68$ ($f_\\nu \\propto \\nu^\\alpha$), given in \\markcite{Sargent89}Sargent {\\it et al}.\\ (1989) can be consistently extrapolated onto our $K$-band spectrum; the spectral index using the emission-free regions in both the optical spectrum (\\markcite{Sargent89}Sargent {\\it et al}.\\ 1989) and our $K$-band spectrum (1330--1380 \\AA, 1430--1460 \\AA, and 5400--5850 \\AA{} in the rest frame) is estimated to be $\\alpha=-0.69$. Since the seeing during our observations was $\\simeq 0\\farcs{}5$ (FWHM) and our slit size was 1\\arcsec{}, we think that we have detected nearly all of the light from the quasar. Though the details of the observing conditions and slit size are not given in ETH (see Table \\ref{tbl-1}), seeing conditions on Mauna Kea are often better than at KPNO, judging from our experience at KPNO (see \\markcite{Kawara96}Kawara {\\it et al}.\\ 1996; \\markcite{Kawara97}Taniguchi {\\it et al}.\\ 1997), and, therefore we expect that our new measurement is more reliable.  \\begin{table} \\dummytable\\label{tbl-1} \\end{table} \\placetable{tbl-1} \\placefigure{fig-2}  \\subsection{The Rest-frame Optical Emission-line Properties of S4{\\ts}0636+68}  The main aim of our current observations is to provide a more accurate measure of the optical \\ion{Fe}{2}/H$\\beta$ ratio in S4{\\ts}0636+68. Figure \\ref{fig-3} compares our result with that of the ETH. (Note that the flux of ETH spectrum is scaled by a factor of two for proper comparison.) Center positions of H$\\beta$, [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007, and the \\ion{Fe}{2} multiplet (42) at a redshift  $z=3.2$ are marked in Figure \\ref{fig-3}. The peak positions of H$\\beta$, [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007, and \\ion{Fe}{2}{\\ts}$\\lambda$5169 (one of \\ion{Fe}{2} multiplet 42 lines), coincide between the two spectra. The $K$-band spectrum of ETH appears to be dominated by very strong \\ion{Fe}{2} emission with weak, or nondetected [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007. ETH actually stated that [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007 was not detected, although they noted that their spectrum had small bumps of low significance at the position of the [\\ion{O}{3}] lines. On the other hand, our $K$-band spectrum clearly shows emission peaks which can be identified with [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007.   \\placefigure{fig-3}  To measure the \\ion{Fe}{2} fluxes in our spectrum, we fit emission-line features simultaneously with a least-squares algorithm. Such fitting results depend on the adopted continuum spectrum. As shown in Boroson \\& Green (1992), a local linear continuum is usually adopted to fit H$\\beta$, [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007, and the \\ion{Fe}{2} features. However, we have already obtained a global power-law continuum using the rest-frame UV and optical spectra as shown in Figure 2. Therefore, we performed spectral fitting for two cases; 1) local linear continuum and 2) global power-law continuum. In the fitting procedure, we assumed $F(\\mbox{[\\ion{O}{3}]}{\\ts}\\lambda{\\rm 5007})/F(\\mbox{[\\ion{O}{3}]}{\\ts} \\lambda {\\rm 4959})$ = 2.97 (\\markcite{Osterbrock89}Osterbrock 1989). We also assumed that the emission line profiles of H$\\beta$ and the [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007 doublet are Gaussian. As for the optical \\ion{Fe}{2} emission features, we used an \\ion{Fe}{2} spectrum of a low-$z$ BAL quasar, PG 0043+039 (\\markcite{Turnshek94}Turnshek {\\it et al}.\\ 1994), as our \\ion{Fe}{2} template. All the emission lines are assumed to have the same redshift. Since it is known that high-ionization broad lines (e.g., \\ion{C}{4}{\\ts}$\\lambda$1549) are often blueshifted with respect to low-ionization lines (Gaskell 1982; Wilkes 1984; Carswell {\\it et al}.\\ 1991; Nishihara {\\it et al}.\\ 1997 and references therein), we use only low ionization lines in our analysis.  The fitting results are presented in Figure \\ref{fig-4} and Table \\ref{tbl-2} for each of the two assumed continua.  The difference in the line flux ratios between the two cases is less than the measurement errors. Although we do not know which continuum case is more realistic, we adopt the results using the linear continuum fit for further  discussion in order to compare our results with those of Boroson \\& Green (1992) and Hill {\\it et al}. (1993) since they also adopted a local linear continuum.  In order to examine whether or not the detection of the [\\ion{O}{3}] lines are real in our spectrum, we compare our fit including the [\\ion{O}{3}] doublet with a fit excluding the [\\ion{O}{3}] doublet, where the local linear continuum has been adopted in both fits. A F-statistics test indicates that the fit with [\\ion{O}{3}] is improved over the 4900--5050 \\AA{} region from the fit excluding [\\ion{O}{3}] at a significance level of  99.8{\\ts}\\%. Therefore, we conclude that the ``bumps'' at the [\\ion{O}{3}] positions are really the [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007 doublet rather than \\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4924,{\\ts}5018 of the 42 multiplet. We obtained an average redshift $z=3.200 \\pm 0.002$.   \\begin{table} \\dummytable\\label{tbl-2} \\end{table} \\placetable{tbl-2} \\placefigure{fig-4}  Our fit assuming a local linear continuum gives the flux ratio $F(\\mbox{\\ion{Fe}{2} } \\lambda 5169)/F({\\rm H}\\beta)=0.28$. This value is significantly smaller than the value of 0.45 that we estimate from the published spectrum of ETH. Although we do not understand this difference, it could reasonably be explained by uncertainties in the continuum fits between the two observations. However, we cannot rule out the possibility of time variation. In fact, such a time variation of \\ion{Fe}{2} is reported for the nearby type 1 Seyfert galaxy NGC 5548 (\\markcite{Wamsteker90}Wamsteker {\\it et al}.\\ 1990; \\markcite{Maoz93}Maoz {\\it et al}.\\ 1993; \\markcite{Sergeev97}Sergeev {\\it et al}.\\ 1997). Thus, monitoring of S4{\\ts}0636+68 may be needed in the future.  The flux ratio, $F(\\mbox{\\ion{Fe}{2} } \\lambda\\lambda\\mbox{3500--6000})/F({\\rm H}\\beta)$, for  S4{\\ts}0636+68 is $3.5\\pm1.1$ (Table \\ref{tbl-2}). This value is greater than the mean value of $1.63 \\pm 0.88$ for the six low-$z$ quasars studied by \\markcite{Wills85}Wills {\\it et al}.\\ (1985) and the value of 2.9 for 3C 273 which is the strongest optical \\ion{Fe}{2} quasar in the sample of \\markcite{Wills85}Wills {\\it et al}.\\ (1985). However, $F(\\mbox{\\ion{Fe}{2}}{\\ts}\\lambda\\lambda\\mbox{4434--4684}) /F({\\rm H}\\beta)$ for S4{\\ts}0636+68 is $0.83 \\pm 0.26$ which is only half of the average value of $1.77 \\pm 0.17$ for the four high-$z$ quasars studied by \\markcite{Hill93}Hill {\\it et al}.\\ (1993). \\markcite{Lipari93}L\\'{\\i}pari {\\it et al}.\\ (1993) defined quasars with $F(\\mbox{\\ion{Fe}{2}}{\\ts}\\lambda\\lambda\\mbox{4434--4684}) /F({\\rm H}\\beta) \\gtrsim 1$ as ``strong'' iron emitters. According to this criterion, we conclude that S4{\\ts}0680+68 is not a strong \\ion{Fe}{2} emitter, contrary to the conclusion of ETH.   \\subsection{Statistical Properties of High-$z$ Quasars vs.\\ Low-$z$ Quasars}  In order to assess the significance of our new result for the ratio $F(\\mbox{\\ion{Fe}{2}}{\\ts}\\lambda\\lambda\\mbox{4434--4684}) /F({\\rm H}\\beta)$ in S4{\\ts}0680+68 we first compare the rest-frame optical emission line properties of low-$z$ and high-$z$ quasars. Figure \\ref{fig-5} shows the relationship of equivalent width (EW) ratios between $EW(\\mbox{[\\ion{O}{3}]}{\\ts}\\lambda4959+\\lambda5007) / EW({\\rm H}\\beta)$ and $EW(\\mbox{\\ion{Fe}{2} }{\\ts}\\lambda\\lambda\\mbox{4434--4684}) / EW({\\rm H}\\beta)$ for low-$z$ and high-$z$ quasars compiled from the literature (\\markcite{Boroson92}Boroson \\& Green 1992; \\markcite{Hill93}Hill {\\it et al}.\\ 1993; ETH; \\markcite{Kawara96}Kawara {\\it et al}.\\ 1996; \\markcite{Taniguchi97}Taniguchi {\\it et al}.\\ 1997). There is a distinct anticorrelation for the low-$z$ quasars as noted before (cf.\\ \\markcite{Boroson92}Boroson \\& Green 1992) although the reason for the anticorrelation between [\\ion{O}{3}]{\\ts}$\\lambda\\lambda$4959,{\\ts}5007 and optical \\ion{Fe}{2}{\\ts}$\\lambda\\lambda\\mbox{4434--4684}$ is still unknown. The low-$z$ radio-loud quasars tend to have  small ratios both in $EW(\\mbox{[\\ion{O}{3}]}{\\ts}\\lambda4959+\\lambda5007) /EW({\\rm H}\\beta)$ and $EW(\\mbox{\\ion{Fe}{2}}{\\ts}\\lambda\\lambda\\mbox{4434--4684}) /EW({\\rm H}\\beta)$. Five of the eight high-$z$ quasars show strong \\ion{Fe}{2} (i.e., \\ion{Fe}{2}/H$\\beta > 1$) emission. The remaining three high-$z$ quasars, which have \\ion{Fe}{2}/H$\\beta < 1$, are all radio-loud and lie within the locus of values traced by low-$z$ radio-loud quasars in Figure 5\\ \\footnote{We note that a radio-loud high-$z$ ($z=2.09$) quasar 1331+170 also appears to lie within the locus of values found for the low-$z$ radio-loud quasars (i.e., 1331+170 appears to have ``quite weak'' optical \\ion{Fe}{2} emission and $EW(\\mbox{[\\ion{O}{3}]}{\\ts}\\lambda4959+\\lambda5007)/EW({\\rm H}\\beta) \\sim{\\ts}0.7$ \\ (Carswell {\\it et al.} \\ 1991).}. Also, the three radio-quiet quasars among the five high-$z$ quasars with strong \\ion{Fe}{2} emission appear to lie within the upper envelope of values observed for low-$z$ radio-quiet quasars. In summary, although over half (5/8) of the high-$z$ quasars appear to be by definition strong \\ion{Fe}{2} emitters, all but two (the radio-loud \\ion{Fe}{2} quasars S5{\\ts}0014+81 and B2{\\ts}1225+317: ETH; \\markcite{Hill93}Hill {\\it et al}.\\ 1993) of the high-$z$ quasars follow a similar trend as that shown by the low-$z$ quasars, i.e.\\ an anticorrelation of $EW(\\mbox{\\ion{Fe}{2}}{\\ts}\\lambda\\lambda\\mbox{4434--4684}) /EW({\\rm H}\\beta)$ versus $EW(\\mbox{[\\ion{O}{3}]}{\\ts}\\lambda4959+\\lambda5007) /EW({\\rm H}\\beta)$, with radio-loud quasars having on average smaller values of $EW(\\mbox{\\ion{Fe}{2}}{\\ts}\\lambda\\lambda\\mbox{4434--4684}) /EW({\\rm H}\\beta)$ than radio-quiet quasars at any given value of $EW(\\mbox{[\\ion{O}{3}]}{\\ts}\\lambda4959+\\lambda5007) /EW({\\rm H}\\beta)$.   \\placefigure{fig-5}  Recently Wang {\\it et al}.\\ (\\markcite{Wang96b}1996b) studied the relation between optical \\ion{Fe}{2} strength and properties of the UV spectra for 53 low-$z$ ($z \\lesssim 0.2$) quasars and found that there is a significant anticorrelation between the equivalent widths of optical \\ion{Fe}{2}{\\ts}{\\ts}$\\lambda\\lambda\\mbox{4434--4684}$ and \\ion{C}{4}{\\ts}$\\lambda$1549.  We examine whether the high-$z$ quasars follow this anticorrelation (Table \\ref{tbl-3} and Figure \\ref{fig-6}). It is perhaps expected that the high-$z$ quasars would have smaller {\\it EW}(\\ion{C}{4}{\\ts}$\\lambda$1549) than low-$z$ quasars simply because of the known anticorrelation between {\\it EW}(\\ion{C}{4}{\\ts}$\\lambda$1549) and UV continuum luminosity (Baldwin effect: \\markcite{Baldwin77}Baldwin 1977; \\markcite{Baldwin78}Baldwin {\\it et al}.\\ 1978).  Not as evident perhaps is that, except for 0933+733, the {\\it EW}(\\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684), appears to show the same range of values as do the low-$z$ quasars at comparable low values of {\\it EW}(\\ion{C}{4}{\\ts}$\\lambda$1549). However, five of the remaining seven high-$z$ quasars (B2{\\ts}1225+317, 1246$-$057, S4{\\ts}0636+68, B{\\ts}1422+231, and PKS{\\ts}1937$-$101) have smaller {\\it EW}(\\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684) than any of the low-$z$ quasars with comparable {\\it EW}(\\ion{C}{4}{\\ts}$\\lambda$1549) (see the lower-left region of the diagram in Figure 6), thus, adding the high-$z$ sample to the low-$z$ sample appears to decrease somewhat the significance of the anticorrelation between {\\it EW}(\\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684) versus {\\it EW}(\\ion{C}{4}{\\ts}$\\lambda$1549), (although it is possible that not having a less luminous high-$z$ sample may cause a selection effect).  We thus consider it possible that the \\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684 emitting region may not have a physical link directly with the \\ion{C}{4}{\\ts}$\\lambda$1549 emitting region.  However, it seems clear that there is no object with large EWs in both \\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684 and \\ion{C}{4}{\\ts}$\\lambda$1549, and furthermore, the upper bound of {\\it EW}(\\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684) still decreases with increasing {\\it EW}(\\ion{C}{4}{\\ts}$\\lambda$1549) for the combined high-$z$ and low-$z$ samples.  Hence, it is suggested that there may still be an indirect relation between the \\ion{Fe}{2}{\\ts}$\\lambda\\lambda$4434--4684 and \\ion{C}{4}{\\ts}$\\lambda$1549 regions.   \\begin{table} \\dummytable\\label{tbl-3} \\end{table} \\placetable{tbl-3} \\placefigure{fig-6}  In summary, the relations among the emission-line properties shown in Figures \\ref{fig-5} and \\ref{fig-6} appear to be valid for both the low-$z$ and high-$z$ quasars with only a few exceptions. This implies that the emission mechanism and the physical properties of the emission-line region in high-$z$ quasars may not be significantly different from those in low-$z$ quasars."
    },
    "9803/astro-ph9803334_arXiv.txt": {
        "abstract": "We have analyzed the polarization properties of pulsars at an observing frequency of 4.9 GHz. Together with low frequency data, we are able to trace polarization profiles over more than three octaves into an interesting frequency regime. At those high frequencies the polarization properties often undergo important changes such as significant depolarization. A detailed analysis allowed us to identify parameters, which regulate those changes. A significant correlation was found between the integrated degree of polarization and the loss of rotational energy $\\dot E$.  The data were also used to review the widely established pulsar profile classification scheme of core- and cone-type beams.  We have discovered the existence of pulsars which show a strongly {\\it increasing} degree of circular polarization towards high frequencies.  Previously unpublished average polarization profiles, recorded at the 100m Effelsberg radio telescope, are presented for 32 radio pulsars at 4.9 GHz. The data were used to derive polarimetric parameters and emission heights. ",
        "introduction": "Polarimetry plays a key role in our understanding of the emission mechanism of pulsars, the ambient conditions in the emission region and the geometrical structure of the magnetic field. Similar to the great variety of profile shapes, the polarimetric features of the radio emission vary strongly from pulsar to pulsar and from frequency to frequency. Virtually every polarization state between totally unpolarized and fully linearly or highly circularly polarized can be found amongst different pulsars and even within one profile. Also the shape of the polarization position angle (hereafter PPA) curve varies between nearly constant, a smooth orderly swing, sudden jumps and nearly chaotic behaviour. The jumps in the PPA--swing often cover precisely $90^\\circ$ and are therefore called orthogonal polarization modes (hereafter OPM, see e.g. Stinebring et al. 1984; Gil \\& Lyne 1995; Gangadhara 1997).  Nevertheless certain common pulse features have been identified in the past, which lead to different classification attempts (e.g. Backer 1976; Rankin 1983; Lyne \\& Manchester 1988). It is widely accepted that two general types of profile components can be identified: Those which are radiated from the outer parts of the emission tube as {\\it conal} profile components and those which are usually emitted from the central part as {\\it core}-beams. This classification was initially formulated systematically by Rankin (1983), for a detailed description we refer to that paper, a short summary is given in Sect. \\ref{types}. The identification of these components is mainly (but not only) based on the frequency development of their polarization and their relative intensity. In general this system has proven to be remarkably successful, although in this paper we discuss some groups of pulsars which do not quite fit into this classification scheme.  One common polarimetric feature of most pulsars is the depolarization towards high frequencies (in the following ``high frequency'' means radio frequencies well above one GHz). Whereas the degree of polarization of pulsars is usually constantly high at low frequencies, it decreases rapidly above a certain frequency (e.g. Manchester 1971; Morris et al. 1981a; Xilouris et al. 1996). This tendency is in contrast to the known properties of other astrophysical objects which have usually stronger polarization at higher frequencies, where the Faraday-depolarization effect is less severe. Therefore, this effect is thought to be inherent to the pulsar magnetosphere, either intrinsic to the emission mechanism or due to a propagation effect within the magnetosphere. The identification of the depolarization mechanism is important as it might help to understand the environmental conditions in the magnetosphere and the relevant physical processes. It is therefore necessary to carry out high frequency observations as we would like to identify parameters which control this effect.  In this paper we also focus on the role of the circular polarization. Theories proposed to explain this type of polarization range from purely intrinsic mechanisms (e.g. \\cite{RR90}) to pure propagation effects (e.g. \\cite{MS77}) and combinations of both (e.g. Kazbegi et al. 1991; Naik \\& Kulkarni 1994). In order to distinguish between the different mechanisms, it is necessary to trace the frequency development of the circular polarization over a large frequency interval. Propagational effects should show a strong frequency dependence.  Whereas the degree of polarization varies strongly with frequency, the measured PPA is very stable over many octaves in frequency. If one allows for the occurrence of OPMs, the PPA swing is therefore thought to reflect the geometry of the pulsar magnetosphere as first noted by Radhakrishnan \\& Cooke (1969). In some cases it is therefore possible to determine the viewing geometry of a pulsar by fitting the geometry dependent theoretical PPA curve -- the rotating vector model (hereafter RVM) -- to the measured PPA (the formula for the RVM is given e.g. by Manchester \\& Taylor (1977)).  Many authors indicate the existence of a radius-to-frequency mapping (hereafter RFM) where the radio emission is narrow band and scales inversely with frequency (e.g. Cordes 1978; Blaskiewicz et al. 1991; Kramer et al. 1997; Kijak \\& Gil 1997; von Hoensbroech \\& Xilouris 1997a). The knowledge of the existence and the strength of the RFM is important as it will help to understand the emission physics. It is therefore necessary to determine the emission height above the pulsar surface, where the emission we observe at a certain frequency, originates. Polarimetry provides one method amongst others to calculate this height. This method was proposed by Blaskiewicz et al. (1991) and we have applied it to our data whenever possible (see Sect. \\ref{Rem}).  Due to their steep radio spectra, pulsars tend to be rather weak sources at centimetre wavelengths. As a result, relatively little published data exists in this part of the spectrum (Morris at al. 1981b; Xilouris et al. 1994; Xilouris et al. 1995; Manchester \\& Johnston 1995; Xilouris et al. 1996; von Hoensbroech et al. 1997b). In this paper we present the polarimetric properties of 32 weaker pulsars at 4.9 GHz which roughly doubles the number of published polarization profiles at this frequency and allows statistical studies at such a high frequency for the first time. ",
        "conclusions": "We have analysed average radio pulsar polarization profiles at high frequencies and present 32 previously unpublished pulsar polarization profiles measured at a frequency of 4.9 GHz. The profiles are also available in EPN-format (\\cite{L98}) through the EPN internet database (see Sect. \\ref{discussion}). The properties of the whole available set of 4.9 GHz average polarization profiles in general and those of individual pulsars in special were compared to lower frequencies.  Investigating the average polarization profiles of individual pulsars with particular respect to their classification in the scheme of Rankin (1983), we found groups of pulsars which deserve additional attention. \\begin{itemize}  \\item Some pulsars, such as PSR B0355+54, have components with very different polarimetric and spectral properties within the same profile. One component is nearly fully polarized and has a flatter spectrum than the profile as a whole, thus dominates the profile at high frequencies. The rest of the profile is hardly polarized and dominates only at low frequencies. As all eight pulsars which form this group so far, show a similar frequency-dependence, it is suggested that an intrinsic correlation exists between a high degree of polarization and a flat spectrum. All these pulsars have been classified as half-cones. Although this classification is tempting, it is important to note that not a single full-cone with similar properties could be found. This indicates that a more general process takes place than just an occasional lack of flux at the position where the line of sight cuts the cone for the second time.  \\item There are three young pulsars (B1800-21, B1823-13 and B1259-63) with a very high loss of rotational energy $\\dot E$ and a very flat spectrum. These pulsars are nearly fully polarized and do not show any significant depolarization. This behaviour shows similarities to the above mentioned highly polarized components of the 0355-like pulsars. Both groups show a correlation between high polarization and a flat spectrum.  \\item We found a number of objects which show a circular polarization, which strongly {\\it increases} with frequency. This is in sharp contrast to the known frequency-dependence of pulsar polarization. As the linear polarization decreases simultaneously, it is suggested that a propagation-effect similar to a $\\lambda/4$-plate is active. If confirmed, this could indicate, that propagation effects influence the polarization within the magnetosphere. Additionally, these pulsars fit hardly into the classification scheme. PSR B0144+59 for instance shows precisely the opposite frequency-development to a $S_t$-pulsar (see Fig. \\ref{0144_freq}).  \\end{itemize}  We would like to point out again the possible role of pulsar evolution on the polarization profile. Within the empirical model for pulsar emission, the profile shape and the polarization properties is nearly exclusively determined by the viewing geometry of pulsar and line of sight and the activity in the different parts of the magnetosphere. But, as it was already noted by Rankin (1983), the classical conal double pulsars (the ``textbook-pulsars'', 0525+21-type) are without exception very old stars which lie close to the ``death-line'' in the $P-\\dot P$-diagram. Contrarily the 1800$-$21-like pulsars mentioned above are very young pulsars with very different properties. For future work it appears to be important to focus stronger on the role of evolution for polarization profile shapes.  Analysing the general properties of pulsar polarization profiles at this frequency of 4.9 GHz, we found a significant correlation between the total degree of polarization with $\\dot E$ and the $\\Phi_\\parallel$ at the polar gap respectively  (see Fig. \\ref{PEdot}, upper plot). Such a correlation does not exist at lower frequencies. The pulsars at low frequencies  rather form groups which have a decreasing (highly polarized, low $\\dot E$ pulsars) and an increasing (weakly polarized, high $\\dot E$ pulsars) degree of polarization to high frequencies. Observations at high radio frequencies therefore yield additional information on the emission physics which is not seen at lower frequencies. This correlation confirms the relation between the depolarization index and $\\Phi_\\parallel$ at 10.5 GHz, which was presented by Xilouris et al. (1995).  The large differences in the degree of polarization between individual pulsars indicate that the ambient physical conditions in the respective emitting region differ significantly among them. As we can see from the correlation, this depends on $\\dot E$. $\\dot E$ again is closely correlated to the polar gap $\\Phi_\\parallel$. As a high degree of polarization seems to be correlated to a flatter spectrum (see above), we speculate that a high $\\Phi_\\parallel$ could induce a flatter energy distribution function of the radiating plasma.  \\onecolumn \\begin{figure}[C] \\epsfysize23cm \\epsffile[25 70 525 770]{datn1.ps} \\caption{Pulsar polarization profiles at 4.85 GHz. The dark-shaded area represents the linear, the light-shaded area corresponds to the circularly polarized intensity ({\\it positive} $\\hat =$ left-hand, {\\it negative} $\\hat =$ right-hand). Total power is represented by the unshaded solid line. The error-box has a height of 2 $\\sigma$ and a width corresponding to the effective time-resolution (see caption of Table~1). When it was possible, the RVM was fitted to the angle (e.g. for B0144+59).} \\label{data1} \\end{figure} \\twocolumn  \\onecolumn \\begin{figure}[C] \\epsfysize23cm \\epsffile[25 70 525 770]{datn2.ps} \\caption{Pulsar polarization profiles at 4.85 GHz. For details see caption of Fig. 8.} \\label{data2} \\end{figure} \\twocolumn  \\onecolumn \\begin{figure}[C] \\epsfysize23cm \\epsffile[25 70 525 770]{datn3.ps} \\caption{Pulsar polarization profiles at 4.85 GHz. For details see caption of Fig. 8.} \\label{data3} \\end{figure} \\twocolumn  \\onecolumn \\begin{figure}[T] \\epsfysize11.5cm \\epsffile[25 220 525 570]{datn4.ps} \\vspace*{-2cm} \\caption{Pulsar polarization profiles at 4.85 GHz. For details see caption of Fig. 8.} \\label{data4} \\end{figure} \\twocolumn"
    },
    "9803/astro-ph9803044_arXiv.txt": {
        "abstract": "We have  derived the masses of central objects ($M_{\\rm BH}$) of nine type 2 Seyfert nuclei using the observational properties of the {\\it hidden} broad H$\\beta$ emission. We obtain the average dynamical mass, log$(M_{\\rm BH} / M_\\odot) \\simeq 8.00 \\pm 0.51 - 0.475 {\\rm log}(\\tau_{\\rm es}/1)$ where  $\\tau_{\\rm es}$ is the optical depth for electron scattering. If  $\\tau_{\\rm es} \\sim 1$, this  average mass is almost comparable with those of type 1 Seyfert nuclei. However, if  $\\tau_{\\rm es} \\ll 1$, as is usually considered, the average mass of type 2 Seyfert nuclei may be more massive than that of type 1s. We discuss implications for issues concerning both the current unified model of Seyfert nuclei and physical conditions of the electron scattering regions. ",
        "introduction": "It is generally considered that active galactic nuclei (AGNs) are powered by single, accreting supermassive black holes (e.g., Rees 1984; Blandford 1990). According to this scenario, the accretion rate onto a black hole is an important parameter to explain the huge luminosity of AGNs.  However, since the luminosity released from this central engine is proportional to the mass of the black hole [ie., the Eddington luminosity, $L_{\\rm Edd} \\sim 10^{46} (M_{\\rm BH}/10^8 M_\\odot)$ erg s$^{-1}$ where $M_{\\rm BH}$ is the black hole mass], the mass itself is considered as another important  parameter (Blandford 1990). Relationships between mass and luminosity of AGNs provide important information about the nature of central engines (e.g., Wandel \\& Yahil 1985; Padovani \\& Rafanelli 1988; Padovani 1989; Koratkar \\& Gaskell 1991b). Further, the mass function of nuclei may place  constraints on the formation and evolution of supermassive black holes in the universe (Padovani, Burg, \\& Edelson 1990; Haehnelt \\& Rees 1993). Therefore, the mass of AGNs is of fundamental importance in understanding the AGN phenomena.  In order to estimate the nuclear mass, the so-called dynamical method has been often used (Dibai 1981, 1984; Wandel $\\&$ Yahil 1985; Wandel \\& Mushotzky 1986; Joly et al. 1985; Reshetnikov 1987; Padovani \\& Rafanelli 1988; Padovani, Burg, \\& Edelson 1990;  Koratkar \\& Gaskell 1991b). If the gas motion in a broad emission-line region (BLR) is dominated by the gravitational force exerted by the central massive object, the line width can be used to estimate the mass of the central object given the radial distance of the BLR (Woltier 1959; Setti \\& Woltier 1966). Since the recent elaborate monitoring observations of AGNs have shown that the gas motion in the BLRs is almost dominated by the gravitation (Gaskell 1988; Koratkar \\& Gaskell 1991a, 1991c; Clavel et al. 1991; Peterson 1993; Robinson 1994; Korista et al. 1995; Wanders et al. 1995; Wanders \\& Peterson 1996), the basic assumption in the dynamical method is considered to be robust. In section 2, we discuss the method in detail.   All the previous estimates of nuclear mass have been made for type 1 Seyfert nuclei (hereafter S1s) and quasars because the dynamical method needs both the flux and the velocity width of broad line emission. This raises the question ^^ ^^ How massive are type 2 Seyfert nuclei (hereafter S2s)  and are they similar to those of S1s ?'' Since the discovery of hidden BLR in the archetypical S2 nucleus of NGC 1068 by Antonucci \\& Miller (1985), it has been considered that S2s are S1s in which the BLR as well as the central engine are hidden from direct view (see, for a review Antonucci 1993). Taking this unified scheme into account, we may expect that there is no systematic difference in the nuclear mass between S1s and S2s. Miller \\& Goodrich (1990; hereafter MG90) made a systematic study of hidden BLRs of high-polarization S2s and found that the properties of the hidden BLRs studied by polarized broad H$\\alpha$ and  H$\\beta$ emission are nearly  the same as those of S1s in the following respects; equivalent widths, line widths, reddening, and luminosities (see also Tran 1995a). The intrinsic H$\\beta$ luminosities\\footnote{We adopt a Hubble constant $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and a deceleration parameter $q_0 = 0$ throughout this paper.} of the S2s amount to $\\sim 10^{43}$ erg s$^{-1}$. MG90 adopted a typical H$\\beta$ luminosity of $10^{43}$ erg s$^{-1}$ for S1s and thus reached the conclusion that the intrinsic H$\\beta$ luminosities are nearly the same between S1s and S2s. However, the typical H$\\beta$ luminosities of S1s are $\\sim 10^{41}$ - $10^{42}$ erg s$^{-1}$  (Yee 1980; Blumenthal, Keel, \\& Miller 1982; Dahari \\& De Robertis 1988). Therefore, the intrinsic H$\\beta$ luminosities of S2s may be significantly more luminous than those of S1s. However, although this suggests that there is a certain systematic difference between S1s and S2s, it is noted that the intrinsic H$\\beta$ luminosities of S2s depend on the estimates of the optical depth for electron scattering, the degree  of {\\it true} polarization, and the covering factor of the scatterers (MG90). Thus the comparison of broad H$\\beta$ luminosities between S1s and S2s must be made carefully. After MG90, several new spectropolarimetric observations of S2s have been published (e.g., Tran, Miller, \\&  Kay 1992; Antonucci, Hurt, \\& Miller 1994; Tran 1995a, 1995b). New interpretations on the observed low polarizations have  been also presented (Cid Fernandes \\& Terlevich 1995; Heckman et al. 1995; Tran 1995c; Kishimoto 1996, 1997). Therefore, it is interesting to revisit the comparison between the hidden BLRs in S2s and the ordinary  BLRs in S1s and then to estimate the dynamical masses  of S2 nuclei. ",
        "conclusions": "We have derived the dynamical masses of S2 nuclei. Comparing them with those of S1s, we have obtained that the nuclear masses of S2s are similar to those of S1s provided that $\\tau_{\\rm es} \\sim 1$ while more massive if $\\tau_{\\rm es} \\ll 1$. For example, if $\\tau_{\\rm es} \\sim 0.1$, the nuclear masses of S2s would be systematically larger by an order of magnitude than those of S1s. We now consider physical conditions in the electron scattering regions, and we discuss some implications for the unified model of Seyfert nuclei.  Low optical depths have been adopted in previous works based on spectropolarimetry of S2s; e.g., $\\tau_{\\rm es} \\simeq$ 0.05 - 0.1 (Antonucci \\& Miller 1985; MG90; Miller et al. 1991). As discussed by MG90, if $\\tau_{\\rm es} \\ll  1$, we would observe many AGN with polarizations higher than 50\\% provided that  the half opening angle of the ionization cone $\\theta_{\\rm c} = \\theta_{\\rm open}/2 \\sim 30^\\circ$ as is observed for many S2s (Pogge 1989; Wilson \\& Tsvetanov 1994; Schmitt \\& Kinney 1996). However, the highest polarization observed so far is 16\\% (NGC 1068: MG90; Antonucci et al. 1994) and the typical polarization is only several percent for the other S2s (MG90; Tran 1995a). The observed polarization is lower than expected, even after  the effect of dilution due to the unpolarized continuum radiation is taken into account (Cid Fernandes \\& Terlevich 1995; Heckman et al 1995; Tran 1995c; Tran, Cohen \\& Goodrich 1995; see also  Kishimoto 1996). Therefore, the optical thick condition cannot be ruled out entirely at present.  We now discuss the physical characteristics of the electron scattering region. The optical depth for electron scattering is estimated by  \\begin{equation} \\tau_{\\rm es} \\sim \\sigma_{\\rm T} {\\overline{n}_{\\rm e}} l_{\\rm eff} \\sim \\sigma_{\\rm T} N_{\\rm e} \\end{equation} where $\\sigma_{\\rm T}$ is the Thomson cross section, 0.66$\\times 10^{-24}$ cm$^2$, ${\\overline{n}_{\\rm e}}$ is the average electron density in the scattering region, $l_{\\rm eff}$ is the effective path length, and $N_{\\rm e}$ is the electron column density. We obtain $N_{\\rm e} \\sim 10^{24}$ cm$^{-2}$ for $\\tau_{\\rm es} \\sim 1$ while $N_{\\rm e} \\sim 10^{23}$ cm$^{-2}$ for $\\tau_{\\rm es} \\sim 0.1$. If the kinetic temperature of free electrons is as high as $\\sim 10^6$ K, significant line broadening would  occur due to the scattering (Antonucci \\& Miller 1985; MG90). Since $v \\sim 2000 (T_{\\rm e}/10^5 {\\rm K})$ km s$^{-1}$, it seems reasonable to assume  $T_{\\rm e} \\sim 10^5$ K at most (see also Miller et al. 1991).       If the gas in the electron scattering region is in pressure equilibrium with the BLR gas ($n_{\\rm e} \\sim 10^9$ cm$^{-3}$ and $T_{\\rm e} \\sim 10^4$ K; Osterbrock 1989), then the average electron density in the scattering region is ${\\overline{n}_{\\rm e}} \\sim 10^8$ cm$^{-3}$,  and the effective path lengths are $\\sim 10^{16}$ cm and $\\sim 10^{15}$ cm for $\\tau_{\\rm es} \\sim 1$  and $\\sim 0.1$, respectively. Since we observe the BLR through the scatterers, the scattering regions are located outside the BLRs and thus the radial distance of scatterers is $r_{\\rm e} > 10^{16}$ cm. In the case of NGC 1068, it is observed that the scattering regions are extended to $\\sim$ 100 pc from the nucleus (Capetti et al. 1995a, 1995b). This large value was indeed suspected from the photoionization consideration by Miller et al. (1991). It is therefore considered that the electron scattering regions are located at $r_{\\rm e} \\sim 10^{16}$ - $10^{20}$ cm. The effective radius of scatterers may be different from object to object because it depends also on how we observe the dusty tori in AGNs (i.e., viewing angle dependent; cf. MG90, Miller et al. 1991; Kishimoto 1996, 1997).  It is worth noting that the above physical conditions are quite similar to those of warm absorbers probed by X-ray spectroscopy of type 1 AGNs (Halpern 1984; Netzer 1993; Nandra \\& Pounds 1992, 1994; Ptak et al. 1994; Reynolds \\& Fabian 1995; Reynolds et al. 1995; Otani et al. 1996; Reynolds 1997). However, the electron column density inferred in this study, $N_{\\rm e} \\sim 10^{23}$ - $10^{24}$ cm$^{-2}$, is higher by one to two orders of magnitude than those of warm absorbers, $\\sim 10^{22}$ cm$^{-2}$. This seems to be  inconsistent with the strict unified model of Seyfert nuclei (Antonucci \\& Miller 1985). We may, however, consider the higher column densities in the S2s as  due to the  effect of multiple scattering if $\\tau_{\\rm es} \\sim 1$. Another interpretation may be that S2s are gas-richer systematically  than S1s (cf. Heckman et al. 1989; Taniguchi 1997).   \\vspace{0.5cm}  We  would like to thank Makoto Kishimoto, Youichi Ohyama, and Takashi Murayama for useful discussion. We also thank the anonymous referee for his/her many useful comments and suggestions which improved this paper significantly. This work was financially supported in part by Grant-in-Aids for the Scientific Research (No. 07044054) of the Japanese Ministry of Education, Culture, Sports, and Science.  \\newpage   \\begin{table} \\caption{Comparison of the FWHM(H$\\beta$)\\tablenotemark{a} ~ between MG90 and Tran(1995a)} \\vspace{5mm} \\begin{tabular}{ccc} \\tableline \\tableline Galaxy  & MG90 & Tran (1995a)  \\\\ \\tableline Mrk 3 & 5400 & 6000$\\pm$500 \\\\ Mrk 463E & 3000 & 2770$\\pm$180 \\\\ NGC 7674 & 1500 & 2830$\\pm$150 \\\\ \\tableline \\tablenotetext{a}{In units of km s$^{-1}$.} \\end{tabular} \\end{table}    \\begin{table} \\caption{The dynamical masses of Seyfert 2 nuclei} \\vspace{5mm} \\begin{tabular}{ccccccc} \\tableline \\tableline Object & FWHM(H$\\beta_{\\rm b}$) & $L({\\rm H}\\beta_{\\rm b})_{\\rm p}$\\tablenotemark{a} & $P$\\tablenotemark{b} & $M_{\\rm BH} (\\tau_{\\rm es}=1)$ & $M_{\\rm BH} (\\tau_{\\rm es}=0.1)$ & Ref.\\tablenotemark{c} \\\\ & (km s$^{-1}$) & (erg s$^{-1}$) & (\\%) & ($M_{\\odot}$) & ($M_{\\odot}$) &  \\\\ \\tableline NGC 1068 & 3030 & $1.11\\times10^{39}$ & 16 & $1.73\\times10^7$ & $5.16\\times10^7$ & 1, 2 \\\\ NGC 7212 & 5420 & $3.40\\times10^{39}$ & 22 & $8.09\\times10^7$ & $2.42\\times10^8$ & 2  \\\\ NGC 7674 & 2830 & $5.36\\times10^{39}$ & 8 &  $4.43\\times10^7$ & $1.32\\times10^8$ & 2  \\\\ Mrk 3    & 6000 & $1.00\\times10^{40}$ & 20 & $1.73\\times10^8$ & $5.17\\times10^8$ & 2  \\\\ Mrk 348  & 9350 & $2.95\\times10^{39}$ & 35\\tablenotemark{d} & $1.81\\times10^8$ & $5.39\\times10^8$ & 2 \\\\ Mrk 463E & 2770 & $3.99\\times10^{40}$ & 10 & $9.89\\times10^7$ & $2.95\\times10^8$ & 2 \\\\ Mrk 477  & 4130 & $7.85\\times10^{40}$ & 2 &  $6.66\\times10^8$ & $1.96\\times10^9$ & 2 \\\\ Mrk 1210 & 3080 & $2.22\\times10^{39}$ & 15 & $2.56\\times10^7$ & $7.63\\times10^8$ & 2 \\\\ Was 49   & 5860 & $3.30\\times10^{40}$ & 20 & $2.91\\times10^8$ & $8.69\\times10^8$ & 2 \\\\ \\tableline \\tablenotetext{a}{Luminosity of polarized, broad H$\\beta$ emission.} \\tablenotetext{b}{Intrinsic polarization corrected for the unpolarized continuum emission as well as interstellar polarization taken from Tran (1995c).} \\tablenotetext{c}{1. Miller \\& Goodrich 1990; 2. Tran 1995a, 1995c.} \\tablenotetext{d}{The intrinsic polarization for H$\\alpha$ emission.} \\end{tabular} \\end{table}   \\begin{table} \\caption{Comparison of the average dynamical masses between S2s and S1s} \\vspace{5mm} \\begin{tabular}{ccc} \\tableline \\tableline Sample & Number & $M_{\\rm BH}$ \\\\ &  & ($M_{\\odot}$) \\\\ \\tableline S2 ($\\tau_{\\rm es}=1$) & 9 & 8.00$\\pm$0.51 \\\\ S2 ($\\tau_{\\rm es}=0.1$) & 9 & 8.47$\\pm$0.51 \\\\ S1 (PR88) & 30 & 7.89$\\pm$0.57 \\\\ S1 (PBE90) & 25 & 7.48$\\pm$0.63 \\\\ \\tableline \\end{tabular} \\end{table}   \\newpage"
    },
    "9803/astro-ph9803272_arXiv.txt": {
        "abstract": "To extract reliable cosmic parameters from cosmic microwave background datasets, it is essential to show that the data are not contaminated by residual non-cosmological signals.  We describe general statistical approaches to this problem, with an emphasis on the case in which there are two datasets that can be checked for consistency.  A first visual step is the Wiener filter mapping from one set of data onto the pixel basis of another.  For more quantitative analyses we develop and apply both Bayesian and frequentist techniques.  We define the ``contamination parameter'' and advocate the calculation of its probability distribution as a means of examining the consistency of two datasets.  The closely related ``probability enhancement factor'' is shown to be a useful statistic for comparison; it is significantly better than a number of $\\chi^2$ quantities we consider.  Our methods can be used: internally (between different subsets of a dataset) or externally (between different experiments); for observing regions that completely overlap, partially overlap or overlap not at all; and for observing strategies that differ greatly.  We apply the methods to check the consistency (internal and external) of the MSAM92, MSAM94 and Saskatoon Ring datasets.  From comparing the two MSAM datasets, we find that the most probable level of contamination is 12\\%, with no contamination only 1.05 times less probable, 50\\% contamination about 8 times less probable and 100\\% contamination strongly ruled out at over $2\\times 10^5$ times less probable.  From comparing the 1992 MSAM flight with the Saskatoon data we find the most probable level of contamination to be 50\\%, with no contamination only 1.6 times less probable and 100\\% contamination 13 times less probable.  Our methods can also be used to calibrate one experiment off of another.  To achieve the best agreement between the Saskatoon and MSAM data we find that the MSAM data should be multiplied by (or Saskatoon data divided by): $1.06^{+0.22}_{-0.26}$. ",
        "introduction": "The cosmic microwave background (CMB) is black body radiation with a mean temperature of $2.728 \\pm 0.002$ K \\cite{firas}.  This mean is modulated by a dipole due to our peculiar motion with respect to the radiation field.  If one removes the dipole, the temperature is uniform in every direction to $\\pm 100 \\muK$.  Precision measurement of these tiny deviations from isotropy can tell us much about the Universe \\cite{forecast}.  Unfortunately, precision measurement of $~100\\muK$ fluctuations is not an easy task.  Even given sufficient detector sensitivity and observing time, one still has to contend with many possible contaminants such as side lobe pickup of the $300^\\circ$ Kelvin Earth and atmospheric noise (even from high-altitude balloons). In addition there can be contamination of CMB observations by astrophysical foregrounds.  Despite these difficulties there is good reason to believe that, at least for some experiments, the signals observed from sub-orbital platforms are not dominated by contaminants.  One of the best reasons for believing this comes from the comparisons that have been done---between FIRS and DMR \\cite{Ganga}, Tenerife and DMR \\cite{Lineweaver}, MSAM and Saskatoon \\cite{nett95}, two years of Python data \\cite{Ruhl}, and two flights of MSAM \\cite{Inman}.  Especially for the case when data being compared are from two different instruments, almost the only thing their acquisitions have in common is that they were observing the same piece of sky--each dataset has entirely different sources of systematic error.  In addition to confirming the astrophysical origin of the estimated signal, comparison can greatly improve the ability to detect foreground contamination.  Perhaps the best evidence for the thermal nature of anisotropy comes from the comparison between the MSAM92 and Saskatoon datasets.  Together, these observations span a frequency range from 36 GHz to greater than 170 GHz.  In \\cite{nett95} it was found that the spectral index $\\beta$ ($\\delta T \\propto (\\nu/\\nu_0)^\\beta$) is constrained to be $\\beta = -0.1 \\pm 0.2$.  For CMB, free-free and dust over this frequency range we expect $\\beta = 0$, $-1.45$ and $2.25$, respectively. The authors conclude that the signals (in the region of overlap) are not dominated by contamination from known astrophysical foregrounds, but are, rather, primarily CMB.  We should not let this apparent success fool us into thinking that going to the next level of precision will be easy.  There is a big difference in the level of toleration of contaminants when the goal switches from detection to precision measurement.  It is likely that there will be significant levels of contamination (from the atmosphere, side lobes, and foregrounds) in future sub-orbital missions.  It may be difficult to convincingly demonstrate that contamination is low without comparison.  Given the importance of comparison, we feel it is worth improving upon the methods used previously.  Past treatments have had to ignore much relevant data, and make uncontrolled approximations.  This is due to the fact that generally the two datasets being compared were obtained from instruments observing the sky in different ways.  The beam patterns and differencing schemes may differ as in the case of the MSAM/Saskatoon comparison.  In \\cite{nett95} one of the MSAM differencing schemes was approximately recreated in software in order to do the comparison.  However, no use of software could change the fact that the MSAM and Saskatoon beam patterns, although they have fairly similar full-widths at half-maximum, differ significantly in shape.  Even when the differencing schemes and beam patterns are the same, there can still be barriers to a direct comparison.  The two MSAM flights took data with essentially the same beam pattern and applied the same differencing, but in this case the direct comparison is frustrated by the fact that the pixels do not all line up exactly.  Therefore in \\cite{Inman}, pixels within half a beam width of each other were approximated as being at the same point, and those pixels with no partner from the other dataset within this distance were ignored. Half of the data were lost this way.  Here we develop methods of comparing datasets that do not require any information to be thrown away.  Differences in demodulation schemes, and effects due to non-overlapping pixels are automatically taken into account.  The inevitable price we pay for this is model-dependence.  However, we generally expect the model-dependence to be small and indeed find it to be so in the case studies shown here.  An extremely useful tool for visual comparison is the Wiener filter. Roughly speaking, it allows us to interpolate the results from one experiment onto the expected results for another experiment that has observed the sky differently.  After some notational preliminaries in section II, in section III we introduce the Wiener filter in the context of the probability distribution of the signal, given the data.  Also in this section we describe the datasets and apply the Wiener filter to them.  When comparing datasets we are testing the consistency of our model of the datasets.  We emphasize that meaningful model consistency testing demands the existence of other models with which to compare. Therefore we extend our model of the data to include a possible contaminant and calculate the probability distribution of its amplitude, given the data.  We find a more limited extension of the model space to also be useful, in which we only consider one alternative to no contamination: complete contamination.  We define the ``probability enhancement factor'' as the logarithm of the ratio of the probability of no contamination to the probability of complete contamination.  This Bayesian approach to comparison is described and applied in section IV.  In section V we discuss and apply frequentist techniques such as $\\chi^2$ tests.  The probability enhancement factor can also be used as the basis for a frequentist test---and it is in fact the well-known likelihood ratio test.  We demonstrate that the probability enhancement factor has more discriminatory power than any of the other tests considered.  After a further look at the data with the probability enhancement factor in section VI, we discuss the fixing of relative calibration in section VII and possible contamination due to dust in section VII. Finally we summarize our results in section IX. ",
        "conclusions": "We have demonstrated the usefulness of the Wiener filter for making visual comparisons of datasets.  We have emphasized that meaningful consistency testing requires alternative models with which to compare. Thus we have explicitly extended our model of the data to include a possible contaminant and calculated the probability distribution of the amplitude of this contaminant.  For purposes of extracting just one number from the comparison we advocate calculating the ratio of the probability of no contamination to the probability of infinite contamination.  Viewed as a statistic, we have shown this ``probability enhancement factor'' to be better than various $\\chi^2$ statistics at discriminating between competing hypotheses.  The utility of our comparison statistics was shown by exercising them on the MSAM92, MSAM94 and SK95 data. We have found from comparing MSAM92 and MSAM94 that the most probable level of contamination is 12\\%, with zero contamination only 1.05 times less probable, and total contamination over $2\\times 10^5$ times less probable.  From comparing MSAM92 and SK95 we have found that the most probable level of contamination is 50\\%, with zero contamination only 1.6 times less probable, and total contamination 13 times less probable.  Looking at subsets of the data we find a region at large RA where the SK and MSAM measurements disagree.  From IRAS and from the MSAM dust measurements we know that this region is also the dustiest region of the overlap between SK and MSAM.  The origin of the discrepancy is unclear and may be due to instrumental artifacts in SK, or foreground contamination of either the SK or MSAM measurements.  A revolution is underway in the quality and quantity of CMB data---a revolution generated by the satellites MAP and Planck \\cite{satellites} as well as by a number of balloon and ground-based programs.  The amount of data may soon be too large for the type of complete statistical analysis described here.  However, any approximate methods developed for extracting the power spectrum or parameters will also be applicable to the statistical procedures introduced here."
    },
    "9803/astro-ph9803208_arXiv.txt": {
        "abstract": "We present analysis of the shape and radial mass distribution of the E4 galaxy NGC 3923 using archival X-ray data from the {\\sl ROSAT} PSPC and HRI. The X-ray isophotes are significantly elongated with ellipticity $\\epsilon_x=0.15 (0.09-0.21)$ (90\\% confidence) for semi-major axis $a\\sim 10h^{-1}_{70}$ kpc and have position angles aligned with the optical isophotes within the estimated uncertainties. Applying the Geometric Test for dark matter, which is independent of the gas temperature profile, we find that the ellipticities of the PSPC isophotes exceed those predicted if $M\\propto L$ at a marginal significance level of $85\\% (80\\%)$ for oblate (prolate) symmetry. Detailed hydrostatic models of an isothermal gas yield ellipticities for the gravitating matter, $\\epsilon_{mass}=0.35-0.66$ (90\\% confidence), which exceed the intensity weighted ellipticity of the $R$-band optical light, $\\langle \\epsilon_R\\rangle = 0.30$ ($\\epsilon_R^{max}=0.39$).  We conclude that mass density profiles with $\\rho\\sim r^{-2}$ are favored over steeper profiles if the gas is essentially isothermal (which is suggested by the PSPC spectrum) and the surface brightness in the central regions $(r\\la 15\\arcsec)$ is not modified substantially by a multi-phase cooling flow, magnetic fields, or discrete sources.  We argue that these effects are unlikely to be important for NGC 3923.  (The derived $\\epsilon_{mass}$ range is very insensitive to these issues.)  Our spatial analysis also indicates that the allowed contribution to the {\\it ROSAT} emission from a population of discrete sources with $\\Sigma_x\\propto\\Sigma_R$ is significantly less than that indicated by the hard spectral component measured by {\\sl ASCA}. ",
        "introduction": "\\label{intro}  The structure of the dark matter halos of galaxies provides important clues to their formation and dynamical evolution (e.g.  Sackett 1996; de Zeeuw 1996, 1997). For example, in the Cold Dark Matter (CDM) scenario (e.g.  Ostriker 1993) there is evidence that the density profiles of halos have a universal form essentially independent of the halo mass or $\\Omega_0$ (Navarro, Frenk, \\& White 1997; though see Moore et al. 1997). The intrinsic shapes of CDM halos are oblate-triaxial with ellipticities similar to the optical isophotes of elliptical galaxies (e.g.  Dubinski 1994). The global shape of a halo also has implications for the mass of a central black hole (e.g. Merritt \\& Quinlan 1997).  At present accurate constraints on the intrinsic shapes and density profiles of early-type galaxies are not widely available (e.g. Sackett 1996; Olling \\& Merrifield 1997)\\footnote{The distribution of dark matter in spiral galaxies is also far from being a solved problem -- see, e.g.  Broeils \\shortcite{broeils}.}. Stellar dynamical analyses that have incorporated the information contained in high order moments of stellar velocity profiles have made important progress in limiting the uncertainty in the radial distribution of gravitating mass arising from velocity dispersion anisotropy (Rix et al. 1997; Gerhard et al. 1997). However, as indicated by the paucity of such stellar dynamical measurements, the required observations to obtain precise constraints at radii larger than $\\sim R_e$ are extensive, and the modeling techniques to recover the phase-space distribution function are complex. It is also unclear whether this method can provide interesting constraints on the intrinsic shapes since only weak limits on the range of possible shapes have been obtained from analysis of velocity profiles out to $\\sim 2$ $R_e$ (e.g.  Statler 1994).  Interesting measurements of the ellipticity of the gravitating mass have been obtained for two Polar Ring galaxies (Sackett et al. 1994; Sackett \\& Pogge 1995) and from statistical averaging of known gravitational lenses (e.g.  Keeton, Kochanek, \\& Falco 1997), but owing to the rarity of these objects it is possible that the structures of their halos are not representative of most early-type galaxies. Moreover, gravitational lenses, which are biased towards the most massive galaxies, only give relatively crude constraints on the ellipticity and radial mass distribution for any individual system and only on scales similar to the Einstein radius (e.g.  Kochanek 1991).  The X-ray emission from hot gas in isolated early-type galaxies (Forman, Jones, \\& Tucker 1985; Trinchieri, Fabbiano, \\& Canizares 1986; for a review see Sarazin 1997) probably affords the best means for measuring the shapes and radial mass distributions in these systems (for a review see Buote \\& Canizares 1997b; also see Schechter 1987 and the original application to the analogous problem of the shapes of galaxy clusters by Binney \\& Strimple 1978). The isotropic pressure tensor of the hot gas in early-type galaxies greatly simplifies measurement of the mass distribution over stellar dynamical methods. Moreover, since the shape of the volume X-ray emission traces the shape of the gravitational potential independent of the (typically uncertain) gas temperature profile (Buote \\& Canizares 1994, 1996a), the shape of the mass distribution can be accurately measured in a way that is quite robust to the possible complicating effects of multi-phase cooling flows and magnetic fields (see Buote \\& Canizares 1997b).  Presently, X-ray measurements of the mass distributions in early-type galaxies are inhibited by limitations in the available data. The {\\sl ROSAT} \\cite{trump} Position Sensitive Proportional Counter (PSPC) \\cite{pf} has inadequate spatial resolution (PSF $\\sim 30\\arcsec$ FWHM) to map the detailed mass distributions for all but the largest nearby galaxies, and the limited spectral resolution and band width complicates interpretation of the measured temperature profiles (Buote \\& Canizares 1994; Trinchieri et al. 1994; Buote \\& Fabian 1997). Although equipped with superior spatial resolution (PSF $\\sim 4\\arcsec$ FWHM), the {\\sl ROSAT} High Resolution Imager (HRI) \\cite{david} has too small an effective area and too large an internal background to provide images of sufficient quality for many galaxies for radii $r\\ga R_e$. Among the few galaxies with detailed measurements of their radial mass profiles are NGC 507 (Kim \\& Fabbiano 1995), NGC 1399 (Rangarajan et al. 1995; Jones et al. 1997), NGC 4472 (Irwin \\& Sarazin 1996), NGC 4636 (Trinchieri et al. 1994), NGC 4649 \\cite{bm}, and NGC 5044 (David et al. 1994).  The shape of the gravitating mass has been measured via X-ray analysis for the E4 galaxy NGC 720 and the E7/S0 galaxy NGC 1332 and found to be at least as elongated as the optical isophotes (Buote \\& Canizares 1994, 1996a, 1997a).  For NGC 720, which has more precise constraints, the ellipticity of the gravitating matter is $\\epsilon_{mass}=0.44-0.68$ (90\\% confidence) compared to the intensity weighted ellipticity of the optical light, $\\langle\\epsilon\\rangle=0.31$ (Buote \\& Canizares 1997a).  In addition, the X-ray isophotes of NGC 720 twist from being aligned with the optical isophotes within $R_e$ to a position $\\sim 30\\degr$ offset at larger radii. This twist, when combined with the ellipticities of the X-ray isophotes, cannot be explained by the projection of a reasonable triaxial matter distribution and thus may implicate a dark matter halo misaligned from the stars (Buote \\& Canizares 1996b; Romanowsky \\& Kochanek 1997).  NGC 720 and NGC 1332 were selected for analysis since they are isolated, significantly elongated in the optical, sufficiently bright, and sufficiently dominated by emission from hot gas in the {\\sl ROSAT} band.  In this paper we present X-ray analysis of the classic ``shell'' galaxy, NGC 3923, which is the last galaxy of which we are aware that satisfies these selection criteria and has deep {\\sl ROSAT} observations.  This isolated E4 galaxy has both archival {\\sl ROSAT} PSPC and HRI data and its {\\sl ASCA} spectrum has been analyzed previously \\cite{bf}. This will serve as our final case study until the impending launch of {\\sl AXAF} revolutionizes this field.  The organization of this paper is as follows. In \\S \\ref{obs} we describe the {\\sl ROSAT} observations and the data reduction. We discuss removal of point sources in \\S \\ref{pt}. Measurements of the ellipticities of the X-ray isophotes and the radial profiles are described in \\S \\ref{e0} and \\S \\ref{radpro} respectively. Analysis of the PSPC spectrum is presented in \\S \\ref{spectra}. We give results for the Geometric Test for dark matter in \\S \\ref{gt} and constraints on the shape and radial mass distribution from detailed hydrostatic models in \\S \\ref{models}. Finally, in \\S \\ref{conc} we give our conclusions. ",
        "conclusions": "\\label{conc}  We have analyzed the gravitating matter distribution of the E4 galaxy NGC 3923 using archival X-ray data from the {\\sl ROSAT} PSPC and HRI. Analysis of the PSPC data, which allows more precise constraints than the HRI data, demonstrates that the X-ray isophotes are significantly elongated with ellipticity $\\epsilon_x=0.15 (0.09-0.21)$ (90\\% confidence) for semi-major axis $a\\sim 10h^{-1}_{70}$ kpc and have position angles aligned with the optical isophotes within the estimated uncertainties. A bright point source located $\\sim 100\\arcsec$ along the major axis inhibits reliable ellipticity constraints for larger radii.  By applying a ``Geometric Test'' for dark matter, which essentially compares the shapes of the observed X-ray isophotes to those predicted if mass traces the optical light $L$ (independent of the poorly constrained temperature profile of the gas), we found that the ellipticity of the PSPC X-ray surface brightness exceeds that predicted by the constant $M/L$ hypothesis at the 80\\%-85\\% confidence level. The ``Geometric Test'' result is conservative since it only considers signatures of dark matter that are distributed differently from the optical light.  Although the evidence for dark matter from the Geometric Test is marginal, the results from models which employ an explicit solution of the hydrostatic equation assuming an isothermal gas (which is supported by the PSPC spectrum -- \\S \\ref{spectra}) indicate that $M\\propto L$ is highly inconsistent with the radial profiles of the PSPC and HRI data ($\\chi^2_{\\rm red}=3.5$ for 16 dof). This particular discrepancy arises because $L$ is too centrally concentrated: the derived scale length of the gravitating matter is approximately 1.5-2 times that of $L$. The ellipticities predicted by this $M\\propto L$ model fall below the PSPC data at a significance slightly greater than the 90\\% level. Modeling the gravitating mass with a density run $\\rho\\sim r^{-2}$ or with a Hernquist profile we find that the ellipticity of the gravitating matter is, $\\epsilon_{mass}\\cong 0.35 - 0.65$ (90\\% confidence), which is larger than the intensity weighted optical ellipticity $\\langle\\epsilon\\rangle = 0.30$.  This evidence for dark matter which is more flattened and more extended than $L$ is similar to our conclusions from previous X-ray studies of two other ellipticals, NGC 720 and NGC 1332, but at somewhat smaller significance level than for NGC 720 (e.g.  Buote \\& Canizares 1997b). These results are consistent with analyses of known gravitational lenses (e.g.  Keeton, Kochanek, \\& Falco 1997), two polar ring galaxies (Sackett et al. 1994; Sackett \\& Pogge 1995), and flaring disks in spiral galaxies (e.g, Olling 1996). The ellipticities of the gravitating matter derived from our X-ray analyses and these other methods are consistent with those of halos produced by CDM simulations (e.g.  Dubinski 1994).  If an isothermal gas is assumed then models with matter density $\\rho\\sim r^{-2}$ are favored over Hernquist models (and similar models like the universal CDM profile of Navarro et al. 1997).  For $r\\sim 100\\arcsec-300\\arcsec$ the $\\rho\\sim r^{-2}$ model marginally fits the data better than the Hernquist model.  However, most of the difference in these models occurs in the central radial bins where the effects of multi-phase cooling flows, magnetic fields, and discrete sources could affect the surface brightness profiles, though we have argued the effects are unlikely to be important (see \\S \\ref{models}). (The derived shape of the gravitating mass is mostly robust to these issues -- Buote \\& Canizares 1997b.) This support for nearly $r^{-2}$ profiles agrees with previous studies of gravitational lenses (e.g. Maoz \\& Rix 1993; Kochanek 1995), although a recent paper finds that density profiles with changing slopes (e.g.  Hernquist and NFW) are preferred \\cite{lilya}.  An emission component that is proportional to $L$ cannot contribute significantly to the {\\sl ROSAT} X-ray emission of NGC 3923, and thus discrete sources should not affect our constraints on the gravitating matter (Buote \\& Canizares 1997a). However, the {\\sl ASCA} spectral data when fitted with two thermal components yield a cold component, $T_C=0.55$ keV, and a hot component, $T_H\\sim 4$ keV, where the relative flux of cold-to-hot is $\\sim 1.9$ in the {\\sl ROSAT} band \\cite{bf}. The conventional interpretation of the hot component (e.g. Matsumoto et al. 1997; Loewenstein \\& Mushotzky 1997) is that it arises from discrete sources. But our analysis (\\S \\ref{radpro}) shows that $\\sim 35\\%$ of the 0.5-2 keV emission cannot be distributed like the optical light which would be expected of discrete sources.  Hence, either the emission from discrete sources is not distributed like $L$, or the hot component obtained from the spectral fits cannot be entirely due to discrete sources as suggested by Buote \\& Fabian \\shortcite{bf}.  The constraints we have obtained for NGC 720, NGC 1332, and now NGC 3923 from analyses of their X-ray isophote shapes and radial surface brightness profiles provide an initial demonstration of the power of X-ray analysis for probing the shape and radial distribution of gravitating matter in early-type galaxies.  The next generation of X-ray satellites, particularly {\\sl AXAF} and {\\sl XMM}, have the capability to accurately map X-ray isophote shapes and orientations from the cores ($r\\sim 1\\arcsec$) out to 10s of kpc for many galaxies\\footnote{The vastly improved spatial resolution of {\\sl AXAF} over the {\\sl ROSAT} PSPC will allow easy exclusion of the bright point source (1) (see Table \\ref{tab.src}) which hindered the present analysis of NGC 3923.}. The spatially resolved spectra provided by these future missions will allow more precise constraints on temperature gradients and the contribution from discrete sources. Thus, unlike most other methods, obtaining interesting X-ray constraints on the shape and radial density profile of the gravitating matter will be possible for a large sample of early-type galaxies since the X-ray analysis is applicable to any isolated early-type galaxy whose soft X-ray emission ($\\sim 0.5-2$ keV) is dominated by hot gas."
    },
    "9803/astro-ph9803178_arXiv.txt": {
        "abstract": "A deep, fuly sampled diffraction limited (FWHM $\\sim$ 70 mas) narrow-band image of the central region in M87 was obtained with the Wide Filed and Planetary Camera 2 of the {\\it Hubble Space Telescope} using the dithering technique.  The \\HaNii\\ continuum subtracted image reveals a wealth of details in the gaseous disk structure described earlier by Ford et al.\\ (1994). The disk morphology is dominated by a well defined three-arm spiral pattern. In addition, the major spiral arms contain a large number of small ``arclets'' covering a range of sizes (0\\as1--0\\as3 = 10--30 pc). The overall surface brightness profile inside a radius $\\sim$ 1\\farcs5 (100 pc) is well represented by a power-law $I(\\mu) \\sim \\mu^{-1.75}$, but when the central $\\sim$ 40 pc are excluded it can be equally well fit by an exponential disk. The major axis position angle remains constant at about PA$_{\\rm disk} \\sim 6^{\\circ}$ for the innermost $\\sim 1''$, implying the disk is oriented nearly perpendicular to the synchrotron jet (PA$_{\\rm jet} \\sim 291^{\\circ}$). At larger radial distances the isophotes twist, reflecting the gas distribution in the filaments connecting to the disk outskirts. The ellipticity within the same radial range is $e = 0.2-0.4$, which implies an inclination angle of $i \\sim 35^{\\circ}$. The sense of rotation combined with the dust obscuration pattern indicate that the spiral arms are trailing. ",
        "introduction": "The disk of ionized gas in the nucleus of M87 is currently the best example of a family of similar small ($r \\sim 100$ pc) gaseous disks found to be common in the centers of elliptical galaxies with active nuclei (for a review see Ford et al.\\ 1998). Several \\HST\\ kinematical studies have shown that in M87 the gas is in Keplerian rotation, orbiting a massive black hole with a mass $M_{\\rm BH} \\sim 2 - 3 \\times 10^9 M_{\\odot}$ (Harms et al.\\ 1994; hereafter H94, Ford et al.\\ 1996a,b; and Macchetto et al.\\ 1997, hereafter M97). The few other galaxies studied kinematically so far (NGC 4261 -- Ferrarese et al.\\ 1996, NGC 6521 -- Ferrrarese et al.\\ 1998, NGC 4374 -- Bower et al.\\ 1998) have further shown that nuclear gaseous disks offer an excellent tool for measuring the central black hole mass.  Recent studies have revealed other important characteristics of the nuclear disk in M87. Its aparent minor axis (F96, M97) is closely aligned with the synchrotron jet ($\\Delta\\theta \\sim 10^{\\circ} - 15^{\\circ}$) suggesting a causal relationship between the disk and the jet. The system of filaments in the center of M87 (Sparks, Ford \\& Kinney 1993; SFK) may also be causally connected to the disk. For example, the filaments extending $\\sim$17$''$ (1200 pc) to the NW at PA $\\sim 315^{\\circ}$ are blue shifted with respect to systemic velocity and show dust absorption implying they are on the near side of M87 as is the jet. These two findings led SFK to conclude that these filaments are streamers of gas flowing away from the center of M87 rather then falling into it. The images in F94 (see also Ford \\& Tsvetanov, this volume, FT98) show an apparent connection between at least some of the larger scale fillaments and the ionized nuclear disk.  Direct spectroscopic evidence for an outflow was found recently. Several UV and optical absorption lines from neutral and very mildly ionized gas were measured in the FOS spectrum of the nucleus (Tsvetanov et al.\\ 1998; T98). These lines are broad (FWHM $\\sim 400$ \\kms) and blue shifted by $\\sim 150$ \\kms\\ with respect to M87's systemic velocity implying both an outflow and turbulence. In addition, non-circular velocity components -- both blue and red shifted -- were found at several locations in the disk (F96, FT98), and observed emission lines are much broader than the expected broadening due to the Keplerian motion accross the FOS aperture. All these properties are best understood if a bi-directional wind from the disk were present. This wind may be an important mechanism for removing angular momentum from the disk to allow accretion through the disk onto the central black hole.  Whatever the physical conditions in the disk it is important to map its morphology in detail. The first \\HST\\ images (F94) have hinted that a spiral pattern could be present, but the signal-to-noise was too low for a definitive conclusion.  In this paper we present deep, fully sampled diffraction limited narrow band images of the nuclear region in M87.  We use these images to characterize the ellipticity, brightness distribution, and morphology of the disk. In this paper we adopt a distance to M87 of 15 Mpc, corresponding to a scale of 1$''$ = 73 pc.  \\vspace{-2mm} ",
        "conclusions": ""
    },
    "9803/astro-ph9803099_arXiv.txt": {
        "abstract": "We present a catalog of 200 clusters of galaxies serendipitously detected in 647 \\ROSAT\\/ PSPC high Galactic latitude pointings covering 158 square degrees.  This is one of the largest X-ray selected cluster samples, comparable in size only to the \\ROSAT\\/ All-Sky Survey sample of nearby clusters (Ebeling et al.\\ 1997). We detect clusters in the inner 17.5\\arcmin\\ of the \\ROSAT\\/ PSPC field of view using the spatial extent of their X-ray emission.  Fluxes of detected clusters range from $1.6\\times10^{-14}$ to $8\\times10^{-12}\\,$\\ergs\\ in the 0.5--2~keV energy band. X-ray luminosities range from $10^{42}~$erg~s$^{-1}$, corresponding to very poor groups, to $\\sim5\\times10^{44}\\,$erg~s$^{-1}$, corresponding to rich clusters. The cluster redshifts range from $z=0.015$ to $z>0.5$. The catalog lists X-ray fluxes, core-radii, spectroscopic redshifts for 73 clusters and photometric redshifts for the remainder. Our detection method, optimized for finding extended sources in the presence of source confusion, is described in detail. Selection effects necessary for a statistical analysis of the cluster sample are comprehensively studied by Monte-Carlo simulations.  We have optically confirmed 200 of 223 X-ray sources as clusters of galaxies.  Of the remaining 23 sources, 18 are likely false detections arising from blends of unresolved point X-ray sources, and for 5 we have not obtained deep CCD images.  Above a flux of $2\\times10^{-13}\\,$\\ergs, 98\\% of extended X-ray sources are optically confirmed clusters. The $\\log N - \\log S$ relation for clusters derived from our catalog shows excellent agreement with counts of bright clusters derived from the \\emph{Einstein}\\/ Extended Medium Sensitivity Survey (Henry et al.\\ 1992) and \\ROSAT\\/ All-Sky Survey (Ebeling et al.\\ 1997). At fainter fluxes, our $\\log N - \\log S$ relation agrees with the smaller-area WARPS survey (Jones et al.\\ 1998). Our cluster counts appear to be systematically higher than those from a 50~deg$^2$ survey of Rosati et al.\\ (1998). In particular, at a flux of $2\\times10^{-13}\\,$\\ergs, we find a surface density of clusters of $0.57\\pm0.07$ per square degree, which is a factor of 1.3 more than found by Rosati et al. This difference is marginally significant at the $\\sim 2$ sigma level. The large area of our survey makes it possible to study the evolution of the X-ray luminosity function in the high luminosity range inaccessible with other, smaller area \\ROSAT\\/ surveys. ",
        "introduction": "Clusters of galaxies are among the most important objects for cosmological studies.  Models of large scale structure formation such as CDM, predict that the abundance of clusters is determined by the spectrum of primordial perturbations and cosmological parameters $\\Omega$ and $\\Lambda$. Observations of clusters at different redshifts can be used to constrain these parameters (e.g., White \\& Rees 1978, Kaiser 1986, White, Efstathiou, \\& Frenk 1993, Henry \\& Arnaud 1991, Viana \\& Liddle 1996, Henry 1997). Following a different approach, observations of the Sunyaev-Zel'dovich effect (Sunyaev \\& Zel'dovich 1972) in a large sample of distant clusters can be used for a direct measurement of the distance to these clusters, and thus provide the values of $H_0$ (e.g., Birkinshaw, Hughes, \\& Arnaud 1991) and~$q_0$.  Up until the present, the largest samples of distant clusters resulted from optical surveys that searched for enhancements in the surface density of galaxies (e.g., Postman et al.\\ 1996). This method suffers seriously from projection effects (e.g., van Haarlem et al.\\ 1997). Distant clusters found by such techniques as galaxy concentrations around distant radio sources (Dickinson 1996) or ``dark'' lenses (Hattori et al.\\ 1997) cannot be considered as statistical samples.  Of all methods for detecting distant clusters, X-ray surveys are the least sensitive to projection, because the X-ray emission is proportional to the square of the density of the hot gas, which must be compressed in a deep potential well for us to detect it. It is noteworthy that unlike optical, X-ray surveys have the possibility of finding interesting objects such as ``fossil'' clusters in which almost all galaxies have merged to form a cD galaxy (Ponman et al.\\ 1994), and hypothetical ``failed'' clusters in which galaxy formation was suppressed (Tucker et al.\\ 1995). To date, the largest published sample of distant X-ray selected clusters is that from the \\emph{Einstein}\\/ Extended Medium Sensitivity Survey (EMSS; Goia et al.\\ 1990, Stocke et al.\\ 1991). However, because of the relatively high flux limit, the EMSS sample contains only 6 clusters at $z>0.5$.  Finding clusters in X-rays is complicated by their rarity among other types of sources.  A comparison of the $\\log N - \\log S$ relations for all sources (Hasinger et al.\\ 1993a) and clusters (this work) shows that at a flux of $10^{-14}\\,$\\ergs\\ in the 0.5--2~keV band, clusters comprise not more than 10--20\\% of the total source population. The large amount of optical identification work needed for cluster selection can be greatly reduced if they are searched for among spatially extended X-ray sources. Even at $z=1$, a rich cluster with a core-radius of 250~kpc has an angular radius of $>20\\arcsec$, which still can be resolved with the \\ROSAT\\/ PSPC on-axis. Detection of extended sources requires new analysis techniques.  Even if the spatial extent is not used for cluster selection, special detection techniques are needed because clusters at $z\\approx 0.2-0.3$ are 3--4 times broader than the \\ROSAT\\/ PSPC point spread function.  The idea of selecting distant cluster samples from various \\ROSAT\\/ surveys was pursued by different groups in the past few years.  Rosati et al.\\ (1995, 1998) searched for clusters in long exposure ($>15$~ksec) \\ROSAT\\/ PSPC pointed observations with a total area of 50~deg$^{2}$, using optical identifications of all extended X-ray sources found by wavelet transform analysis. Their sample consists at present of 70 clusters.  The Wide Angle \\ROSAT\\/ Pointed Survey (WARPS, Scharf et al.\\ 1997, Jones et al.\\ 1998) uses the Voronoi Tessellation and Percolation technique to detect both point-like and extended sources, followed by optical identifications of all sources. The WARPS cluster sample consists at present of 46 clusters found in \\ROSAT\\/ pointings with exposures $>8$~ksec, covering 16.2~deg$^{2}$.  A small sample of 15 clusters at $0.3<z<0.7$ was identified by the SHARC survey (Collins et al.\\ 1997). The RIXOS cluster sample (Castander et al.\\ 1995) consists of 13 clusters, detected using a technique which was optimized for point sources. Their results on cluster evolution appear to contradict other \\ROSAT\\/ surveys (Collins et al.\\ 1997), probably because the point source detection algorithm had a low efficiency for detecting extended cluster emission.  Finally, important information about the surface density of clusters at very low fluxes is provided by several very deep \\ROSAT\\/ pointings in which complete optical identifications are performed (e.g.\\ McHardy et el.\\ 1997).  Note that because of the small area, none of the aforementioned surveys is able to study the luminosity function of distant clusters above $3\\times10^{44}$~ergs~s$^{-1}$, where the deficit of high redshift EMSS clusters was reported (Henry et al.\\ 1992).   In this paper, we present a sample of distant clusters selected from 647 \\ROSAT\\/ PSPC observations of high Galactic latitude targets, covering a solid angle of 158 square degrees, a factor of three larger than the largest of the other \\ROSAT\\/ surveys. The source catalog includes 200 optically confirmed clusters, and thus is one of the largest X-ray selected samples, comparable in size only to the \\ROSAT\\/ All-Sky Survey sample of nearby clusters (Ebeling et al.\\ 1997).  We detect cluster candidates as extended X-ray sources using the wavelet decomposition technique described in this paper and Maximum Likelihood fitting of the surface brightness distributions to determine the significance of the source extent. We then identify only significantly extended sources with optical follow-up observations. Optical observations confirm that 90\\% of our sources are indeed clusters of galaxies.  Various selection effects such as the fraction of clusters which remain unresolved or undetected, are studied using extensive Monte-Carlo simulations. Comparison of the $\\log N - \\log S$ relation for clusters derived from our and other \\ROSAT\\/ surveys shows that our cluster counts at the bright end are in excellent agreement with those from the \\ROSAT\\/ All-Sky Survey sample of Ebeling et al.\\ (1997). At a flux of $2\\times10^{-13}\\,$\\ergs, our $\\log N - \\log S$ relation agrees well with the WARPS survey (Jones et al.\\ 1998), but is somewhat higher than that found by Rosati et al.\\ (1998).  Cluster size and flux estimates throughout the paper use $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0=0.5$. All X-ray fluxes and luminosities are reported in the 0.5--2~keV energy band. ",
        "conclusions": "We present a catalog of 200 clusters detected as extended X-ray sources in 647 \\ROSAT\\/ PSPC observations covering a solid angle of 158 square degrees. To detect these sources, we used a novel detection algorithm combining a wavelet decomposition to find candidate extended sources and Maximum Likelihood fitting to evaluate the statistical significance of the source extent.  Optical identifications demonstrate a high success rate of our X-ray selection: 90\\% of detected sources in the total sample, and 98\\% in the bright subsample are optically confirmed as clusters of galaxies. We present X-ray parameters of all detected sources and spectroscopic or photometric redshifts for optically confirmed clusters. Extensive Monte-Carlo simulations of our source detections are used to derive the sky coverage of the survey necessary for a statistical study of X-ray properties of our clusters. We present the $\\log N - \\log S$ relation derived from our cluster catalog.  This relation shows a general agreement with other, smaller area surveys.  In a subsequent paper (Vikhlinin et al.\\ 1998) we use this sample to constrain the evolution of cluster luminosities and radii at high redshift."
    },
    "9803/astro-ph9803050_arXiv.txt": {
        "abstract": "We investigate the application of neural networks to the automation of MK spectral classification.  The data set for this project consists of a set of over 5000 optical (3800--5200\\,\\AA) spectra obtained from objective prism plates from the Michigan Spectral Survey.  These spectra, along with their two-dimensional MK classifications listed in the Michigan Henry Draper Catalogue, were used to develop supervised neural network classifiers. We show that neural networks can give accurate spectral type classifications (\\sig68 = 0.82 subtypes, \\sigrms = 1.09 subtypes) across the full range of spectral types present in the data set (B2--M7). We show also that the networks yield correct luminosity classes for over 95\\% of both dwarfs and giants with a high degree of confidence.  Stellar spectra generally contain a large amount of redundant information. We investigate the application of Principal Components Analysis (PCA) to the optimal compression of spectra. We show that PCA can compress the spectra by a factor of over 30 while retaining essentially all of the useful information in the data set. Furthermore, it is shown that this compression optimally removes noise and can be used to identify unusual spectra.  This paper is a continuation of the work done by von Hippel et~al.\\ (1994) (Paper I)\\nocite{vonhippel_94a}. ",
        "introduction": "The MK classification of stellar spectra (Morgan, Keenan \\& Kellman 1943\\nocite{morgan_43a}; Keenan \\& McNeil 1976\\nocite{keenan_76a}; Morgan, Abt \\& Tapscott 1978\\nocite{morgan_78a}) is an important tool in stellar and galactic astrophysics.  In addition to providing fundamental stellar information it was, for example, central to the discovery of nearby Galactic spiral arms (Morgan, Sharpless \\& Osterbrock 1952; Morgan, Whitford \\& Code 1953).\\nocite{morgan_52a}\\nocite{morgan_52a}  MK classification is usually performed by a trained expert visually matching the overall appearance of a spectrum to the `closest' MK standard spectrum.  Such a qualitative method of classification suffers from subjective decisions and may differ from person to person: what is deemed as `close' by one person may not be `close' for another.  In addition, visual classification is very time consuming, with an expert classifying a few $10^5$ stars in a dedicated lifetime. Spectra collected from large spectral surveys, often as a by-product of other surveys (e.g.\\ the Sloan Digital Sky Survey (Kent 1994)\\nocite{kent_94a}) will have to be classified by automated means. Thus if stellar classification is to continue to be useful to the astronomical community, it has to be made faster and put on a more quantitative and objective basis.  In this paper we investigate the application of neural networks to the MK classification of optical stellar spectra.  The so-called `supervised' neural networks used in this project are implemented to yield an accurate mapping between a data domain (the stellar spectra) and a classification domain (the MK classifications).  While visual classifiers have mentally determined this mapping, they have not quantified it.  This mapping is, however, present intrinsically in a large set of classified spectra.  The neural network's resultant classification criteria will be essentially equivalent to the human's criteria. However, whereas a human's criteria may vary from adverse physiological and psychological factors such as health and mood, the network will retain a consistent set of classification criteria.  We will also demonstrate how the technique of Principal Components Analysis (PCA) can be used to optimally compress the spectra.  This has a number of advantages including the preferential removal of noise and an ability to isolate bogus spectra. Furthermore, using PCA-compressed spectra (rather than complete spectra) in the neural network classifiers leads to reduced training times and better convergence stability.  While MK classification will continue to be a useful tool to astronomers, it becomes increasingly desirable to obtain physical parameters (\\teff, \\logg, etc.)\\ for stars. Bailer-Jones et~al.\\ (1997b)\\nocite{bailerjones_97b} describe a neural network approach to the parametrization of stellar spectra by training a neural network on synthetic spectra. ",
        "conclusions": "We have produced a system for the automated two-parameter classification of stellar spectra over a wide range of spectral types (B2--M7) based on a large ($> 5000$), homogenous set of spectra. We have shown that we can achieve classification errors of \\sig68\\ = 0.82 subtypes (\\sigrms\\ = 1.09 subtypes) over this complete range of spectral subtypes. This result compares favourably with the intrinsic errors of \\sig68 = 0.63 subtypes in our training data.  Once a neural network has been trained, its classification results are completely reproducible. Moreover, the low values of their internal errors ($<0.4$ spectral subtypes) demonstrate that networks can be re-trained to give sufficiently consistent classifications.  We have achieved correct luminosity class classification for over 95\\% of dwarfs (class V) and giants (class III).  Results for luminosity class IV spectra were considerably worse.  It is believed that the data themselves could be a limiting factor and methods for improving these results were discussed. Despite the correlation in the data set between spectral type and luminosity class, it was demonstrated that the neural networks were using luminosity features to do dwarf-giant discrimination.  Network with two hidden layers performed considerably better ($\\approx 0.2$ subtypes) than ones with only one hidden layer.  The best classification results were achieved by tackling the spectral type and luminosity class problems separately, using continuous and probabilistic networks respectively.  We used Principal Components Analysis to compress the spectra by a factor of over 30 while retaining 96\\% of the variance in the data. It was shown that this compression predominantly removes noise. In addition the PCA preprocessing reduces the dimensionality of the data and can be used to filter out bogus spectral features or identify unusual spectra. However, PCA has the drawback that very weak or rare features will not be well-reconstructed.  More complex non-linear preprocessing schemes could no doubt be devised, but the strength of PCA is its analytic simplicity and its robustness.  The automated classifiers presented in this paper have been used to produce classifications for several thousand stars which do not have classifications listed in the MHD catallogue. These will be presented in a future paper (Bailer-Jones 1998)."
    },
    "9803/astro-ph9803320_arXiv.txt": {
        "abstract": "We have made a series of joint spectral fits for two blank fields, the Lockman Hole and the Lynx-3A field, where a significant amount of both {\\it ASCA} and {\\it ROSAT} PSPC data exist after thorough screenings. The {\\it ASCA} SIS, GIS and {\\it ROSAT} PSPC spectra from these fields have been  fitted simultaneously. Comparison at $E>1$ keV shows general agreement within 10\\% in the Lockman Hole data and a $20-30\\%$ disagreement in the Lynx-3A data, indicating remaining observation-dependent systematic problems. In both cases, satisfactory fits have been found for the overall 0.1-10 keV spectrum with an extragalactic power-law component (or a broken power-law component with steepening at $E<1$ keV), a hard thermal component with plasma temperature of $kT^{\\rm h}\\approx0.14$ keV and a soft thermal component  $kT^{\\rm s}\\approx0.07$ keV. ",
        "introduction": "\\label{sec:intr}  The global spectrum of the cosmic X-ray background is a primary piece of information for understanding its origin. The 3-50 keV CXRB spectrum observed with HEAO-1 A2 can be well described by a  $kT=40$ keV thin thermal plasma-like spectral shape (Marshall et al. \\cite{marshall80}; Boldt \\cite{boldt87}), which can be approximated by a power-law with a photon index of $\\Gamma = 1.4$ in $E\\la 10 keV$. {\\it BBXRT} (Jahoda et al. \\cite{jahoda92}) and {\\it ASCA} (Gendreau et al. \\cite{gend_spec}; Ishisaki et al. \\cite{ishi98}) measurements show that this power-law component extends down to 1 keV, below which an excess is observed.  A number of authors report {\\it ROSAT} measurements in the 0.5-2 keV band (e.g. Hasinger \\cite{has92}; Georgantopoulos et al. \\cite{georg96}) and show about 30\\% larger flux than the Gendreau et al.'s (\\cite{gend_spec}) {\\it ASCA} SIS result at 1 keV, with different slopes (see Hasinger \\cite{has96} for review). The disagreement may be contributed by the differences in the position/solid angle of the measured sky, problems arising from incomplete modelings, and/or calibration problems.  In order to separate these effects, we have made a series of joint spectral fits of {\\it ROSAT} PSPC, {\\it ASCA} GIS, and {\\it ASCA} SIS spectra from two fields of the sky, where sufficient amount of blank-sky data exist after thorough screening. Because of the limited data meeting the criteria, the work presented in this paper is not intended to determine the current best estimate of the global CXRB spectrum, but rather a comparison of measurements among {\\it ASCA} and {\\it ROSAT} instruments in the same parts of the sky with consistent modelings.  In Sect.\\ref{sec:data}, we describe the {\\it ASCA} and {\\it ROSAT} data used in the analysis. Joint spectral fits are described in Sect.\\ref{sec:fit}. The results are discussed in Sect.\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  There is a bright variable source in LH with [1.2$\\pm$.2] and [2.5$\\pm$.7] $\\times 10^{-13} [{\\rm erg\\,cm^{-2}\\,s^{-1}}]$ for the 0.7-2 keV and 2-7 keV bands respectively (Ogasaka \\cite{oga_t}), consisting about 10\\% of the total ASCA fluxes in both bands. This source was much fainter during the PSPC observation. Thus one should decrease the GIS and SIS normalizations about $10\\%$ lower when comparing with the PSPC data. In this case, the agreement between PSPC and GIS is excellent and falls well within statistical errors of each other. For both LH and LX, the SIS data consistently show $\\approx 10\\%$ lower normalizations compared to GIS. This might be caused by incomplete calibration for the radiation damage of the SIS with the 4CCD mode, which can even exist at this level after a few months after the launch, when LH and LX were observed (Dotani et al. \\cite{dotani95}).  In the LX observation, a larger discrepancy exists. The GIS and SIS normalizations are lower than the PSPC value by  $\\approx20\\%$ and $\\approx30\\%$ respectively and slopes are shallower. The ASCA LX normalizations are also significantly lower than those of LH.  There is no variable source which can cause this amount of discrepancy in LX.  The fact that the 0.1-10 keV fit still show the disagreement of the normalization (see $N_{\\rm G}$ in B2) shows that this is not a modeling problem (e.g. leak of the $E>1$ keV excess with the PSPC energy resolution). One possible explanation is the LTE (e.g Snowden et al. \\cite{snow94}), which is usually apparent in the $E<0.5$ keV channels of the PSPC data, but sometimes extends above 1 keV when the activity is high. There may also be an over-subtraction of the NXB background from the {\\it ASCA} data. Furthermore, the instruments are not looking at exactly the same part of the sky. Due to the stray-light and PSF of the ASCA instruments, $\\approx 40\\%$ of the GIS/SIS flux comes from outside of the designated FOV (estimated using our ray-tracing program).  These effect may also contribute to this discrepancy.  \\begin{figure}[t] \\psfig{file=lo_comp2.ps,width=\\hsize,angle=270} \\caption[]{The PSPC and GIS $E\\,I(E)$ spectra (using a two power-law model as an appropriate smooth function for unfolding purpose) of LH are shown and compared with previous measurements: the thick solid bowtie is from Hasinger (\\cite{has92}); the dot-dashed bowtie from Georgantopoulos et al. (\\cite{georg96}), both used {\\it ROSAT} PSPC. The dotted line is from rocket measurements (McCammon \\& Sanders \\cite{maccamon90}). The long-dashed horn is from an {\\it ASCA} SIS measurement by Gendreau et al. \\cite{gend_spec} and the thin solid bowtie is a joint {\\it ROSAT} PSPC/{\\it ASCA} SIS analysis of QSF3 by Chen et al. (\\cite{chen97}) for $E>1$ keV. The thick solid line represents the HEAO-1 A2 measurement by Marshall et al. (\\cite{marshall80}).} \\label{fig:lo_comp} \\end{figure}  The best consistency for a certain region of the sky from this work is seen for the PSPC and GIS data on LH. Thus it is instructive to compare the observed spectra of these with previous CXRB measurements. The comparison is  shown in Fig. \\ref{fig:lo_comp}. Fig. \\ref{fig:lo_comp} shows a large excess at $E\\sim 0.6$ keV on the PSPC data, inconsistent with Gendreau et al.'s {\\it ASCA} SIS data. This may be partially due to the low Galactic column density of LH. The LH GIS data for $E\\ga 2$ keV are above the HEAO-1 A2 and Gendreau et al. (\\cite{gend_spec}) SIS values. We note, however, that about 10 \\% of source fluctuation is expected over this small area. Since the ASCA LH data contains a bright source ($\\sim 5\\times 10^{-13} {\\rm erg\\,s^{-1}\\,cm^{-2}}$ in 2-10 keV), this field should be one of the brighter ones. We also note, however, that an integration from the brightest source in the field to the faintest source excluded in the collimator experiments (e.g. $\\approx2\\times 10^{-11} {\\rm erg\\,s^{-1}\\,cm^{-2}}$ in 2-10 keV for HEAO-1 A2 measurement by Marshall et al.) would add $\\approx10\\%$ of intensity. A thorough treatment of source fluctuations using a larger area and comparing spectra with appropriate source removal will be presented in a future paper.  In summary, a close look at {\\it ROSAT} and {\\it ASCA} spectra for the same regions of the sky have revealed systematic errors caused by response calibration problems and non cosmic background subtraction of up to $\\approx 20-30\\%$ for one set of observations. These probably caused the reported disagreements between {\\it ASCA} and {\\it ROSAT} measurements (Hasinger 1996), while modelings and sky selection can also contribute. We have obtained a fair description of the CXRB spectrum over 0.1-10 keV range cosisting of a extragalactic power-law component (either single or broken below 1 keV), hard and soft thermal components with a satisfactory fit to all instruments."
    },
    "9803/astro-ph9803116_arXiv.txt": {
        "abstract": " ",
        "introduction": "Session B.3 received a partisan organisation, and was divided into sections corresponding to the main paradigms pervading modern cosmology. Three sub-sessions were allocated to cover inflationary cosmology, pre-big-bang scenarios, and topological defects in cosmology. Anything not fitting into these topics makes up the last section, covering miscellaneous topics.  Below I briefly review the current status in each subject covered in a sub-session after which I summarise the talks presented. These summaries reflect my personal understanding of the talks, and I apologise to the speakers if I accidentally missed the entire point. ",
        "conclusions": ""
    },
    "9803/astro-ph9803184_arXiv.txt": {
        "abstract": "Blue- and red-shifted Hydrogen and Helium satellite recombination lines have recently been discovered in the optical spectra of at least two supersoft X-ray sources (SSSs), RX~J0513-069 and RX~J0019.8+2156, and tentatively also in one short-period cataclysmic variable star (CV), the recurrent nova T~Pyx. These features are thought to provide evidence for the presence of highly collimated jets in these systems. No similar spectral signatures have been detected in the spectra of other short-period cataclysmic variables, despite a wealth of existing optical data on these systems. Here, we ask if this apparent absence of ``jet lines'' in the spectra of most CVs already implies the absence of jets of the kind that appear to be present in the SSSs and perhaps T~Pyx, or whether the current lack of jet detections in CVs can still be ascribed to observational difficulties.  To answer this question, we derive a simple, approximate scaling relation between the expected equivalent widths of the observed jet lines in both types of systems and  the accretion rate through the disk, $EW(line) \\propto  \\dot{M}_{acc}^{\\frac{4}{3}}$. We use this relation to predict the strength of jet lines in the spectra of ``ordinary'' CVs, i.e. systems characterized by somewhat lower accretion rates than T~Pyx. Making the assumption that the features seen in T~Pyx are indeed jet lines and using this system as a reference point, we find that if jets are present in many CVs, they may be expected to produce optical satellite recombination lines with EWs of a few hundredths to a few tenths of Angstroms in suitably selected systems. A similar prediction is obtained if the SSS RX~J0513-069 is used as a reference point. Such equivalent widths are small enough to account for the non-detection of jet features in CVs to date, but large enough to allow them to be detected in data of sufficiently high quality, if they exist. ",
        "introduction": "\\label{introduction}  One of the most intruiging empirical connections that has emerged from observations of accretion-powered objects on all astrophysical scales is that between accretion disks and powerful bipolar outflows and jets. Indeed, the occurrence of mass loss in this form has been established in members of essentially all classes of (presumed) disk-accretors, including close binary systems such as X-ray binaries, supersoft X-ray sources (SSSs) and non-magnetic cataclysmic variable stars (CVs), young stellar objects such as T-Tauri and FU~Orionis stars, and active galactic nuclei and quasars.  However, a complete theoretical understanding of this empirical disk-wind (or disk-jet) connection has not yet been achieved. Recent reviews of the subject with references to the relevant observational literature may be found in Livio \\shortcite{livio4,livio5}.  A potentially vital clue to the origin of mass loss from accretion disk systems is provided by the fact that in at least some members of all but one class of objects the outflow collimation appears to be very tight, with half-opening angles, $\\theta_{max}$, of no more than a few degrees (we will refer to such flows as {\\em jets} hereafter). The sole exception to this rule are CVs (see however below). In these systems, the presence of mass loss has nevertheless been clearly established, based on the shapes and eclipse behavior of the ultraviolet (UV) resonance lines (e.g. Drew 1991), but the inferred collimation of the corresponding outflows is weak ($\\theta_{max} \\gtappeq 45^o$; Shlosman, Vitello \\& Mauche 1997; Knigge \\& Drew 1997).  The apparent absence of jets in CVs may hold the key to an improved theoretical understanding of the disk-wind connection. If confirmed, any ``generic'' model (in the sense of being applicable to more than a particular class of objects) for the origin of mass loss from accretion disks must be able to explain why this mass loss takes the form of jets in all systems but CVs. This would be quite a restrictive constraint, particularly in the light of two recent observational developments.  The first of these is the detection of blue- and red-shifted satellite emission features to the optical Hydrogen and Helium recombination lines in the SSSs RX~J0513-069, RX~J0019.8+2156 and (possibly) CAL~83 \\cite{crampton1,crampton2,southwell1,southwell2,becker1}. These features cannot be attributed to any other ionic species and are therefore thought to arise in some kind of bipolar outflow. The combination of small width and large displacement from line center exhibited by these satellite lines further suggests that the corresponding outflows are very highly collimated, i.e. that they are jets (c.f. the prototypical jet system SS~433; Vermeulen~1993).  If this identification is correct (which we will assume throughout this paper) it is significant, because SSSs and CVs are extremely similar types of objects: both are semi-detached binary systems in which a Roche-lobe filling secondary star transfers material via an accretion disk onto a white dwarf (WD) primary. The main difference between SSSs and CVs is thought to be the rate at which this mass transfer proceeds. In SSSs, the accretion rate is believed to be high enough ($\\dot{M}_{acc} \\sim 10^{-7} - 10^{-6}$~M$_{\\sun}$~yr$^{-1}$) to fuel steady nuclear hydrogen burning on the surface of the WD (e.g. van den Heuvel et a. 1992). By contrast, the highest accretion rates encountered in CVs are about one order of magnitude lower than this ($\\dot{M}_{acc} \\sim 10^{-8}$~M$_{\\sun}$~yr$^{-1}$) and thus insufficient to initiate steady nuclear burning on the WD. Instead, the accretion can result in shell flashes which are responsible for nova outbursts. This difference is one of the factors that led Livio (1997a) to propose that the formation of powerful jets (as opposed to more weakly collimated bipolar outflows) may {\\em require} the presence of an additional wind/energy source at the center of the accretion disk.  The second recent observational finding of significance in the present context is the tentative detection of similar H$\\alpha$~satellite lines in the optical spectrum of {\\em one} CV, the recurrent nova T~{Pyx} \\cite{shahbaz1}. It is important to stress that no similar features have ever been detected in any other short-period CV, despite the fact that optical spectra of high enough quality to detect satellite lines of similar strength as in T~Pyx (which, on average, have an equivalent width of about 1~\\AA) should be available for many of them. While it is possible that in some cases they might have been overlooked (previous studies would not have expected to see such features), it would nevertheless appear that satellite lines of the strength seen in T~Pyx are not a common feature among CVs. For example, Shahbaz et al. (1997) did not detect similar satellite lines in the spectrum of another recurrent nova, U~Sco. (It should be noted that one-sided satellite lines have been seen in the spectra of a few CVs, e.g. S193, V795~Her, BT~Mon [Szkody 1995; Haswell et al. 1994; Seitter 1984]; however, these probably arise in other types of high velocity flows in these systems.)  The physical similarities between CVs and SSSs, on the one hand, and the observational similarity between the satellite features in T~Pyx and those in SSSs such as RX~J0513-069, on the other, suggest that T~Pyx may also harbor a well collimated jet. A fundamental assumption we adopt in the present paper is that this is indeed the case. Note that, under this assumption, we would expect T~Pyx's binary inclination to be higher than that of the SSSs listed above, since in T~Pyx the ratio of satellite line widths to their displacements from line center is somewhat smaller. A high inclination for T~Pyx ($i \\sim 70^o$) is also indicated if one demands that the displacement of the satellite features from line center should roughly correspond to the escape velocity from the WD, as expected if they are formed in a jet (c.f. Shahbaz et al. 1997; Livio 1997a). It is acknowledged, however, that Shahbaz et al. (1997) also derived an inclination estimate from the peak-to-peak separation of the double-peaked H$\\alpha$ line core, which, by contrast, turns out to be very low ($i \\sim 10^o$). While this estimate should be regarded as a lower limit \\cite{shahbaz1}, the case for a well collimated jet in T~Pyx would have to be critically reexamined if the inclination of this system were shown to be $\\ltappeq 50^o$ in the future. The analytic scaling relation derived in Section~2 would of course nevertheless be valid and, we believe, useful, even if T~Pyx should eventually turn out not to contain a collimated jet.  Actually, the existence of a jet in T~Pyx is not entirely unexpected, since it is thought that recurrent novae in general, and T~Pyx in particular, are characterized by accretion rates that are higher (i.e. $\\dot{M}_{acc} \\gtappeq 10^{-8}$~M$_{\\sun}$~yr$^{-1}$) than those in other types of CVs, such as dwarf-novae (DNe) and nova-like variables (NLs). In fact, Webbink et al. (1987) have suggested that intermittent nuclear burning might take place on the surface of the WD in T~Pyx even during quiescence, i.e. between nova outbursts. If so, Livio's (1997a) hypothesis could be used to reconcile the presence of jet lines in SSSs and T~Pyx with the absence of similar features in the spectra of other CVs.  The main goal of the present paper is to determine whether such reconciliation is actually required at present. Thus, we will ask whether the apparent absence of jet lines in the existing optical spectra of most CVs, particularly NLs and DNe in outburst (i.e. systems in which an optically thick accretion disk is present but no nuclear burning occurs) already implies the {\\em absence} of jets of the kind seen in the SSSs and (possibly) T~Pyx, or whether the current lack of jet detections in CVs can still be ascribed to observational difficulties. In our attempt to answer this question, we will take as our fundamental working hypothesis that the main difference between CVs and the SSSs (as well as perhaps T~Pyx) -- the presence of an extremely hot WD at disk center -- is irrelevant to the formation of the observed jets. The corollary of this hypothesis is that CVs must actually drive jets of the same kind that appear to be present in the SSS and T~Pyx. Our goal is to estimate the expected equivalent widths of CV jet lines within the framework of this hypothesis. In so doing, we will try to make as few references as possible to specific models for the jet formation and line emission mechanisms, and instead use only simple and general physical arguments.  As already noted above, under our working hypothesis, jets akin to those in the SSSs must actually be present in CVs. This is not in direct conflict with the relatively wide outflow opening angles that have been inferred for CV winds from modeling the UV resonance lines: these spectral features probe only the near-disk regime of the outflow -- out to at most a few hundred $R_{WD}$ -- leaving collimation at larger distances as a distinct possibility. ",
        "conclusions": "\\label{BLA}  \\subsection{A prediction for the strength of jet lines in CVs} \\label{BLA2}  We are now in a position to use Relation~(\\ref{final2}) to predict the jet line EWs we would expect to see in ``ordinary'' CVs, according to our working hypothesis. Since T~Pyx {\\em is} in fact a CV, its system parameters are more typical of ``normal'' CVs than are those of the SSSs exhibiting jet lines. It is therefore preferable to use T~Pyx as a reference point in making predictions for other CVs, because the ratios of the relevant factors in Relation~(\\ref{final2}) will be closer to unity and the associated uncertainties will be smaller. However, in Section~\\ref{BLA3} below we will check whether Relation~(\\ref{final2}) is at least consistent with the observed accretion rate and EW ratios of T~Pyx and the SSS RX~J0513-069. (See also the note at the end of the manuscript, in which we show that the prediction derived in this section by using T~Pyx as a reference datum is consistent with what is be obtained if RX~J0513-069 is used instead.)  Concerning T~Pyx, Webbink et al. (1987) give $\\dot{M}_{acc}(T~Pyx) \\gtappeq 10^{-8}$~M$_{\\sun}$~yr$^{-1}$ based on the short recurrence time scale of its eruptions and $\\dot{M}_{acc}(T~Pyx) \\sim 5 \\times 10^{-8}$~M$_{\\sun}$~yr$^{-1}$ based on its optical colors. Here, we adopt the latter estimate, which assumes that the optical light is due to the accretion disk, rather than to direct and/or reprocessed light from a hot WD. This is in line with our working hypothesis that a wind/energy source at disk center is not required to drive jets from accretion disks. T~Pyx's orbital period is thought to be $P_{orb}(T~Pyx) \\simeq 1.8$~hrs \\cite{schaefer1}, placing the system below the period gap. As noted in Section~\\ref{introduction}, the inclination angle of T~Pyx is not well constrained observationally, although the appearance of the satellite recombination lines themselves suggests a high value $i(T~Pyx) \\simeq 70^o$ if these features are formed in a well collimated jet. Finally, the equivalent widths of the H$\\alpha$~jet lines in T~Pyx can be measured from the data of Shahbaz et al. (1997) and turn out to be about EW(T~Pyx) $\\simeq 1$~\\AA~on average, with the strongest feature in any one of the observing epochs reaching about twice this value (Shahbaz, private communication).  We now need to make some assumptions about the typical properties of ``ordinary'' CVs. The form of Relation~(\\ref{final2}) shows that if jets are present in these objects, the associated satellite recombination lines are likely to be strongest in systems with high mass accretion rates, short orbital periods and high inclinations (though not so high as to shift the jet lines into the line core). Since it would be sufficient to falsify our working hypothesis for CVs with these properties, we adopt $\\dot{M}_{acc}(CV) \\simeq 1 \\times 10^{-8}$~M$_{\\sun}$~yr$^{-1}$ (appropriate to NLs and DNe in outburst), $P_{orb}(CV) \\simeq P_{orb}(T~Pyx)$, and $i(CV) \\simeq i(T~Pyx)$. (We note in passing that there are actually no well-established non-magnetic NL variables with periods shorter than 3.2~hrs, although there is a fair number of DNe with $P_{orb} \\leq 2$~hrs.)  We can now use Relation~(\\ref{final2}) to predict the jet line EWs we expect to see in this most favorable sub-group of ``ordinary'' CVs, according to our working hypothesis. To this end, we take the ratio of the two separate relations (one for the normal CVs, one for T~Pyx), solve for $EW(CV)$ and substitute our adopted parameters. This yields $EW(CV) \\simeq 0.1^{+0.1}_{-0.04}$~\\AA. While it may be possible to increase the upper limit implied by this result somewhat -- by taking $P_{orb}(CV)$ to be shorter or $i(CV)$ to be higher, for example -- it is clear that CV jet lines, if they exist, would be at best marginally detectable in typical optical spectra. As a result, we are forced to conclude that our working hypothesis and its corollary -- that a hot central object is inessential to the formation of jets and that CVs do in fact drive jets -- {\\em cannot} yet be ruled out.  Let us take a step back at this point to make it clear what we are -- and are not -- claiming. We started by adopting the working hypothesis that the formation of powerful jets does {\\em not} require the presence of an additional energy source at disk center. As a corollary, we assumed that ``ordinary'' CVs harbor the same kind of jets that may be present in T~Pyx (as indicated by the satellite recombination lines that are observed in that object). We then showed that based on these assumptions one can derive a simple scaling law which can be used to predict the expected strength of these jet lines in the optical spectra of ordinary CVs. The predicted jet line EWs for these systems turned out be very small, even for objects with nearly optimal system parameters. We therefore concluded that the lack of jet line detections in the optical spectra of ordinary CVs is not yet in conflict with our working hypothesis, i.e. that jets {\\em may} be present in ordinary CVs. Note that we do {\\em not} claim to have shown that ordinary CVs actually {\\em do} contain jets. After all, an inability to falsify a hypothesis does not prove it. Summarized succinctly, our conclusion is that {\\em the non-detection of jet lines in existing optical spectra of ``ordinary'' CVs should not yet be taken to imply that these systems cannot harbor collimated jets}.  Two further points need to be made regarding this statement. First, even though we have been unable to rule out the presence of jets in ``ordinary'' CVs on the basis of existing data, the predicted EWs of a few hundredths up to a few tenths of Angstroms may not be beyond the reach of high resolution, high signal-to-noise optical spectra. Thus we strongly encourage observers to search for the signatures of jets in the spectra of appropriately selected CVs.  Second, it was assumed above that T~Pyx's optical continuum is dominated by the radiation field emitted by a standard accretion disk. However, if (intermittent) nuclear burning really does take place in T~Pyx, the surface of the WD at the center of the disk will be extremely hot. It is therefore worth asking whether (some of) the optical continuum could actually be direct or reprocessed radiation emitted by the WD, and what effect this may have on our conclusions.  A numerical calculation similar to that described following Relation~(\\ref{continuum1}) in Section~\\ref{scale} shows that direct light from the WD is unlikely to be of any importance, even if the temperature of the WD is as high as a few times $10^5$~K, and the accretion rates as low as $10^{-8}$~M$_{\\sun}$~yr$^{-1}$. To judge the potential significance of reprocessed WD radiation, we rely on the recent work of King (1997), who derived a simple condition that can be used to estimate the relative importance of dissipation and reprocessing in a CV accretion disk. More specifically, King (1997) showed that reprocessing of WD radiation will begin to have a dominant effect on the local disk temperature if $L_{WD} \\gtappeq 2.5 L_{acc} (1-\\beta)^{-1}$, where $L_{WD}=4\\pi R_{WD}^2 \\sigma T_{WD}^4$ and $L_{acc} = GM_{WD} \\dot{M}_{acc}/R_{WD}$ are the WD and total accretion luminosities, respectively, and $\\beta$ is the albedo of the disk surface. To give a numerical example, we note that if reprocessing is assumed to be efficient ($\\beta \\simeq 0$), the temperature distribution in a disk around a $1M_{\\sun}$~WD accreting at a rate of $\\dot{M}_{acc} = 10^{-8} M_{\\sun}$~yr$^{-1}$ will be dominated by reprocessing if $T_{WD} \\gtappeq 2 \\times 10^5$~K.  If reprocessed WD radiation is in fact contributing significantly to T~Pyx's optical continuum, then our previous prediction for the strength of jet lines in other CVs no longer applies, since our continuum scaling law, Relation~(\\ref{continuum}), ceases to be valid. Qualitatively, the effect of this will be to increase the predicted EWs significantly, since (a) the adopted accretion rate for T~Pyx is almost certainly an overestimate in this case, and (b) the extra contribution to the continuum that is ultimately due to nuclear burning on the WD (and not to accretion) is making the jet lines appear weaker than if only the disk were producing the continuum. Quantitatively, these effects can be corrected for by multiplying the predicted EWs by a factor of $f_{\\dot{M}}^{2}$, where $f_{\\dot{M}}>1$ is the factor by which T~Pyx's accretion rate has been overestimated. The dependence on $f_{\\dot{M}}$ {\\em squared} arises because the part of correction (a) that is related to the scaling of the continuum flux with accretion rate exactly cancels correction (b). This leaves the scaling of the line luminosity with accretion rate as the only relevant factor.  It is now easy to see that if irradiation is very important in T~Pyx and has caused us to overestimate the accretion rate by a significant amount, then the non-detection of jet lines in the spectra of other CVs does become inconsistent with the presence of jets in these systems. Indeed, if $f_{\\dot{M}}\\gtappeq 4$, then even the previously derived lower limit of 0.06~\\AA~on the jet line EWs in (suitably selected) CVs becomes as large as 1~\\AA~and thus comparable to the strength of the same features in T~Pyx. In practical terms, this means that studies of T~Pyx aimed at deriving $T_{WD}$ (or, more precisely, $L_{WD}/L_{acc}$) for this system may provide yet another way to falsify our working hypothesis observationally in the future.   \\subsection{The scaling relation applied to T~Pyx and the SSSs} \\label{BLA3}  Given that the jets in T~Pyx and the SSSs are presumably of the same type, it is natural to try and use these systems to check our scaling relation for the jet line EWs. Unfortunately, the accretion rates of the relevant SSSs are only poorly constrained and, in addition, disk irradiation by the hot WD is likely to be very strong in the SSSs. As a consequence, a rigorous test of Relation~(\\ref{final2}) via this route is not possible. However, we will nevertheless proceed to apply our scaling relation to T~Pyx and the SSS RX~J0513-069, partly to illustrate these problems, and partly to perform at least a rough consistency check.  In their study of RX~J0513-069, Southwell et al. (1996) state that $\\dot{M}_{acc} \\sim 10^{-5}$~M$_{\\sun}$~yr$^{-1}$ is required if the optical luminosity of this system is to be ascribed entirely to a standard accretion disk. An accretion rate this high is of the order of the Eddington value, and Southwell et al. (1996) therefore conclude that it is almost certainly an overestimate. They argue that irradiation of the disk and secondary star, as well as (perhaps) direct light from the hot WD are likely to contribute significantly to the optical light. Consequently, they prefer a lower value of about $10^{-6}$~M$_{\\sun}$~yr$^{-1}$ for the accretion rate. To make progress in the face of this uncertainty, we will adopt the higher value to start with and then check {\\em a posteriori} what value this implies for the correction factor $f_{\\dot{M}}^2$. Regarding RX~J0513-069's other relevant parameters, Southwell et al. (1996) give values of $P_{orb} \\simeq 18$~hrs for the orbital period, and, based on the mass function of the system, $i \\simeq 10^o$ for the inclination.  Adopting these parameters for RX~J0513-069, and using the same parameters as above for T~Pyx, we would predict a best-bet ratio for the EWs of the jet lines in these two systems of about 50 (in favor of the SSS). Now, Southwell et al. (1996) measure the equivalent widths of the blue and red H$\\alpha$~jet satellite lines in RX~J0513-069 to be EW(SSS,blue) $\\simeq 1.6$~\\AA~and EW(SSS,red) $\\simeq 2.6$~\\AA, respectively. Thus the actual ratio of the jet line EWs in RX~J0513-069 and T~Pyx is only about 2. If we interpret this as a result of disk irradiation in the SSS, then the correction factor $f_{\\dot{M}}^2 \\simeq 25$ and $f_{\\dot{M}} \\simeq 5$. Consequently, we would predict the true accretion rate in RX~J0513-069 to be about $\\dot{M}_{acc} \\sim 2 \\times 10^{-6}$~M$_{\\sun}$~yr$^{-1}$, which is in line with the value of $10^{-6}$~M$_{\\sun}$~yr$^{-1}$ preferred by Southwell et al. (1996). We do not attach too much weight to this apparent consistency, because there are large observational uncertainties associated with the ratios constructed from two of the relevant parameters (accretion rate and inclination). Moreover, the accretion rate and orbital period ratios are so large for these systems that the theoretical uncertainties expressed by the ``errors'' in Relation~(\\ref{final2}) also become rather large.  It is finally interesting to consider briefly the implications of adopting the complement of our working hypothesis. Specifically, we may ask whether a consistent physical picture capable of accounting for the relative strengths of the jet lines in T~Pyx and RX~J0513-069 can also be found if we assume that a hot, central object is in fact present in both systems and is crucial for driving the observed jets. To answer this question, we take $\\dot{M}_{jet} \\propto L_{WD}$ and assume the extreme case of $L_{WD} >> L_{acc}$. The disk is then still quite likely to dominate the optical flux (c.f. the numerical estimates for the direct WD contribution given previously), but its local temperature distribution will be dominated by irradiation, not dissipation (see Section~\\ref{BLA2}). Since the disk will be extremely hot in this case, we may further assume that the optical waveband lies on the Rayleigh-Jeans tail of the disk spectrum now, i.e. $F_{opt} \\propto L_{disk}^{1/4} \\propto L_{WD}^{1/4}$ (the latter holds since $L_{disk}$ is now dominated by $L_{WD}$). We can then replace the dependence on $\\dot{M}_{acc}$ in Relation~(\\ref{final2}) with one on $L_{WD}$, giving $EW(line) \\propto L_{WD}^{7/4} \\propto T_{WD}^7$. Adopting again an EW ratio of 2 for RX~J0513-069 and T~Pyx, we find that $T_{WD}(SSS) \\simeq 2~T_{WD}(T~Pyx)$ in this simplistic picture, if the remaining parameter dependences in Relation~(\\ref{final2}) are assumed to stay unchanged.  This reasonable looking result should of course not be taken too seriously. However, the moral of this simple calculation is that it is certainly possible to account for the jet line EW differences between T~Pyx and RX~J0513-069 in the context of a model in which the presence of an energy source at disk center {\\em is} a crucial ingredient in driving the observed jets. This prompts us to stress again that our analysis in this paper has only shown that the presence of jets in CVs should not be ruled out simply because no jet lines have so far been detected in the optical spectra of these systems. We have by no means demonstrated that jets are actually present, or are even likely to be present, in ordinary CVs.  {\\bf Note added:} After this paper was accepted for publication, we received a draft of a work by Margon \\& Deutsch, in which it is argued that the satellite lines seen in T~Pyx are in fact due to [N~{\\sc ii}]~$\\lambda\\lambda$6548,6584 and are formed in the complex velocity field of T~Pyx's nova shell(s). While the analytic scaling relation we derived in Section~2 retains its validity (and, we believe, usefulness) if this interpretation turns out to be correct, the same is not true for the prediction we made for the jet line EWs in ordinary CVs (since this is based on the assumption that T~Pyx's satellite lines are jet features). The best we can do in this case is to derive a new prediction by scaling down directly from one of the SSSs to CVs. To do this, we use Southwell~et al.'s (1997) inclination, orbital period and accretion rate estimates for RX~J0513-069 ($i\\simeq 10^o$; $P_{orb} \\simeq 18$~hrs; $\\dot{M}_{acc}(apparent)\\sim 10^{-5}$~M$_{\\sun}$~yr$^{-1}$ with $f_{\\dot{M}}=10$) and, as before, parameters appropriate to an optimally selected, ``ordinary'' CV ($i\\simeq 70^o$; $P_{orb} \\simeq 1.8$~hrs; $\\dot{M}_{acc} \\sim 10^{-8}$~M$_{\\sun}$~yr$^{-1}$). Ignoring limb-darkening ($\\eta(i)=1$), we obtain a new prediction of $EW(CV) \\sim 0.3$~\\AA. Even though the uncertainties on this number are substantial and hard to quantify (see Section~2.3), this estimate still suggests it would be premature to rule out the presence of jets in CVs completely at this stage. \\footnote{Note that if jets are present in CVs but jet lines are not seen in T~Pyx, the {\\em absence} of the latter would have to be attributed to one or both of the following: (i) T~Pyx's inclination is much lower than $70^o$; (ii) irradiation is increasing the brightness of the accretion disk in T~Pyx substantially.} We therefore suggest that an optical survey of suitably selected CVs to search for jet lines is called for, regardless of the nature of the satellite lines in T~Pyx."
    },
    "9803/astro-ph9803071_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "These two measurements of D/H in QSO absorption systems are the best and most robust measures to date.  Deuterium has been identified and analyzed in a number of other QSO absorption systems\\cite{dhother} We have found another two systems which place a strong upper limit on D/H at D/H $< 10^{-4}$.  Combined with the two measurements described above, the four independent systems support a low primordial abundance of deuterium, and together give D/H = 3.4 $\\pm \\, 0.3 \\times 10^{-5}$. If this represents the primordial value, nucleosynthesis calculations from standard BBN models with three light neutrinos give $\\eta =  5.1 \\pm \\, 0.3 \\times 10^{-10}$ and $\\Omega_b\\,h_{100}^2 = 0.019 \\pm \\, 0.001$.  The constraints from D/H can be utilized to constrain cosmological models, quantify dark matter both in unobserved baryons and non-baryons, specify the zero point for models of deuterium evolution\\cite{chemevol}, test directly the predictions of standard BBN by comparing with other light element abundances\\cite{hat}, and limit the amount of small scale entropy fluctuations in the early universe\\cite{jed}."
    },
    "9803/astro-ph9803247_arXiv.txt": {
        "abstract": "We have obtained spectra of the Galactic center at energies 400--600 keV from high-resolution data acquired by the TGRS Ge spectrometer on board the {\\em WIND\\/} mission during 1995--1997.  The data were obtained using an on-board occulter, and are relatively free from systematics and backgrounds.  Analysis of the spectra reveals a well-resolved electron-positron annihilation line at 511 keV and the associated continuum due to annihilation via positronium formation.  Measurements of the line width and the continuum-to-line ratio allow some constraints to be placed on the interstellar sites where annihilation occurs. ",
        "introduction": "The line at 511 keV from the annihilation of electrons and positrons in the region of the Galactic center (GC) is the best-studied line in $\\gamma$-ray astronomy.  Over 20 years of observations (reviewed by Tueller 1993) have established that there is an extensive diffuse line source of total intensity $\\sim 2 \\times 10^{-3}$ photon cm$^{-2}$ s$^{-1}$.  This source has recently been mapped in considerable detail by OSSE on board the {\\em Compton Observatory\\/} (Purcell et al. 1997), which has revealed a third spatial component in addition to the well-known Galactic disk and bulge components.  This new component is extended and is centered at $l = -2^{\\circ}$, $b = +9^{\\circ}$, well above the Galactic plane. It is unclear whether there are any point sources superimposed on this diffuse distribution; recent results do not show any variability in the flux.  The line is known to be narrow and centered at 511 keV (Leventhal, MacCallum, \\& Stang 1978). The annihilation spectrum also includes a lower-energy continuum arising from $3 \\gamma$ annihilation via the formation of positronium (Ps).  In principle, spectral lines contain much information about the physical conditions in the line formation region.  The next step in the study of the 511 keV line will be to extract the information contained in the line profile and in the ratio of line to Ps continuum amplitudes.  The key requirement is for sensitive long-term measurements with fine spectral resolution.  The measurements described above were mostly made with low-resolution scintillator detectors.  In this paper, we describe observations made over more than 2 years with the high-resolution Ge spectrometer TGRS on board the {\\em WIND\\/} spacecraft. ",
        "conclusions": "The results of our measurements of the annihilation spectrum are given in Table 1.  These results supersede the preliminary measurement made by Teegarden et al. (1995), which reported a 511 keV line flux $1.64 \\times 10^{-3}$ photon cm$^{-2}$ s$^{-1}$.  Two features of the earlier analysis contributed to this overestimate.  First, the line flux was obtained from the count rate in the narrow 506--516 keV band, without fitting the shape of the spectrum, and is thus overestimated by including the other spectrum components (Fig. 1).  Second, the present analysis uses an improved model of the instrument spectral response, instead of simply dividing by photopeak effective area as was done in the earlier work.  \\subsection{Comparison with OSSE results}  We calculated the flux, dimension and centroid of the OSSE model of the GC 511 keV emission (Purcell et al. 1997) as folded through the TGRS occulted response (Table 1).  Since the occulter passes close to the centers of two of the three spatial components of the model (exactly crossing the GC, and $2.8^{\\circ}$ from the high-latitude feature) a test of these model features becomes possible in principle. The measurements of the flux and centroid are in good agreement; the offset of the TGRS centroid measurement from the GC is in the same direction as the offset of the OSSE centroid due to the new high-latitude feature, but is also compatible with the GC.  The spatial extension found by TGRS slightly exceeds that found by OSSE, but to draw any conclusion from this would be premature since improved modeling of the occultation response is required.  Our results agree with OSSE in finding a lack of variability on 90 d time-scales (Fig. 3).  \\subsection{Source physics: Line width}  Our measurements of the total line width $\\sigma_{tot}$, and of the background line widths, are shown in Fig. 4.  The interpolated instrument intrinsic width $\\sigma_{inst}$ is also shown.  It is clear that $\\sigma_{tot}$ at all times exceeds $\\sigma_{inst}$; this is necessary if the cosmic line width is to be obtained from $\\sigma_{tot}^{2} = \\sigma_{inst}^{2} + \\sigma_{gal}^{2}$.  The result (Table 1) is somewhat narrower than the average of four balloon measurements by the GRIS Ge detector (Leventhal et al. 1993), but the difference is not very significant.  The line width reflects the convolved widths of components due to different annihilation mechanisms predominating in different phases of the ISM.  These mechanisms were treated by Guessoum, Ramaty \\& Lingenfelter (1991).  They may be divided into two classes.  Firstly, annihilation by charge-exchange in flight produces a broad line (FWHM 6.4 keV), and is predominant in cold molecular clouds.  The second class contains all other processes, which produce lines narrower than the instrument resolution.  We can therefore hope to test two alternative suggestions by Guessoum et al. --- annihilation occurring uniformly in all phases of the ISM, and otherwise-uniform annihilation excluding cold clouds.  We therefore repeated our analysis under the assumption that the width $\\sigma_{gal}$ had two components, a broad component of width 6.4 keV, and a narrow unresolved component.  Instead of line width, we now have the amplitude of the 6.4-keV broad component as a fitted parameter.  The spectra were fitted equally well by this model; though there were no significant improvements in the $\\chi^2$ values, we hope this procedure yields physical insight into the meaning of $\\sigma_{gal}$.  Assuming the presence of a 6.4-keV broad line component, we found that $11$\\%$\\pm 9$\\% of the total line intensity was due to this broad line. This is much closer to the prediction when positrons are excluded from molecular cloud cores (in which case the broad line contributes only 11\\%: Guessoum et al. 1991) than to the maximum predicted broad-line contribution of 59\\% when positrons penetrate all phases of the ISM equally.  \\subsection{Source physics: Positronium fraction}  The fraction $f$ of positrons which annihilate through the formation of Ps can be written $f = 2/[2.25(I_{511}/I_{Ps})+1.5]$, where $I_{511}$ and $I_{Ps}$ are the line and Ps continuum intensities (Brown \\& Leventhal 1989); our result from Table 1 is $f = 0.94 \\pm 0.04$.\\footnote{ The uncertainty does not include those systematic errors in $I_{511}$ and $I_{Ps}$ in Table 1 which are positively correlated.}  This is in good agreement with the most recent OSSE result $f = 0.97 \\pm 0.03$ (Kinzer et al. 1996).  However, predicted values from annihilation in most of the phases of the ISM cluster in the range $f \\sim 0.9$--1.0, so small discrepancies in measured $f$ may be important.  Our result falls roughly in the middle of this range, and is consistent with annihilation in cold molecular clouds ($f = 0.9$: Brown, Leventhal \\& Mills 1986), the warm neutral or ionized ISM ($f = 0.9$--0.95: Bussard, Ramaty \\& Drachman 1979), or any combination of these (Guessoum et al. 1991).  It is not compatible with annihilation in the hot phase, nor with any scenario in which grains are important sites of annihilation. These two statements are in fact equivalent, since in the hot phase grains become the most important location for annihilation in the absence of H atoms.  The corresponding value of $f$ is expected to be very low ($\\le 0.5$: Guessoum et al. 1991).  \\subsection{Summary}  We have measured the 511 keV line from the GC and also the Ps continuum associated with it during 1995--1997.  Our values for the intensities of these features agree with the most recent OSSE measurements.  Our preliminary results for the spatial distribution of the line are consistent with the OSSE mapping, but require further analysis of the instrument response.  The 511 keV line is resolved, and, if a specific model for its width is assumed (an underlying broad component from annihilation through charge-exchange in flight), then our result favors a scenario in which annihilation in cold molecular clouds is suppressed.  Our measurement of the Ps fraction $f$ from the Ps continuum is consistent with this, and suggests further that annihilation in the hot phase of the ISM is of minor importance."
    },
    "9803/astro-ph9803137_arXiv.txt": {
        "abstract": "We report the discovery of \\lya\\ emission from a galaxy at $z=5.34$, the first object at $z>5$ with a spectroscopically confirmed redshift. The faint continuum emission (${\\rm m_{AB}(8000{\\rm \\AA})\\approx 27}$), relatively small rest-frame equivalent width of the emission line ($W_{\\rm Ly\\alpha}^{rest}\\approx 95$\\AA), and limits on the \\ion{N}{5}/\\lya\\ ratio suggest that this is a star--forming galaxy and not an AGN. The star--formation rates implied by the UV continuum emission and the \\lya\\ emission are (in the absence of dust extinction) fairly modest ($\\sim 6~h_{50}^{-2}\\ \\Msun~yr^{-1}$ for \\qnot=0.5). The continuum luminosity is similar to that of sub-$L^*_{1500}$ star--forming galaxies at $z\\sim3$, and the width of the \\lya\\ line yields an upper limit to the mass of $< 2.6\\times 10^{10}\\Msun$.  The strong emission line detected in this low-luminosity galaxy provides hope for the discovery of higher luminosity primeval galaxies at redshifts $z>5$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803301_arXiv.txt": {
        "abstract": "Using the Hubble Space Telescope and WFPC2 we have imaged the central 20pc of the giant H~II region 30 Doradus nebula in three different emission lines.  The images allow us to study the nebula with a physical resolution that is within a factor of two of that typical of ground based observations of Galactic H~II regions.  We present a gallery of interesting objects within the region studied.  These include a tube blown by the wind of a high velocity star and a discrete H~II region around an isolated B star.  This small isolated H~II region appears to be in the midst of the champagne flow phase of its evolution.  Most of the emission within 30 Dor is confined to a thin zone located between the hot interior of the nebula and surrounding dense molecular material.  This zone appears to be directly analogous to the photoionized photoevaporative flows that dominate emission from small, nearby H~II regions.  For example, a column of material protruding from the cavity wall to the south of the main cluster is found to be a direct analog to elephant trunks in M16.  Surface brightness profiles across this structure are very similar to surface brightness profiles taken at ground based resolution across the head of the largest column in M16.  The dynamical effects of the photoevaporative flow can be seen as well.  An arcuate feature located above this column and a similar feature surrounding a second nearby column are interpreted as shocks where the photoevaporative flow stagnates against the high temperature gas that fills the majority of the nebula.  The ram pressure in the photoevaporative flow, derived from thermal pressure at the surface of the column, is found to balance with the pressure in the interior of the nebula derived from previous x-ray observations.  By analogy with the comparison of ground and HST images of M16 we infer that the same sharply stratified structure seen in HST images of M16 almost certainly underlies the observed structure in 30 Dor.  30 Doradus is a crucial case because it allows us to bridge the gap between nearby H~II regions and the giant H~II regions seen in distant galaxies.  The real significance of this result is that it demonstrates that the physical understanding gained from detailed study of photoevaporative interfaces in nearby H~II regions can be applied directly to interpretation of giant H~II regions.  Stated another way, interpretation of observations of giant H~II regions must account for the fact that this emission arises not from expansive volumes of ionized gas, but instead from highly localized and extremely sharply stratified physical structures. ",
        "introduction": "30 Doradus is a giant ionized complex in the Large Magellanic Cloud (LMC), located at a distance of 51.3 kpc (eg. Panagia et al 1991).  The nebula is centered on a dense cluster of newly formed stars, the most dense component of which is called R136.  The nebula itself is more than 180 parsecs across, qualifying it as a smaller member of the elite class of nebulae termed Giant Extragalactic H~II Regions (GEHR's).  If 30 Doradus was placed at the distance of the Orion Nebula from the Earth, it would appear to be more than 20 degrees across, and would fill more than 4\\% of the night sky.  The central cluster is very dense and is comprised of several hundred OB stars with a small number of W-R stars (Hunter et al 1995b).  The integrated ultraviolet flux from this cluster is intense: more than fifty times that being produced in the center of the Orion Nebula (Campbell et al 1992).  Radiation from the cluster combined with strong stellar winds from the most massive stars in the cluster has eroded a large cavity in the nearby molecular complex, producing the nebula we see today.  Hunter et al 1995b showed that the majority of the stars in the cluster were formed in a single star formation event more than 2-3 million years ago.  The census performed yielded a ``head count\" of more than 3000 stars with more than 300 OB stars capable of producing the intense UV radiation and stong stellar winds responsible for forming and shaping the H~II regions we observe in galaxies.  The level of star formation exhibited by the 30 Doradus region and the neighboring LMC complex are the closest example of starburst-like star formation.  As such we are getting a unique view of the star formation environment in the middle of an ongoing starburst.  The average reddening along the line of sight to the Large Magellanic Cloud and 30 Doradus is very low (Panagia et al 1991). However, within several H~II complexes in the LMC comparison between optical and radio measurements suggest a large variation in the local reddening. Kennicutt \\& Hodge 1986 found a variation in these estimates between 0 and 1 magnitude in A$_{V}$.  Hunter et al 1995b also found substantial variation in the reddening across the face of the 30 Doradus nebula, and derived a mean estimate for this reddening of 1.4 magnitudes in A$_{V}$ at 555 nm, and 0.8 magnitudes at 814 nm.  For the purposes of this paper we will adopt an extinction of 1 magnitude in A$_{V}$ for the emission lines we observe.  The 30 Doradus nebula plays a key role in our understanding of H~II regions. Nearby regions are close enough for the physical processes at work within the nebula to be studied in detail.  The work by Hester et al 1996 (hereafter H96) on M16, for example, shows that emission within the nebula arises predominantly within a narrow region at the interface between the H~II region and the molecular cloud.  They follow Hester 1991 in describing this thin region as a photoionized photoevaporative flow.  However, an H~II region like M16 is tiny in comparison with giant H~II regions, and no giant H~II regions are close enough to allow the stratified ionization structure of the photoevaporative flow to be studied directly.  30 Doradus alone offers an opportunity to bootstrap the physical understanding of small nearby H~II regions into the context of the giant regions seen in distant galaxies.  In this paper we present Hubble Space Telescope images of the ionization structure we observe around the central cluster R136.  The wealth of spatial information contained in these pictures is daunting to consider, but we attempt to summarize the most telling points by selecting and presenting several examples of distinct structures around the field of view that provide insight into how the interface with the local gas and dust is evolving.  In \\S 2 we discuss the observations themselves and the general structure of the nebula, as well as presenting full-field mosaics of the data.  In \\S 3 we discuss the conditions apparent in the ionized hydrogen along the walls of the H~II region cavity, as well as comparing the ionization structure we observe with models we derive from the H$\\alpha$ surface brightness. ",
        "conclusions": "We have presented high resolution narrow-band imagery of the 30 Doradus nebula. There are many interesting localized structures within the nebula, a number of which appear to be associated with winds and UV from stars that are not part of the main 30 Doradus cluster.  However the majority of the emission from the nebula is due to photoionization by the flux from the central cluster.  This emission is largely concentrated in thin regions located at the interface between dense molecular material and the shock-heated interior of 30 Dor.  At the resolution of the {\\it HST} data we find that the structure in 30 Doradus is remarkably similar to what is seen in ground-based observations of nearby H~II regions.  This similarity is not surprising given that despite an overall difference in scale, locally the physical conditions in 30 Doradus are not much different than those found in smaller H~II regions.  We demonstrate this point above by focussing on one particular region in 30 Doradus and showing that it is a very direct analog of the Galactic H~II region M~16. Taking this same argument a step further we are lead to the conclusion that underlying the observed structure in M~16 is the same sort of extremely localized and sharply stratified structure seen in the {\\it HST} images of M~16.  Thus, even though the 30 Doradus nebula spans hundreds of parsecs, the emission from this giant H~II region arises largely in the same sorts of sharply stratified photoionized photoevaporative flows seen in nearby H~II regions.  The 30 Doradus nebula is a crucial case.  The fact that at {\\it HST} resolution 30 Doradus is so similar to ground based images of nearby H~II regions has allowed us to bootstrap our physical understanding based on detailed study of nearby regions into the physical context of a giant H~II region surrounding a massive young cluster.  Similarly, preliminary analysis of {\\it HST} images of more distant giant H~II regions suggests that they compare favorably with 30 Dor when that nebula is viewed at the same physical resolution.  This indicates that conditions in 30 Doradus are probably typical of those found in these distant H~II regions as well.  Bootstrapping first from nearby H~II regions to 30 Doradus in this paper, and we anticipate from 30 Doradus to more distant regions in later work, we are approaching the conclusion that the emission from giant H~II regions megaparsecs distant is determined by the physics of photoevaporative flows in which relevant physical scales can be as small as 100 AU or less.  The significance of this work lies in the conclusion that the detailed study of nearby, well-resolved H~II regions is directly applicable to distant giant H~II regions in much the same way that an understanding of radiative shocks that is tested in nearby supernova remnants can be applied in a variety of contexts in which the shock itself is not resolved.  Viewed from a different perspective, interpretation of observations of distant giant H~II regions must take into account the fact that much of this emission arises not in vast expanses of ionized or even clumpy gas, but instead in well defined and highly stratified photoevaporative flows localized to the surfaces of molecular clouds."
    },
    "9803/astro-ph9803315_arXiv.txt": {
        "abstract": "The process of molecule formation in the primordial gas is considered in the framework of Friedmann cosmological models from redshift $z=10^4$ to $z=0$. First, a comprehensive analysis of 87 gas phase reaction rates (both collisional and radiative) relevant in the physical environment of the expanding universe is presented and critically discussed.  On this basis, calculations are carried out of the abundance of 21 molecular species as function of redshift, for different values of the cosmological parameters $\\Omega_0$, $\\eta$ and $H_0$, evaluating consistently the molecular heating and cooling due to H$_2$, HD and LiH molecules.  One of the major improvements of this work is the use of a better treatment of H recombination that leads to a reduction of a factor 2--3 in the abundance of electrons and H$^+$ at freeze-out, with respect to previous studies. The lower residual ionization has a negative effect on the chemistry of the primordial gas in which electrons and protons act as catalysts in the formation of the first molecules.  We find that in the standard model ($h=0.67$, $\\eta_{10}=4.5$, $\\Omega_0=1$ and [D/H] $=4.3\\times 10^{-5}$), the residual fractional ionization at $z=1$ is $[{\\rm e/H}]=3.02\\times 10^{-4}$, and the main molecular species fractional abundances $[{\\rm H}_2/{\\rm H}]=1.1\\times 10^{-6}$, $[{\\rm HD/H}_2]=1.1\\times 10^{-3}$, $[{\\rm HeH}^+/{\\rm H}]=6.2\\times 10^{-13}$, $[{\\rm LiH}^+/{\\rm H}]=9.4\\times 10^{-18}$ and $[{\\rm LiH/LiH}^+]=7.6\\times 10^{-3}$.  We devise a reduced chemical network that reproduces with excellent accuracy the numerical results of the complete model and allows to  follow the chemical compositions and the thermal properties of a primordial gas in the presence of an external radiation field. Finally, we provide accurate cooling functions of H$_2$, HD and LiH in a wide range of density and temperature that can be conveniently used in a variety of cosmological applications. ",
        "introduction": "The study of molecule formation in the post-recombination epoch has grown considerably in recent years. Saslaw \\& Zipoy (1967) and Peebles \\& Dicke (1968) were the first to realize the importance of gas phase reactions for the formation of the simplest molecule, H$_2$. They showed that trace amounts of molecular hydrogen, of order 10$^{-6}$--10$^{-5}$, could indeed form via the intermediaries species H$_2^+$ and H$^-$ once the radiation field no longer contained a high density of photons with energies above the threshold of dissociation (2.64 and 0.75 eV, respectively).  The presence of even a trace abundance of H$_2$ is of direct relevance for the cooling properties of the primordial gas which, in its absence, would be an extremely poor radiator: cooling by Ly-$\\alpha$ photons is in fact ineffective at temperatures $\\la 8000$~K, well above the matter and radiation temperature in the post-recombination era.  Since the evolution of primordial density fluctuations is controlled by the ability of the gas to cool down to low temperatures, it is very important to obtain a firm picture of the chemistry of the dust-free gas mixture, not limited to the formation of H$_2$, but also to other molecules of potential interest. In this regard, Lepp \\& Shull (1984) and Puy et al. (1993) have computed the abundances of H$_2$, HD and LiH as a function of redshift for various cosmological models. Although the final abundances of H$_2$ and HD agree in the two calculations, their evolution with redshift is markedly different, since the epoch of formation varies by a factor of $\\sim 2$. Also, the LiH abundance shows a large discrepancy of about two orders of magnitude.  More recently, Palla et al. (1995) have analyzed the effects on the chemistry of the pregalactic gas of a high primordial D abundance in the light of the controversial results obtained towards high redshift quasars (see e.g. Tytler \\& Burles 1997).  They found that the abundance of H$_2$ is rather insensitive to variations in the cosmological parameters implied by a factor of $\\sim 10$ enhancement of primordial [D/H], while HD and LiH abundances vary by larger amounts. However, the abundance of LiH, obtained with simple estimates of the radiative association rate, was largely overestimated. Because of the potential relevance of the interaction of LiH molecules with the cosmic background radiation (CBR) (Maoli et al. 1994), a proper treatment of the lithium chemistry was necessary.  Dalgarno et al. (1996) and Gianturco \\& Gori Giorgi (1996a,b) provided accurate quantum-mechanical calculations of the main reaction rates.  The chemistry of lithium in the early universe has been then studied by Stancil et al. (1996) and Bougleux \\& Galli (1997). Finally, useful reviews of the chemistry of the early universe can be found in Dalgarno \\& Lepp (1987), Black (1991), Shapiro (1992), and Abel et al. (1997).  The latter two studies, in particular, focus on the nonequilibrium H$_2$ chemistry in radiative shocks which is thought to be of primary importance during the gravitational collapse of density fluctuations (see also Anninos et al. 1997).  In spite of such a wealth of specific studies, a comprehensive analysis of the subject and a critical discussion of the reaction paths and rates are still lacking. To overcome this limitation, in this paper we present a complete treatment of the evolution of {\\em all} the molecular and atomic species formed in the uniform pregalactic medium at high redshifts ($z<10^4$).  The structure of the paper is as follows: in Sect.~2 we describe the H, D, He, and Li chemistry, with a critical discussion of the most important rates; the evolutionary models are presented in Sect.~3, and the results for the standard model and the dependence on the cosmological parameters are given in Sect.~4; Sect.~5 introduces a minimal model which highlights the dominant reactions for the formation of H$_2$, HD, HeH$^+$, LiH and LiH$^+$; a comparison with the results of previous studies is given in Sect.~6, and the conclusions are summarized in Sect.~7. Also, the Appendix provides the collisional excitation coefficients for HD and H$_2$ and cooling function of H$_2$, HD and LiH which are needed for the computation of the thermal evolution of the primordial gas. ",
        "conclusions": "The main results of the present study can be summarised as follows:  1) We have followed the chemical evolution of the primordial gas after recombination by computing the abundances of 21 species, 12 atomic and 9 molecular, by using a complete set of reaction rates for collisional and radiative processes. The rates which are critical for a correct estimate of the final molecular abundances have been analysed and compared in detail.  2) One of the major improvements of this work is the use of a better treatment of H recombination that leads to a reduction of a factor 2--3 in the abundance of electrons and H$^+$ at freeze-out, with respect to previous studies. The lower residual ionization has a negative effect on the chemistry of the primordial gas in which electrons and protons act as catalysts in the formation of the first molecules.  3) In the standard model ($h=0.67$, $\\eta_{10}=4.5$, $\\Omega_0=1$ and [D/H] $=4.3\\times 10^{-5}$), the residual fractional ionization at $z=1$ is $[{\\rm e/H}]=3.02\\times 10^{-4}$, and the main molecular species have fractional abundances $[{\\rm H}_2/{\\rm H}]=1.1\\times 10^{-6}$, $[{\\rm HD/H}_2]=1.1\\times 10^{-3}$, $[{\\rm HeH}^+/{\\rm H}]=6.2\\times 10^{-13}$, $[{\\rm LiH}^+/{\\rm H}]=9.4\\times 10^{-18}$ and $[{\\rm LiH/LiH}^+]=7.6\\times 10^{-3}$.  4) As for molecular hydrogen, its final abundance does not depend on the model parameters, making this molecule a poor diagnostic of cosmological scenarios. The largest uncertainty resides in the accurate knowlwedge of the photodissociation rate of H$_2^+$. A detailed treatment of the reaction kinetics of this reaction would be required.  5) We have presented a minimal model consisting of 11 reactions for H$_2$, 6 for HD, 3 for HeH$^+$ and 14 for LiH which reproduces with excellent accuracy the results of the full chemical network, regardless of the choice of the cosmological parameters.  6) Finally, we have computed accurate expressions for the cooling functions of H$_2$, HD and LiH in a wide range of density and temperature that can be conveniently used in a variety of comological applications."
    },
    "9803/astro-ph9803123_arXiv.txt": {
        "abstract": "We have obtained  WFPC2 images of 256 of the nearest (z$\\leq$0.035) Seyfert 1, Seyfert 2, and starburst galaxies. Our 500-second broadband (F606W) exposures reveal much fine-scale structure in the centers of these galaxies, including dust lanes and patches, bars, rings, wisps and filaments, and tidal features such as warps and tails. Most of this fine structure cannot be detected in ground based images. We have assigned qualitative classifications for these morphological features, a Hubble type for the inner region of each galaxy, and also measured quantitative information such as 0.18 and 0.92 arcsecond aperture magnitudes, position angles and ellipticities where possible.  There is little direct evidence for unusually high rates of interaction in the Seyfert galaxies.  Slightly less than 10\\% of all the galaxies show tidal features or multiple nuclei. The incidence of inner starburst rings is about 10\\% in both classes of Seyfert galaxies. In contrast, galaxies with H II region emission line spectra appear substantially more irregular and clumpy, because of their much higher rates of current star formation per unit of galactic mass.  The presence of an unresolved central continuum source in our {\\it HST} images is a virtually perfect indicator of a Seyfert 1 nucleus as seen by ground-based spectroscopy. Fifty-two percent (52\\%) of these Seyfert 1 point sources are saturated in our images; we use their wings to estimate magnitudes ranging from 15.8 to 18.5. The converse is not universally true, however, as over a third of Seyferts with direct spectroscopic evidence for broad Balmer wings show no nuclear point source. These 34 resolved Seyfert 1's have fainter nonstellar nuclei, which appear to be more extinguished by dust absorption. Like the Seyfert 2's, they have central surface brightnesses consistent with those expected for the bulges of normal galaxies.  The rates for the occurrences of bars in Seyfert 1's and 2's and non-Seyferts are the same. We found one significant morphological difference between the host galaxies of Seyfert 1 and Seyfert 2 nuclei. The Seyfert 2 galaxies are significantly more likely to show nuclear dust absorption, especially in lanes and patches which are irregular or reach close to the nucleus. A few simple tests show that the difference cannot be explained by different average redshifts or selection techniques. It is confirmed by our galaxy morphology classifications, which show that Seyfert 1 nuclei reside in earlier type galaxies than Seyfert 2 nuclei. If, as we believe, this is an intrinsic difference in host galaxy properties, it would undermine one of the postulates of the strong unification hypothesis for Seyfert galaxies, that they merely appear different due to the orientation of their central engine. The excess galactic dust we see in Seyfert 2's may cause substantial absorption which obscures their hypothesized broad-emission-line regions and central nonstellar continua. This galactic dust could produce much of the absorption in Seyfert 2 nuclei which had instead been attributed to a thick dusty accretion torus forming the outer part of the central engine. ",
        "introduction": "Several causal connections have been proposed between an active galactic nucleus (AGN) and the host galaxy in which it resides. The principal ways in which the latter could affect the former are through influencing a) the formation of a nonstellar central engine; b) its fueling; and c) obscuring it from our view, (which can alter the central engine's appearance even if it is not physically affected.)  It is widely believed that active galactic nuclei (AGNs) are powered by non-spherical accretion onto massive black holes. This is partly because this model has the lowest fuel supply requirements: an AGN's luminosity is proportional to its mass accretion rate, which would be about 0.01 \\msun year$^{-1}$ for a bright Seyfert nucleus. It is not known how this rate of fuel supply can be brought from the host galaxy down to several thousand Schwarzschild radii (of order 10$^{17}$ cm for a ``typical\" Seyfert galaxy black hole mass of 10$^8$ \\msun (Malkan \\markcite{a60} 1983)) at which point viscous processes are supposed to drive the final accretion onto the black hole.  One speculation is that a close interaction with another galaxy can distort the galactic potential and disturb the orbits of gas clouds sufficiently to carry a significant mass of fuel into the galaxy's center (Shlosman \\etal\\markcite{a1}\\markcite{a2}1989, 1990, Hernquist and Mihos \\markcite{a3}1996). More indirect scenarios are also possible, in which a tidal galaxy interaction stimulates a burst of star formation which in turn stimulates nonstellar nuclear activity. A further possibility is that special conditions in isolated galaxies may trigger the feeding of fuel to an active nucleus, such as a bar instability. (Schwartz \\markcite{a31}1981; Shlosman, Frank, \\& Begelman \\markcite{a2}1990; Mulchaey and Regan \\markcite{a32}1997)  However, direct observational evidence that galaxy encounters stimulate the luminosity of an AGN has been ambiguous (Adams \\markcite{a4}1977, Petrosian \\markcite{a65}1983, Kennicutt and Keel \\markcite{a5}1984, Dahari \\markcite{a6}\\markcite{a7}1985a, 1985b, Bushouse \\markcite{a8}1986, Fuentes-Williams and Stocke \\markcite{a66}1988).  One difficulty is that the most dramatic morphological indications of the encounter may have subsided by the time that the newly injected fuel reaches the nucleus. In any case, the weak correlation between galaxy interactions and Seyfert activity is stronger for type 2 Seyferts than for type 1's.  Conversely, the presence of an AGN could alter the appearance of the central regions of its host galaxy, principally by its injection of substantial energy, both radiative and mechanical, over many millions of years. A further question is whether the particular type of active nucleus, Seyfert 1 or 2, is related to any property of the host galaxy.  We have therefore used the superior spatial imaging resolution of the post-repair {\\it Hubble Space Telescope} to make a snapshot survey of nearby active galaxies to investigate the morphological implications of different theories on the formation and fueling of AGN. ",
        "conclusions": "Our large sample of high-resolution images of the centers of nearby Seyfert 1, 2 and HII galaxies has allowed us to search for statistical differences in their morphologies.  The Seyfert galaxies do not, on average, resemble the HII galaxies. The latter have more irregularity and lumpiness associated with their high rates of current star formation. Conversely, none of the HII galaxies have the filaments or wisps which are sometimes seen in Seyfert 1 and 2 galaxies, and are evidently gas filaments photoionized by the active nucleus.  Sixty-three percent (63\\%) of the galaxies classified as Seyfert 1 have an unresolved nucleus, 52\\% of which are saturated. Some (6\\%) have such dominant nuclei that they would appear as ``naked quasars\" if viewed at somewhat higher redshifts. The presence of an unresolved nucleus, particularly a saturated one, is anti-correlated with an intermediate spectroscopic classification (such as Seyfert 1.8 or 1.9) and is also anti-correlated with the Balmer decrement.  This implies that those Seyfert 1's with weak nuclei in the PC2 images are extinguished and reddened by dust.  The vast majority of the Seyfert 2 galaxies show no central point source. In fact, the only two of these that do (IRAS 1832-594 and IC 4870) are mis-classified galaxies. If all Seyfert 2's actually harbor point-like continuum sources like those in Seyfert 1's, they are at least an order of magnitude fainter on average. In those galaxies without any detectable central point source (37\\% of the Seyfert 1's; 98\\% of the Seyfert 2's, and 100\\% of the H II's), the central surface brightnesses are statistically similar to those observed in the bulges of normal galaxies.  Seyfert 1's and 2's  both  show circumnuclear rings in about 10\\% of the galaxies. We identified strong inner bars as often in Seyfert 1 galaxies (27\\%) as in Seyfert 2 galaxies (22\\%). In some cases we see a strong assymetry of the dust absorption across the major axis, which allows us to infer which half of the disk is nearer to us: the side which more strongly absorbs the smooth light of the bulge behind it.  The Seyfert 2 galaxies are more likely than Seyfert 1's to show irregular or disturbed dust absorption in their centers as well as galactic dust lanes which pass very near their nuclei. They also, on average, tend to have latter morphological types than the Seyfert 1's. This difference remains in Seyfert 1 and 2 subsamples matched for redshift, [OIII] and radio luminosities.  It also holds true when we restrict our consideration to sub-samples of the data which are less biased by selection effects.  Thus it appears that the host galaxies of Seyfert 1 and 2 nuclei are {\\it not} intrinsically identical. A galaxy with more nuclear dust and in particular more irregularly distributed dust is more likely to harbor a Seyfert 2 nucleus. This indicates that the higher dust-covering fractions in Seyfert 2's are the reason for their spectroscopic classification: their compact Seyfert 1 nucleus may have been obscured by galactic dust. This statistical result contradicts the simplest and most popular version of the unified scheme for Seyfert galaxies. We suggest that the obscuration which converts an intrinsic Seyfert 1 nucleus into an apparent Seyfert 2 often occurs in the host galaxy hundreds of parsecs from the nucleus. If so, this obscuration need have no relation to a hypothetical fat dust torus surrounding the equator of the central engine. Also then the orientation of the central engine with respect to our line-of-sight does {\\it not} determine whether an active nucleus will appear to us as a Seyfert 1 or as a Seyfert 2."
    },
    "9803/astro-ph9803065_arXiv.txt": {
        "abstract": "Wind-blown bubbles, from those around massive O and Wolf-Rayet stars, to superbubbles around OB associations and galactic winds in starburst galaxies, have a dominant role in determining the structure of the Interstellar Medium. X-ray observations of these bubbles are particularly important as most of their volume is taken up with hot gas, $10^{5} \\ltsimm T (\\K) \\ltsimm 10^{8}$.  However, it is difficult to compare these X-ray observations, usually analysed in terms of single or two temperature spectral model fits, with theoretical models, as real bubbles do not have such simple temperature distributions. Spectral fits, and the properties inferred from them, will depend in a complex way on the true temperature distribution and the characteristics and limitations of the X-ray observatory used.  In this introduction to a series of papers detailing the {\\em observable} X-ray properties of wind-blown bubbles, we describe our method with which we aim to solve this problem, analysing a simulation of a wind-blown bubble around a massive star.  Our model is of a wind of constant mass and energy injection rate, blowing into a uniform ISM, from which we calculate X-ray spectra as would be seen by the {\\it ROSAT} PSPC. Analysing these spectra in the same way as a real observation would be, we compare the properties of the bubble as would be inferred from the {\\it ROSAT} data with the true properties of the bubble in the simulation.  We find standard spectral models yield inferred properties that deviate significantly from the true properties, even though the spectral fits are statistically acceptable, and give no indication that they do not represent to true spectral distribution. For example, single temperature spectral fits give best fit metal abundances only 4\\% of the true value. A cool bubble has best fit temperatures significantly higher than a bubble twice as hot. These results suggest that in any case where the true source spectrum does not come from a simple single or two temperature distribution the ``observed'' properties cannot naively be used to infer the true properties. In this situation, to compare X-ray observations with theory it is necessary to calculate the {\\em observable} X-ray properties of the model. ",
        "introduction": "Bubbles blown in the Interstellar Medium (ISM) by massive stars are a common astrophysical phenomenon. X-ray observations can provide information regarding the density, metal abundance, temperature, ionisation state and physical structure in the hot bubbles surrounding Wolf-Rayet (WR) and O stars (Wrigge, Wendker \\& Wisotski 1994), Luminous Blue Variables (LBV's) such as $\\eta$ Carinae (Corcoran \\etal 1995) and planetary nebulae (PN) (Kreysing \\etal 1992; Leahy, Zhang \\& Kwok 1994; Arnaud, Borkowski \\& Harrington 1996; Leahy \\etal 1996). On the larger scale, superbubbles are created by the sum of the winds and SN within OB associations (Belloni \\& Mereghetti 1994) and giant star forming regions in young starburst galaxies.  Superbubbles within starburst galaxies such as M82 eventually break out the galaxy to form spectacular galactic winds (Watson, Stanger \\& Griffiths 1984; Heckman, Armus \\& Miley 1987). In many cases the X-ray emission probes different regions of the object in question to that revealed by optical observations, increasing the importance of the X-ray data.  Analytic solutions to the development and structure of astrophysical bubbles must rely on simplifying assumptions, and increasingly attention has turned to the use of numerical hydrodynamics. These simulations have been enlightening with respect to the nonlinear processes occurring, with some degree of quantitative agreement with observation, but generally lack predictive power. This is partially due to the difficulty in comparing them with observations, in particular X-ray observations.  The problem is that X-ray observations are usually analysed by fitting a single or two temperature spectral model to the observed spectra (see for example the references above), and the best-fit results are used to infer the physical properties of the object.  However, for wind-blown bubbles such as those mentioned above, the true situation is more complex, and the results of the spectral fits may be influenced by, for example: projection of different physical regions along the line of sight; the presence of a wide range of temperatures; interstellar absorption; unknown or non-standard elemental abundances; non-ionisation equilibrium conditions; low numbers of observed photons and the limitations of the current X-ray telescope optics and detectors.  All of these make the interpretation of what is normally a one or two temperature spectral fit to the data difficult to relate to the underlying physical conditions, and conversely, make it difficult to predict the {\\em observable} properties of a model or simulation.  To our knowledge there has been no study of the influence of the complexities mentioned above on the best-fit properties of a spectral fit to the observable X-ray data, and in particular not for wind blown bubbles. As we shall show, the combination of the physical effects above and the properties (and limitations) of real X-ray observatories, can significantly affect the results of simple spectral fitting.  Previous authors (Weaver \\etal 1977;  Zhekov \\& Perinotto 1996) have calculated theoretical X-ray spectra from their 1-D models, but did not consider particular instruments or fit models to those spectra. In general only X-ray luminosities are calculated (\\eg Volk \\& Kwok 1985; Mellema \\& Frank 1995; Garcia-Segura \\& Mac Low 1995).   The aim of this paper is to introduce a method of analysing numerical simulations in the same way as actual X-ray observations are analysed, \\ie predict the {\\em observable} X-ray properties. This method can be applied to a wide range of phenomenon where X-rays are important, from wind-blown bubbles around WR stars and PN, through the larger bubbles around clusters of massive stars to starburst-driven galactic winds.  We simulate a wind blown bubble using a 2-dimensional hydrodynamic code, concentrating on the properties of the hot X-ray emitting gas. The hydrodynamic model is used to generate artificial X-ray spectra and images, in particular simulated {\\it ROSAT} spectra. We then analyse these spectra in the same way as real {\\it ROSAT} spectra would be, in order to determine what the observationally determined properties of the bubble would be, and how those relate to its actual structure.  This synthesis is necessary to a) determine the physical processes that are observationally important, and b) {\\em allow a direct comparison between observation and theory}.  Our model of a wind-blown bubble is deliberately chosen to be the simplest applicable model with an analytic solution, in order to simplify the analysis of our results, and avoiding added complications that a more physically accurate model of a wind blown bubble (\\eg Garc\\'{\\i}a-Segura, Mac Low \\& Langer 1996)  would introduce into the interpretation of our results. Later papers will consider more realistic models, with additional physics such as time varying energy and mass injection rates, along with spatial variation of the X-ray properties. This will be necessary to understand the properties of the extended emission from galactic winds (see for example Strickland, Stevens and Ponman 1997).  In Section~\\ref{sec:num_method} we describe the numerical code used to produce the results shown in Section~\\ref{sec:results}. Section~\\ref{sec:disc}  discusses the implications of these results, and we briefly sum up in Section~\\ref{sec:conclusions}.         \\begin{figure*} \\vspace{16.0cm} \\special{psfile=fig1_top.eps hoffset=0 voffset=-30 hscale=90 vscale=90 angle=0} \\special{psfile=fig1_bot.eps hoffset=0 voffset=-240 hscale=90 vscale=90 angle=0} \\caption{Logarithm of the gas number density during the simulation at $t = 3500$, $7700$, $10170$ and $14630 \\yr$. At $t = 3500 \\yr$ the bubble has suffered no significant radiative energy loss. Shell collapse is underway at $7700 \\yr$, approximately the time of maximum soft X-ray luminosity, and has just finished at $10170 \\yr$. The bubble then enters the self-similar phase, its properties at $t = 14630 \\yr$ being typical of this stage.} \\label{fig:dens_4t} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}   We have shown that in order to compare X-ray observations to theory, it is necessary to consider the {\\em observable} X-ray properties of the theory. The results of a spectral fit are a complex function of the the density and temperature distributions of the source, absorption, the properties of the detector used and the spectral fitting procedure. As such they should not be considered as ``real'' values, but as characteristic values, and specific to the instrument used. The normal method of fitting a simplistic model to the observed data, and then treating the best-fit parameters as the real properties can easily give answers an order of magnitude out from the truth.  This technique will allow the first direct comparison between observation and theoretical models of superbubbles and starburst driven outflows.\\\\  We would like to thank Trevor Ponman, Robin Williams and the referee for constructive criticism. DKS and IRS acknowledge financial support from PPARC. This work was performed on the Birmingham node of the {\\sc Starlink} network."
    },
    "9803/astro-ph9803229_arXiv.txt": {
        "abstract": "Recently gathered observational data on a sample of Type Ia Supernovae (SNe~Ia) reveal a wide distribution of expansion velocities of the Fe cores, measured from the width of the nebular lines. Moreover, the velocity appears to correlate with the luminosity decline rate after maximum light, $\\Delta m_{15}(B)$. Since it has been shown that for SNe~Ia $\\Delta m_{15}(B)$ correlates with the absolute magnitude at maximum, this then implies a relation between the expansion velocity of the Fe nebula and the luminosity at maximum.  Physically, the maximum luminosity is related to the mass of synthesized $^{56}$Ni, whereas the $FWHM$ of the lines is related to the kinetic energy of the ejecta. Our finding constitutes observational proof of the theoretical prediction that the two quantities have to be related. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/physics9803022_arXiv.txt": {
        "abstract": "A fully nonlinear, time-asymptotic theory of resonant particle trapping in large-amplitude quasi-parallel Alfv\\'en waves is presented. The effect of trapped particles on the nonlinear dynamics of quasi-stationary Alfv\\'enic discontinuities and coherent Alfv\\'en waves is highly non-trivial and forces to a significant departure of the theory from the conventional DNLS and KNLS equation models. The virial theorem is used to determine the time-asymptotic distribution function. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803159_arXiv.txt": {
        "abstract": "We present a high-resolution spectrum of the high redshift, lensed quasar Q1208+1101, obtained with the echellette spectrograph on the Multiple Mirror Telescope.  We examine the new and published spectra and provide an updated list of high-confidence metal-line absorption systems at $z=1.1349, 2.8626, 2.9118, 2.9136, 2.9149$.  Combining this with a simple model of the gravitational lens system allows us to constrain the possible lens redshifts.  The high-redshift ($z > 2.5$) and low-redshift ($z < 0.4$) candidates can be ruled out with high confidence. The current spectra effectively probe about 40\\% of the redshift range in which the lens is expected.  In that range, there is only one known metal-line absorption system, an MgII absorber at $z=1.1349$.  We consider the possibility that this system is the lensing galaxy and discuss the implied parameters of the galaxy. ",
        "introduction": "The bright, high redshift (z=3.815) radio quiet quasar Q1208+1011 has been identified as a gravitational lens by Maoz et al. (1992) and Magain et al. (1992).  The lens consists of two components (V=18.3 and 19.8 mag, Bahcall et al. 1992) separated by 0.\\arcsec 47, with a 4:1 intensity ratio.  The FOS HST spectra (Maoz et al., 1992) show that both components have the same redshift and similar spectra.  There are three key aspects in studying gravitational lenses: 1) understanding the lens geometry; 2) understanding the properties of the lensing galaxy; and 3) understanding the properties of the background source. Determining the amount of magnification allows us to understand the intrinsic quasar emission. Given a limiting observed magnitude, lensing allows us to probe both to lower intrinsic luminosities (at a certain redshift) or to higher redshifts (at a certain luminosity).  Q1208+1011 is apparently an extremely high luminosity source with an observed optical luminosity of $\\sim 10^{48}$~ergs~s$^{-1}$.  The true intrinsic luminosity is likely to be much lower, which affects the modeling and influences the parameters of the quasar models such as required black hole mass or accretion rates (Czerny 1994, Antonucci 1994, Siemiginowska et al 1996).  Precise lens modeling and evaluation of the quasar magnification requires detailed information about the lens, including its exact position relative to the quasar images, its morphology or mass distribution, and its redshift (Kochanek 1991). The SIS lens model predicts an average magnification of about 4 (Turner et al. 1984), however we cannot give a correct value for Q1208+1011 until the lens is detected.  Bechtold (1994) investigated the proximity effect in the spectra of Q1208+1101 and concluded that the data were consistent with a magnification of 1.  Fontana at al. (1997) give a factor of 20 magnification for Q1208+1101 based on high resolution Lyman alpha forest data.  Lens detection combined with the proximity effect could give stronger constraints on the magnification factor, and therefore allow more accurate modeling of this very luminous source.   Thus far the lensing galaxy for Q1208+1011 has not been directly imaged, consistent with the expectation that it is 4-6 magnitudes fainter than the quasar (Bahcall et al. 1992, Kochanek 1991).  The small separation indicates that a galaxy at relatively high redshift, $z \\ga 0.5$, is likely responsible for the lensing (Turner et al. 1984).  For this system and others with suspected high-redshift lensing galaxies, it may be possible and even necessary to identify the lens by its {\\em absorption} properties, rather than by its emission. With few exceptions, galaxies within $\\sim 30 h^{-1}$\\,kpc of the quasar line of site cause MgII or CIV absorption (Steidel 1997; Steidel 1993), so one would expect an metal-line absorption system at the lens galaxy redshift.  For Q1208+1011, the only published analysis of possible lens redshifts based on absorption lines has been by Magain et al. (1992).  They re-analyzed the absorption line data presented by Steidel (1990) and suggested at least 18 possible metal-line absorption systems, spanning redshifts from 0.3741 to 2.9157, with the majority in the range $2.5 < z < 3.1$.  They proposed that the most likely lens system was at redshift 2.9157 and derived a corresponding mass estimate for the lens. However, lens models indicate that such a high redshift location is highly unlikely (see Section~\\ref{sec:lens_model}).  Furthermore, most of low redshift identifications were based on just two doublet lines within the Ly-$\\alpha$ forest, whereas Bechtold and Yee (1995) have shown that the false detection rate for doublets in the forest is quite high (see also Section~\\ref{sec:abs_lines}).  To constrain the lens redshift more reliably in Q1208+1011, we obtained a high-resolution spectrum in March 1996 with the echellette spectrograph on the Multiple Mirror Telescope (MMT).  In this paper we examine the new and published spectra and provide an updated list of high-confidence metal-line absorption systems (Section~\\ref{sec:abs_lines} and Section~\\ref{sec:lens_model}).  We then combine this with a simple model of the gravitational lens system to constrain the possible lens redshifts (Section~\\ref{sec:lens_model}).  We show that the high-redshift ($z > 2.5$) and low-redshift ($z < 0.4$) candidates proposed by Magain et al. can be ruled out with high confidence, and that the current spectra effectively probe about 40\\% of the redshift range in which the lens is expected.  In that range, there is only one known metal-line absorption system, an MgII absorber at $z=1.1349$.  In Section~\\ref{sec:discussion} we consider the possibility that this system is the lensing galaxy and discuss the implied parameters of the galaxy. We also calculate the expected galaxy IR luminosity. ",
        "conclusions": "Because the separation between the two quasar images is only 0.\\arcsec 47, the mass of a normal galaxy is adequate to produce the lensed images. Assuming the Singular Isothermal Sphere (SIS) model for the lensing galaxy we can estimate the velocity dispersion and the enclosed mass for a given redshift.  In the redshift range which we have searched, there is only one candidate lens redshift, the $z=1.1349$ MgII system.  However, since there is a significant probability that this is not the lens redshift, we also calculate the parameters for several other interesting redshifts: $z=0.4$ low-redshift case (low end of the 90\\% probability range); $z=2.4$ high-redshift case (high end of the 90\\% probability range); and $z=2.9$ C\\,IV case, corresponding to the known C\\,IV absorption systems. In all the calculations below, we assume $\\rm \\Omega_0=0.1$ and $\\rm H_0=100h~km~sec^{-1} Mpc^{-1}$.    The mass of the galaxy can be obtained from:  $$ M  \\sim {4 \\theta^2  \\over 9} \\, {D_l D_s \\over D_{ls}}$$  \\noindent where $\\theta$ is the image separation, M is mass of the lens, $D_l, D_s, D_{ls}$ are the angular diameter distances to the lens, to the quasar and between the lens and the quasar respectively (see the review by Blandford \\& Narayan, 1992).  The corresponding velocity dispersion is related to the image separation by:  $$ \\theta = 4 \\pi {\\sigma^2 \\over c ^2} \\, {D_{ls} \\over D_s} = 2.6 \\arcsec \\sigma ^2 _{300} {D_{ls} \\over D_s} $$  \\noindent where $\\sigma = 300 \\times \\sigma _{300} $~km~s$^{-1} $ is a velocity dispersion.  Table 2 contains the calculated mass and velocity dispersions for the four considered lens redshifts.  The main uncertainty on the mass is related to cosmology and the uncertainty on the Hubble constant.    The required mass ($\\sim 2.8 \\times 10^{11} M_{\\odot}$) and the velocity dispersion ( $\\sim 202$~km~s$^{-1}$ for the lens at $z=1.1349$ are quite typical of normal galaxies.  We believe the Mg II absorption system at z=1.1349 is a strong candidate to be the lensing galaxy.  Absorption of the kind and strength we see is, with few exceptions, associated with a galaxy within $\\sim30 h^{-1}$\\,kpc of the line of sight (Le Brun, et al. 1995; Steidel, 1993). This implies that there is a galaxy within $\\sim 4$\\arcsec\\ of the Q1208+1011 pair.  Additionally, only in rare cases is there a galaxy within $\\sim 30h^{-1}$\\,kpc which {\\em does not} cause Mg\\,II absorption (Steidel 1993).  Thus far the lensing galaxy has not been detected -- pre-Costar HST imaging with the PC (Bahcall et al. 1992) limits the galaxy to have V$> 20.7$ if it is more than 0.\\arcsec 5 from the brighter component, or V$>19$ if the galaxy is between the two images.   Given the mass estimate for this system, we can predict its brightness using the Tully-Fisher and Faber-Jackson relations. If the galaxy is a disk system, the velocity dispersion implies that it is about 1 magnitude brighter than $L^*$; if it is an elliptical, it is 0.3 magnitudes fainter than $L^*$.  Figure 5 shows predicted lens galaxy magnitudes in HST $V$ (F555W) and $H$ (F160W) bands.  The luminosities were estimated by combining an SIS lens model with the Faber-Jackson or Tully-Fisher relation, and the magnitudes were then estimated by applying $K$ and evolutionary corrections computed from the spectral evolution models of Bruzual \\& Charlot (1993).  (See Keeton, Kochanek \\& Falco 1997 for details.) The predicted apparent magnitude in the visible is V=24.1-25.4 (see Figure~\\ref{fig:prob_dist}), much fainter than the limit $V\\sim20.7$ placed by Bahcall et al (1992) from the pre-Costar PC on HST.  The predicted near-IR magnitude is $H\\approx19.2-20.6$ (see Figure~\\ref{fig:prob_dist}), or $K\\approx20.2-21.6$.  This is near the faint end of the range of luminosities of galaxies selected by the presence of Mg\\,II absorption and described by Steidel \\& Dickinson (1995).  They presented data which showed that, for 5 Mg\\,II systems with $1.0 < z < 1.2$, the galaxy causing the absorption had K magnitude between 18.5 and 20.0.  We have simulated NICMOS observations to determine whether such a galaxy will be easily visible.  We assume the galaxy is centered between the quasar images, synthesize a test image, and remove the quasar images using a synthesized point-spread-function.  We find that a four-orbit exposure with the low background H-band (F160W) filter might give a detection with sufficient signal to estimate a lens model and the corresponding magnification.  A single-orbit exposure such as the one planned for Cycle~7 (Falco et al.)  \\footnote{Preliminary NICMOS images of Q1208+1011 have recently become available on the CASTLE Web page: {\\texttt http://cfa-www.harvard.edu/glensdata/1208.html}. The galaxy is not apparent in the image consistent with the predicted magnitude.} migth detect the core of the galaxy but will not likely trace the profile very far beyond the quasar image. We note that the image separation of 0.$\\arcsec$47 corresponds to $\\sim 3h^{-1}$~kpc, slightly smaller than typical scale lengths and effective radii of $L^*$ galaxies.  \\bigskip  We have combined our high-resolution spectra of the metal-line absorption systems towards the lensed quasar Q1208+1101 with gravitational lensing models.  We find the MgII absorber at z=1.1349 to be a plausible candidate for the lensing galaxy."
    },
    "9803/astro-ph9803190_arXiv.txt": {
        "abstract": "We report the detection of luminous extended X-ray emission in NGC 6240 on the basis of \\ros HRI observations of this ultraluminous IR galaxy.  The spatial structure and temporal behavior of the X-ray source were analyzed. We find that $\\ga 70\\%$ of the soft X-ray emission is extended beyond a radius of 5\\arcsec. Strong emission can be traced out to a radius of 20\\arcsec~ and weaker emission extends out to $\\sim$50\\arcsec. With a luminosity of at least $L_{\\rm x} \\simeq$ 10$^{42}$ erg/s this makes NGC 6240 one of the most luminous X-ray emitters in {\\em extended} emission known. Evidence for a nuclear compact variable component is indicated by a drop of 32\\% in the HRI count rate as compared to the PSPC data taken one year earlier. No short-timescale variability is detected. The HRI data, which represent the first high-resolution study of the X-ray emission from NGC 6240, complement previous spectral fits to \\ros PSPC data that suggested a two-component model consisting of thermal emission from shocked gas immersed in a starburst wind plus a powerlaw source attributed to scattered light from an obscured AGN.  We discuss several models to account for the extended and compact emission. Although pushed to its limits the starburst outflow model is tenable for the essential part of the {\\em extended} emission. For the AGN-type component we propose a model consisting of a near-nuclear `warm scatterer' that explains the apparent fading of the X-ray flux within a year as well as the strong FeK$\\alpha$ complex seen in an \\asca spectrum. ",
        "introduction": "With a far-infrared luminosity of $\\sim 10^{12} L_\\odot$ (Wright et al.\\ 1984) and a redshift of $z=0.024$, NGC 6240 is one of the nearest members of the class of ultraluminous infrared galaxies (hereafter ULIRG). The basic, as yet unsolved, enigma of these objects is the nature of the primary power source that has to be situated inside the central few arcseconds (Wynn-Williams \\& Becklin 1993). An amount of $\\sim10^{10} M_{\\sun}$ of cold molecular gas (e.g. Solomon et al. 1997), a record 2.121$\\mu$m-line luminosity from shocked `warm' H$_2$ (e.g. van der Werf et al. 1993), earthbound IR spectra (e.g. Joseph \\& Wright 1985, Rieke et al. 1985, Schmitt et al. 1996), MAMA (Smith et al. 1992) and HST (Barbieri et al. 1993) observations, and recent ISO-SWS spectra (Lutz et al. 1996) all point towards the presence of hidden prodigious star formation after onset of a galactic collision, which could be responsible for the FIR power. Heckman et al. (1987, 1990) found indications for superwind and supershell activity, a well-known result of strong starbursts.  However, the smallness of the recombination line flux (de Poy et al. 1986), the detection of a high-excitation component in HST images (Barbieri et al. 1995, Rafanelli et al. 1997) and general arguments valid for ULIRGs as a class (e.g. Sanders et al. 1988) suggest that a dust-shrouded AGN contributes significantly to the heating of the dust that emits the FIR radiation.  The unambiguous detection and investigation of an AGN in NGC 6240 and other interacting ULIRGs would be of prime importance for our understanding of the formation and evolution of AGN in general. It has been proposed that starbursts are the germ cell for the formation of AGN (Weedman 1983, Barnes \\& Hernquist 1991, Mihos \\& Hernquist 1996) and interaction may provide the triggers and fuel for both kinds of activity (see Sanders \\& Mirabel 1996 for a recent review). A large fraction of ULIRGs indeed turned out to be interacting systems (e.g. Andreasian \\& Alloin 1994, Clements et al. 1996).  As outlined above, there are indications for a starburst in NGC 6240, but what would be the best evidence for a hidden AGN? The far-infrared emission is reprocessed black-body like radiation containing no direct clue on the nature of the primary source. Near-IR and mid-IR line spectra provided signatures for a red giant population and a younger burst. A few high-excitation features in IR spectra and in optical HST narrow-band images could be due to an AGN but not necessarily. The optical emission-line spectrum (Fosbury \\& Wall 1979, Zasov \\& Karachentsev 1979, Fried \\& Schulz 1983, Morris \\& Ward 1988, Keel 1990, Heckman et al. 1987, Veilleux et al. 1995, Schmitt et al. 1996) is dominated by LINER-like line ratios over the central $\\sim10$ kpc. Its large extent and little variation in excitation tracers is more easily attributed to shock-heating rather than to a central photoionizing AGN continuum.  X-rays are an important tool for studying both, an AGN as well as starburst components. In the \\ros band, AGN tend to be dominated by strong powerlaw (hereafter PL) emission while starbursts can usually be represented by thermal spectra. In a recent analysis of \\ros PSPC spectra from NGC 6240, Schulz et al. (1998; hereafter paper I) found good fits by either a single thermal Raymond-Smith (hereafter RS) spectrum with $L_{\\rm 0.1-2.4 keV} = 3.8\\,10^{43}$ erg/s (for a distance of 144 Mpc) or a hybrid model consisting of 80\\% PL plus 20\\% thermal RS (dubbed 0.8PL+0.2RS below) contributions and a total luminosity of $5.2\\,10^{42}$ erg/s. Since the spectral shape with PSPC resolution is not sufficiently distinctive the luminosity information was taken as an additional constraint. Due to the unprecedented high luminosity of the single RS model and additional severe difficulties to explain it in terms of scalable superbubble models, the hybrid model was favored. This is also supported by the {\\asca} detection of a strong FeK$\\alpha$ line in the X-ray spectrum of NGC 6240 (Mitsuda 1995; the same data indicate further emission lines around 1--2 keV). The powerlaw was attributed to the electron scattered X-ray flux from a hidden AGN so that an AGN-plus-starburst scenario was proposed for the ultimate power source of NGC 6240 (paper I).  The deep HRI observations which are discussed below represent the first high spatial resolution study of the X-ray emission from NGC 6240. They allow to trace the emission from a thermal starburst source that is expected to be appreciably spatially extended while an AGN-induced powerlaw source should be much more compact unless there is extensive large-scale scattering. Further, they provide information on the long- and short-term X-ray variability of the source.  Luminosities given below are calculated assuming a distance $d = 144$ Mpc of NGC 6240. This yields a scale perpendicular to the line of sight in which 1\\arcsec~ corresponds to 700 pc in the galaxy. ",
        "conclusions": "We detected luminous extended X-ray emission in NGC 6240 in \\ros HRI data. At the given spatial resolution the source looks nearly spherical and contains its most significant emission within a radius of 20\\arcsec~(or 14 kpc for a distance d=144 Mpc) where the total 0.1--2.4 keV X-ray luminosity amounts to at least $\\sim 10^{42}$ erg s$^{-1}$. At the epochs of the observations at most 40\\% of this luminosity arises within the innermost region of 5\\arcsec~radius.  The extended emission can be consistently described by crude supershell models thereby explaining it as the result of a super-starburst with a total luminosity close to $10^{12} L_{\\sun}$.  The presence of an additional compact AGN component is in X-rays indicated by (i) a decrease in the count rate between Feb. 1993 (last PSPC observation) and Feb. 1994 (first HRI observation) by 32\\%, (ii) a probable powerlaw component necessary to fit PSPC spectra and (iii) a strong FeK$\\alpha$ complex detected in \\asca spectra. We propose a model in which near-nuclear warm gas ionized by the AGN powerlaw continuum emits FeK$\\alpha$ which is seen superposed on the reflected continuum.  Both components, the starburst as well as the AGN provide enough power to explain the luminous FIR emission in NGC 6240 and it seems that both contribute with comparable strength."
    },
    "9803/astro-ph9803049_arXiv.txt": {
        "abstract": "s{Several conclusions have been reached over the last few years concerning high-redshift galaxies: (1)~The excess of faint blue galaxies is due to dwarf galaxies.  (2)~Star formation peaks at redshifts $z\\approx 1-2$.  (3)~It appears to occur piecemeal in any given galaxy and there is no evidence for starbursting throughout a large $\\sim 10\\kpc$ galaxy.  (4)~There is significant and sharp diminution in the number of $L_\\star$ spiral galaxies at redshifts $1<z<2$ and elliptical galaxies at redshifts $2.5<z<4$. (5)~It is increasingly more difficult to ``hide'' large high-redshift galaxies in universes with larger volumes per unit redshift, i.e., open or $\\lambda$ models, which have lower deceleration parameters.} ",
        "introduction": "This paper reviews high redshift galaxies as we understand them.  It is not meant to be comprehensive.  Instead, we ask the questions that seem most important to us.  Why do we observe high redshift galaxies, \\S\\ref{sec:why}? How do we observe them, \\S\\ref{sec:how}?  What have we learned from the observations, \\S\\ref{sec:learned}?  And how should we continue our research, \\S\\ref{sec:where}? ",
        "conclusions": ""
    },
    "9803/astro-ph9803080_arXiv.txt": {
        "abstract": "\\indent Using a limited, but representative sample of sources in the ISM of our Galaxy with published spectra from the {\\it Infrared Space Observatory}, we analyze flux ratios between the major mid-IR emission features (EFs) centered around 6.2, 7.7, 8.6 and 11.3\\um, respectively.  In a flux ratio-to-flux ratio plot of EF(6.2\\um)/EF(7.7\\um) as a function of EF(11.3\\um)/EF(7.7\\um), the sample sources form roughly a $\\Lambda$-shaped locus which appear to trace, on an overall basis, the hardness of a local heating radiation field.  But some driving parameters other than the radiation field may also be required for a full interpretation of this trend.  On the other hand, the flux ratio of EF(8.6\\um)/EF(7.7\\um) shows little variation over the sample sources, except for two HII regions which have much higher values for this ratio due to an ``EF(8.6\\um) anomaly,'' a phenomenon clearly associated with environments of an intense far-UV radiation field. If further confirmed on a larger database, these trends should provide crucial information on how the EF carriers collectively respond to a changing environment. ",
        "introduction": "\\label{sec1}  Since they were first discovered about two decades ago (Gillett \\etal 1973), the mid-IR emission features (EFs), \\ie those broad emission bands centered respectively at 6.2, 7.7, 8.6 and 11.3\\um,  have been detected from a variety of sources in our own Galaxy as well as in some other galaxies.   This makes it rather clear that the carriers of these EFs play a significant role in regulating the physical conditions in the ISM.  One concept that has gained popularity is that these EFs arise from the vibrational modes of the so-called aromatic hydrocarbon molecules (hereafter PAHs;  L\\'eger \\& Puget 1984; Allamandola \\etal 1985).  However, this is largely based on a wavelength coincidence between the observed EFs and the absorption bands measured in laboratories, a criterion that other candidate EF carriers (\\eg Papoular \\etal 1989; Sakata \\etal 1984) also seem to satisfy.   It is therefore important to test this PAH and other candidate scenarios in as many ways as possible.  One possible way is to observe how the strength ratios between the EFs react to a changing heating radiation field in the ISM.  This approach is viable because, for example, if EFs arise from PAHs, some of their feature-to-feature strength ratios should depend on the hardness of the heating radiation field via factors such as the fraction of PAH cations (\\ie singly ionized PAHs;  Langhoff 1996) and the degree of dehydrogenation (\\eg Jourdain de Muizon \\etal 1990).  Because of the limited sensitivity and spectral coverage associated with sub-orbital platform observations, studies on how features respond to a changing environment have been carried out so far only to a limited extent (\\eg Cohen \\etal 1986; Joblin \\etal 1996).   With its unprecedented sensitivity and continuous spectral coverage, the Infrared Space Observatory (ISO; Kessler \\etal 1996) has been used to obtain mid-IR spectra of sources in a variety of environments ranging over far-UV intense to typical diffuse ISM.   This makes it possible for the first time to study the relative strengths among all the major EFs under a wide range of physical conditions.  We gather here the published ISO mid-IR spectra on a limited, but representative set of sources in the ISM of our Galaxy to seek possible correlations between the feature-to-feature flux ratios and the local heating radiation field.  While the former parameters are measured directly from the ISO spectra, the latter is inferred from the properties of the stars that likely dominate the radiation field at the locations where the ISO spectra were taken.   In Sect.~2 we describe the sample, the data and how we evaluate the integrated flux of a feature.  In Sect.~3 we analyze feature-to-feature flux ratios and identify some possible systematic trends.  Some implications from these trends are discussed in Sect.~4, when applicable, in comparison with the current knowledge of PAHs. ",
        "conclusions": "\\label{sec4}   The result of Fig.~1a that the ratios of EF(8.6)/EF(7.7) for EF(8.6)-normal sources stay nearly constant is already quite secure at this point.  This suggests either (a) that there is roughly a fixed ratio between the strengths of these two EFs or (b) that EF(7.7) has a long wavelength tail/plateau that dominates the integrated flux of EF(8.6) as a result of our dividing the two EFs at 8.1\\um\\ in evaluating their fluxes.  Under the current framework of PAHs, however, the ratio for EF(8.6)/EF(7.7) is expected to decrease with an increasing degree of dehydrogenation associated with an increasing hardness of the heating radiation field (Jourdain de Muizon \\etal 1990). So the laboratory results on PAHs favor (b).   On the other hand, additional and improved data are probably needed for further confirmation on the details of the distribution pattern in Fig.~1b.  There is a slight possibility that, once more data points are inserted in Fig.~1b, a simpler pattern [\\eg EF(6.2)/EF(7.7) increases somewhat as EF(11.3)/EF(7.7) decreases] could emerge. However, if this $\\Lambda$-shaped distribution pattern is further confirmed, it could provide new constraints on the identification of the EF carriers. For example, the right side of the $\\Lambda$-shaped curve in Fig.~1b may be consistent with a picture where one has a combination of rising PAH temperature and increasing ionization effect (Langhoff \\etal 1996) as one goes up along the curve. However, the current knowledge of PAHs does not seem to suggest a turnaround in EF(6.2)/EF(7.7).  An increasing photodestruction effect does suggest a decreasing ratio of EF(6.2)/EF(7.7) as smaller PAHs are easier to be destroyed than their larger cousins (Allamandola \\etal 1989; Allain, Leach, \\& Sedlmayr 1996a, 1996b) and as larger PAHs reach lower peak temperatures after absorbing a UV photon (L\\'eger \\& Puget 1989). But the wavelength difference between the two features is so small that the dependence of their ratio on the PAH size distribution seems inadequate to explain the observed large change in EF(6.2)/EF(7.7).  Besides, this effect has to be counter-balanced by the fact that in an increasingly harder UV-rich environment, the average PAH temperature also gets hotter.  Perhaps, this implies that something in addition to the radiation field is needed to explain the distribution pattern in Fig.~1b.   The laboratory counterpart of the EF(8.6)-anomaly phenomenon is unknown at this point.  One speculation is that both EF(7.7) and EF(8.6) are stronger relative to the other features in a far-UV rich environment.  Perhaps, the two features sit on top of an emission plateau whose contribution to the fluxes of EF(7.7) and EF(8.6) becomes relatively important only under an intense far-UV radiation field.  In fact, such an emission plateau may put the EF(8.6)-anomaly phenomenon in a sequence between those EF(8.6)-normal sources and those with even more ``unusual'' EF spectra, \\eg an ISO spectrum dominated by a previously unknown broad feature around 8\\um\\ seen in an ultra compact HII region (Cesarsky \\etal 1996c) to those featureless spectra seen in many active galactic nuclei (\\eg Aitken \\& Roche 1985).   Regardless of how consistent they are with the physical properties of a specific class of candidates for the EF carriers, the trends in Fig.~1 should inevitably lead us to a more complete picture on how the relative EF strengths change from one type of environment to another. A direct application of this would be to help interpreting the global EF spectra from galaxies.  For example, a preliminary analysis shows that in a plot such as Fig.~1b, most normal galaxies scatter within a region close to that occupied by those data points taken along the inner Galactic plane (Lu \\etal 1998), suggesting that the global EFs of a galaxy may be dominated by the emission from the diffuse ISM component."
    },
    "9803/astro-ph9803339_arXiv.txt": {
        "abstract": "We investigate the dynamics and radiation from a relativistic blast-wave which decelerates as it sweeps up ambient matter.  The bulk kinetic energy of the blast-wave shell is converted into internal energy by the process of accreting external matter.  If it takes the form of non-thermal electrons and magnetic fields, then this internal energy will be emitted as synchrotron and synchrotron self-Compton radiation.  We perform analytic and numerical calculations for the deceleration and radiative processes and present time-resolved spectra throughout the evolution of the blast-wave.  We also examine the dependence of the burst spectra and light curves on various parameters describing the magnetic field and non-thermal electron distributions. We find that for bursts such as GRB~910503, GRB~910601 and GRB~910814, the spectral shapes of the prompt gamma-ray emission at the peaks in $\\nu F_\\nu$ strongly constrain the magnetic fields in these bursts to be well below ($\\la 10^{-2}$) the equipartition values.  These calculations are also considered in the context of the afterglow emission from the recently detected gamma-ray burst counterparts. ",
        "introduction": "The recent Beppo-SAX observations of fading X-ray afterglows coincident with the positions of gamma-ray bursts GRB~970228 (Costa et al.\\ 1997\\markcite{Costa97}) and GRB~970508 (Piro et al.\\ 1998\\markcite{Piro98}) have led to the first identifications of possible burst counterparts in the optical wave band (van Paradijs et al.\\ 1997\\markcite{vanParadijs97}; Djorgovski et al.\\ 1997\\markcite{Djorgovski97}).  In the optical spectrum associated with the latter burst, absorption lines have been detected with a redshift of $z = 0.835$ providing, for the first time, direct evidence for a gamma-ray burst at cosmological distances (Metzger et al.\\ 1997\\markcite{Metzger97}). The gamma-ray fluence of GRB~970508 measured by BATSE is $\\sim 3 \\times 10^{-6}\\,$erg~cm$^{-2}$ (Kouveliotou et al.\\ 1997\\markcite{Kouveliotou97}).  Therefore, if the absorption line measurements give a lower limit on the redshift of the burst, this implies an isotropic burst energy of $> 10^{51}\\,$erg~s$^{-1}$.  An impulsive event releasing this amount of energy in a compact region naturally leads to a fireball and thence to a relativistic blast-wave.  Blast-wave models for gamma-ray bursts have been examined previously in the literature in several contexts.  The most extensive body of work on this topic has been produced by \\Meszaros, Rees and collaborators (e.g., \\Meszaros\\ \\& Rees 1992a; Rees \\& \\Meszaros\\ 1992; \\Meszaros, Laguna \\& Rees 1993; Wijers, Rees \\& \\Meszaros\\ 1997; Panaitescu \\& \\Meszaros\\ 1998a, 1998b; et al.).  Other recent papers include Sari \\& Piran (1995), Vietri (1997), Waxman (1997), and Katz \\& Piran (1997).  The basic fireball/blast-wave model consists of some triggering event---either coalescing neutron stars or black holes (\\Meszaros\\ \\& Rees 1992b) or the collapse of a massive star (Paczy\\'nski 1998) or a failed type II supernova (Woosley 1993)---depositing a large amount of energy, $E_0 \\sim 10^{51}$--$10^{55}$~ergs, in a small region with radius $r_0 \\sim 10^6$--$10^7$~cm.  Because of the unavoidable presence of baryonic material, it is expected that the initial fireball energy will be transformed into kinetic energy of these baryons rather than escape as radiation.  This material expands until the internal motions of the baryons become sub-relativistic in the co-moving frame of the material, at which point it forms a cold shell with bulk Lorentz factor $\\G_0 \\simeq E_0/M_0 c^2$, where $M_0$ is the rest mass of the contaminating baryons.  This shell continues to expand freely into the surrounding ambient medium until the integrated momentum impulse upon the shell by the swept-up matter is equal to the rest mass of the original material, $\\G_0 4\\pi r_d^3 \\rho_\\ext /3 \\approx M_0$ where $\\rho_\\ext$ is the mass density of the external medium.  This defines the so-called deceleration radius $r_d$ (Rees \\& \\Meszaros\\ 1992).  Beyond this radius, the shell can no longer be regarded as freely expanding, and the bulk kinetic energy of the blast-wave begins to be reconverted into internal energy.  If this internal energy is radiated promptly, then the deceleration of the blast-wave shell can be approximately described by $\\G(r) \\propto r^{-3}$, and the expansion is said to be in the radiative regime.  On the other hand, if the internal energy is radiated on a time scale which is long compared to the expansion time scale, then the expansion is in the non-radiative regime, and $\\G(r) \\propto r^{-3/2}$.  In either case, for large initial Lorentz factors, $\\G_0 \\sim 10^2$--$10^3$, relativistic effects (e.g., Rees 1966) compress the time scale for the radiation such that the bulk of the blast-wave energy is emitted in the first tens of seconds in the observer's frame following the initial detonation event, thus producing the observed gamma-ray burst.  Several authors have pointed out that well after the prompt gamma-ray burst event the blast-wave shell will continue to decelerate and radiate (Vietri 1997; Waxman 1997).  The recent detections by the X-ray, optical and radio communities of the aforementioned fading X-ray, optical and radio counterparts within the error boxes of GRBs appear to support this picture.  Furthermore, model estimates of the temporal decay of these transients yield time-dependencies which agree with those observed, $F_\\nu \\sim t^{-1}$ (Wijers et al.\\ 1997); and for the one of the bursts for which optical data are available (GRB~970508), the optical spectral index is consistent with synchrotron emission from a power-law distribution of electrons, $dN/d\\g \\propto \\g^{-p}$ with $p \\simeq 2$--2.3 (Djorgovski et al.\\ 1997) indicating, for example, a shock-accelerated electron population.  Despite the successes of the blast-wave model in accounting for the prompt burst properties and its prediction of fading afterglows, several important theoretical questions must be addressed in order to have a reasonably complete model: \\begin{itemize} \\item What is the nature of the coupling between the electrons, protons and magnetic field (Panai\\-tescu \\& \\Meszaros\\ 1998b)?  To what extent can these components be in equipartition given that only the electrons can efficiently radiate away their energy? \\item How does the magnetic field change as the blast-wave decelerates? If it is initially formed through equipartition processes, but is not strongly coupled to the non-thermal electron energy density, how does it evolve? \\item What is the nature of the acceleration mechanism?  Are the electrons energized by repeated diffusion across the shocks themselves or by gyroresonant scattering with disturbances in the post-shock turbulent MHD fluid? \\item What is the proper form for the injected electron energy distributions?  Is it well described by a typical power-law spectrum?  If so, what determines the characteristic energies of the particles? \\end{itemize}  In order to begin to address these questions, it is worthwhile to go beyond the simple, though useful, back-of-the-envelope estimates which have generally prevailed in the literature thus far.  In this paper, we present a detailed calculation of the blast-wave deceleration and the evolution of the magnetic fields and electron distributions under various assumptions relevant to the above issues.  We compute model light curves and spectra and use the available afterglow data to help discriminate between the various options.  The format of the paper is as follows: In \\S~2, we describe the basic dynamics of an impulsively driven blast-wave which decelerates by accretion of ambient material. In \\S~3, we discuss the physical processes responsible for producing the observed radiation including prescriptions for magnetic field generation, the formation of non-thermal particle distributions and the relevant radiation processes.  The numerical procedures for computing the deceleration of the blast-wave and the integration of the blast-wave shell emission are described in \\S~4.  Model spectra and light curves for gamma-ray bursts are presented and analyzed in \\S~5.  Lastly, in \\S~6, we discuss these results, explore prospects for further research and present our conclusions. ",
        "conclusions": "In this paper, we have attempted a more realistic calculation of the dynamics and synchrotron and synchrotron self-Compton emission for the blast-wave model of gamma-ray bursts.  By matching the detailed characteristics of burst spectra, we have found relations (eqs.~\\ref{Gamma_limit} \\&~\\ref{xi_B_limit}) which place constraints on magnetic field strengths and bulk Lorentz factors.  If these relations are to be believed, then burst data can have a significant impact on models of magnetic field generation in turbulent plasmas.  The apparent deficiencies of this calculation point towards areas of further research.  In particular, the detailed temporal structure of individual bursts is not explicitly dealt with in this model.  In the context of external shocks, it may be due to inhomogeneities in the external medium, or fluctuations in the electron injection and/or magnetic field equipartition parameters (Panaitescu \\& \\Meszaros\\ 1998a).  For $s > 3$, the generalized expression for the luminosity for energies $\\e \\ge \\e_\\peak$ (eq.~\\ref{peak_luminosity}) is \\begin{equation} \\e^2 \\frac{dN}{d\\e dt} = 2\\pi m_p c^2 \\xi_e r^2 n_\\ext(r) \\G^2(r) (\\e/\\e_\\peak)^\\lambda \\end{equation} where $\\lambda = (2-s)/2$ applies for relatively strong magnetic fields $\\xi_B \\ga 10^{-4}$ when the electrons just above the break are efficiently cooled and $\\lambda = (3-s)/2$ applies for relatively weak fields and uncooled electrons.  From this expression we see that any burst light curve substructure must be due to variations in $\\xi_e$ and $n_\\ext$, and indirectly, due to variations in $\\xi_B$ through $\\e_\\peak$ (eq.~\\ref{peak_energy}).  Our treatment of the dynamics also ignores the structure of the shock region itself, and our approach essentially only considers the emission from the forward shock and neglects the reverse shock. Panaitescu \\& \\Meszaros\\ (1998a) have performed calculations similar to our own, but from a hydrodynamical perspective, and found that the reverse shock only makes a significant contribution to the emission at optical and UV energies.  Therefore, neglecting the reverse shock should not affect our results for the gamma-ray emission, but it could have a significant impact on the optical and radio afterglow emission. We also neglect the thickness, $\\Delta r$, of the shock shell in integrating the emission for a given observer time $\\dt$.  This should not be important at early times when $\\Delta r = r_0/\\G_0^2$ (in the lab frame), but it could affect the afterglow emission at late times.  It is unlikely that the blast-wave itself is spherical.  If the initial fireball is created by the coalescence of two compact objects, then the orbital plane defines a natural axis of symmetry along which the blast-wave will propagate (\\Meszaros\\ \\& Rees 1992b).  This sort of asymmetry could be accounted for in our model by a non-unity collimation factor, $f_b$.  Furthermore, if the observer line-of-sight does not lie within the opening angle of the blast-wave cone, then other effects due to relativistic beaming and the gradual deceleration of the shock front would be introduced.  In this respect, highly anisotropic blast-waves would share properties with relativistic jets in blazars.  This analogy can be take even further by noting the similarity of the burst spectra we derive compared to that of gamma-ray blazars.  Like our model spectra, the spectral energy distributions (SEDs) of these objects tend to have two peaks, one in the UV--X-ray range and one at gamma-ray energies.  If the lower peak in blazar SEDs is due to synchrotron emission and corresponds to the $\\sim 1$~MeV peak in gamma-ray bursts, we can apply a similar analysis as we have discussed above to derive bulk Lorentz factors and equipartition parameters for blazars.  In particular, the recent ASCA observations of Mrk~421 (Takahashi et al.\\ 1996) provide sufficient information to get actual values rather than simply upper or lower limits.  Using the light curves of Mrk~421 measured in different X-ray energy bands, Takahashi et al.\\ (1996) performed a cross-correlation analysis and found that the longer relative time lags of the lower energy data versus the higher energy data are consistent with synchrotron cooling of the underlying electron distribution.  Several authors have noted this effect and have calculated this temporal dependence for the cases of bursts and blazars (e.g., Tashiro et al.\\ 1995; Tavani 1996; Dermer 1998).  Takahashi et al.\\ use the TeV variability time scale (Kerrick et al.\\ 1995) to estimate a Doppler factor and find $\\D = 5$ (cf.\\ Takahara 1994).  Using this estimate and their time lag measurements, they derive a magnetic field of $B = 0.2$~G. From non-simultaneous data (Shrader \\& Wehrle 1997), the synchrotron portion of the SED of Mrk~421 peaks at about $\\sim 130~$eV. Using \\begin{equation} \\e_\\peak = \\frac{B}{B_{\\rm crit}} \\g^2 {\\cal D} \\end{equation} (cf.\\ eq.~\\ref{peak_energy}), and $\\g = (m_p/m_e)\\G$, we find $\\G \\approx 60$ and an observer angle $\\theta \\approx 5^\\circ$. We also obtain an equipartition field strength of $B_{eq} \\approx 10 n_1^{1/2}$~G implying an equipartition parameter of $\\xi_B \\sim 10^{-2}$.  Although the above value for the bulk Lorentz factor is substantially larger than the mean value of $\\langle \\G \\rangle \\sim 10$ found by applying the beaming model to a sample of radio-loud objects (Urry \\& Padovani 1995), its large value may indicate the special nature of gamma-ray loud blazars which are characterized not only by small observing angles but also by larger than typical bulk Lorentz factors.  Despite the crudeness of this calculation, it illustrates the potential applicability of this sort of analysis to blazars as well as bursts."
    },
    "9803/astro-ph9803175_arXiv.txt": {
        "abstract": "We report   the detection of  four  images in the  recently discovered lensed QSO RX~J0911.4+0551.   With   a maximum angular   separation of $3.1$\\arcsec, it  is the  quadruply imaged  QSO  with the widest known angular separation.  Raw and deconvolved data reveal an elongated lens galaxy.  The observed reddening in at least two of the four QSO images suggests differential extinction by this lensing galaxy.  We show that both an ellipticity of  the galaxy ($\\epsilon_{\\rm min}=0.075$) and an external shear ($\\gamma_{\\rm min}=0.15$) from  a nearby mass has to be included  in the lensing  potential in order  to reproduce the complex geometry observed in  RX~J0911.4+0551.  A  possible galaxy cluster  is detected about  38\\arcsec\\, from  RX~J0911.4+0551 and could contribute to the X-ray emission  observed by ROSAT in  this field.  The color of these galaxies indicates a plausible redshift in the range of 0.6-0.8. ",
        "introduction": "RX~J0911.4+0551, an AGN candidate  selected   from the ROSAT   All-Sky Survey (RASS)  (Bade et al.  1995, Hagen  et  al.  1995), has recently been classified  by Bade   et al.  (1997;    hereafter B97) as   a new multiply imaged  QSO.   B97 show that  it consists  of  at least three objects: two barely  resolved    components and a third   fainter  one located 3.1\\arcsec\\ away from the other two.  They  also show that the spectrum of this  third fainter component is  similar to the  combined spectrum  of the two bright components.   The lensed source is a radio quiet  QSO at $z=2.8$.   Since RASS detections  of distant radio quiet QSOs are rare,   B97 pointed out   that the observed X-ray  flux might originate from a galaxy cluster at $z \\geq 0.5$ within the ROSAT error box.  We present here  new optical and near-IR  high-resolution images of RX~J0911.4+0551  obtained with the   2.56m Nordic Optical Telescope (NOT) and   the  ESO  3.5m New  Technology   Telescope  (NTT). Careful deconvolution of the data allows us to clearly resolve the object into four QSO components  and a lensing galaxy.   In  addition, a candidate galaxy cluster is detected in the vicinity of the four QSO images.  We estimate   its redshift from the   photometric  analysis of its member galaxies. ",
        "conclusions": "Thanks    to   our   new high-resolution    imaging    data,   the QSO RX~J0911.4+0551   is  resolved  into   four   images.  In    addition, deconvolution with the  new MCS algorithm  reveals  the lensing galaxy, clearly  confirming the   lensed nature  of  this system.    The image deconvolution provides precise photometry and  astrometry for all  the components of the system.  Reddening in components A2 and A3 relative to A1  is observed from our $U$, $V$,  and $I$ frames  that were taken  within  three hours on the same    night.  The  absence  of reddening    in  component  B and the difference  in    reddening  between components    A2 and   A3 suggest extinction by  the deflecting galaxy.  Note that although  our near-IR data were obtained from  15 days to  6 weeks after the optical images, they appear to be consistent with  the optical fluxes measured for the QSO images, i.e.   flux ratios increase continuously  with wavelength, from $U$ to $K$, indicating extinction by the lensing galaxy.  We have discovered a good galaxy cluster  candidate in the SW vicinity of RX~J0911.4+0551 from our field photometry  in the $I$, $J$, and $K$ bands.  Comparison of our color-magnitude diagram with that of a blank field (e.g., Moustakas et  al.  1997) shows  that the galaxies  around RX~J0911.4+0551  are  redder  than  field-galaxies   at an  equivalent apparent   magnitude.    In   addition,  the  brightest    galaxies in Fig.~\\ref{fig:cmd} lie on a red sequence at $I-K\\sim 3.3$, typical for the early  type members of  a distant galaxy  cluster.  The two dashed lines indicate   our $\\pm0.4$ color error   bars at $K\\sim  19$ around $I-K\\sim3.3$. Most of these galaxies  are grouped in the region around a double elliptical at a distance of  $\\sim38$\\arcsec\\, and a position angle of $\\sim204^{\\circ}$  relative to A1.   This can also be seen in Fig.~\\ref{fig:field} which shows a group  of red galaxies with similar colors centered    on the double  elliptical  (in  the  center  of the circle).  Consequently, there  is considerable evidence  for at least one galaxy cluster in the field.  The redshift of our best candidate cluster (the one circled in Fig.~\\ref{fig:field}) can be estimated from the $I$ and $K$ band photometry.  We have  compared the $K$-band magnitudes of the brightest   cluster galaxies with  the    empirical $K$ magnitude  vs. redshift  relation found by  Arag{\\'o}n-Salamanca  et al.  (1998).  We find that our cluster candidate,  with its brightest $K$ magnitude  of about $\\sim17.0$, should have  a  redshift of $z\\sim0.7$.   A  similar comparison  has been done in the  $I$-band without taking into account galaxy  morphology.  We compare the  mean $I$ magnitude of the cluster members with the  ones found by Koo et  al.  (1996)  for galaxies with known redshifts in the Hubble Deep Field and obtain a cluster redshift between  0.6 and 0.9.  Finally, comparison  of the  $I-K$ color of the galaxy sequence   with data  and models from   Kodama  et  al.  (1998) confirm the redshift estimate of 0.6-0.8.  In  order to calculate physical  quantities  from the model parameters found in section 4, we assume a simple model for the cluster which may be responsible  for the external shear.   For an isothermal potential, the true shear and convergence are of the same order of magnitude.  As the  convergence is not explicitly  included in the model, the deduced shear is a reduced shear leading to an absolute convergence of $\\kappa =   \\gamma/(1+\\gamma)  = 0.241$.  For  a   cluster redshift of $z_{\\rm d}=0.7$ and  with cosmological parameters $\\Omega=1$, $\\lambda=0$ this corresponds to   a velocity dispersion  of about   $1100\\,\\kms$ if the cluster is  positioned at  an angular  distance  of 40\\arcsec\\,.   See Gorenstein, Falco \\& Shapiro (1988) for a discussion of the degeneracy preventing a direct determination of $\\kappa$.   From the direction of the shear $\\phi$,  (see  Table ~\\ref{tab:bestmod}) we can  predict the position angle of the cluster as seen from the QSO to be $12^\\circ$ or $192^\\circ$.  The  latter value agrees well  with  the position of our cluster candidate SW of the QSO images.   Note also the good agreement between  the position angle $\\thg$    derived from the observed  light distribution,  and the predicted  position angle  corresponding to our best fitting model of the lensing potential. Interestingly, this is in good agreement  with Keeton,  Kochanek \\& Falco  (1998) who  find that projected mass distributions  are generally aligned with the projected light distributions to less than $10^{\\circ}$.  The color of the  main lensing galaxy is very  similar to that of  the cluster members, suggesting that it might be a  member of the cluster. Using the same model for the cluster  as above, assuming the galaxy at the same redshift as the cluster, and neglecting the small ellipticity of $\\epsilon<0.05$, the velocity dispersion  of the lensing galaxy can be predicted from the calculated deflection angle  $\\alpha_0$ to be of the  order of $240\\,\\kms$.  Since the  galaxy profile is sharp towards the  nucleus in $K$, we   cannot rule out  the  possibility of a fifth central image of  the   source,  as predicted for  non-singular   lens models.   Near-IR    spectroscopy  is   needed  to    get  a  redshift determination of the  lens and to  show  whether it is blended  or not with a fifth image of the (QSO) source.  Some 10\\arcsec\\,  SW from the  lens, we detect a   small group of even redder   objects.    These     red    galaxies  can   be      seen  in Fig.~\\ref{fig:field} a few arcseconds to the left  and to the right of the cross. They  might be part  of a  second galaxy-group at  a higher redshift, and   with  a position in  better   agreement with the X-ray position  mentioned by B97.  However,  since the measured X-ray signal is near the detection limit, and the 1-$\\sigma$ positional uncertainty is at least 20\\arcsec\\,, the  X-ray emission  is compatible with  both the QSO  and  these galaxy  groups  in  the field.  Furthermore,  this second group, at $z>0.7$, would most likely  be too faint in the X-ray domain to be detected in  the RASS.  In fact,  even our lower redshift cluster candidate would need to have  an X-ray luminosity of the order of   $\\rm  L_{0.1-2.4 \\rm   keV}\\sim  7.10^{44}  \\rm erg\\,\\rm  s^{-1}$ (assuming    a   6    keV    thermal     spectrum,  $\\rm    H_{0}=50\\, \\rm{km\\,s}^{-1}\\,\\rm Mpc^{-1}$,   $\\rm q_{0}=0.5$),  in  order to   be detected  with  0.02 $\\rm cts\\,\\rm   s^{-1}$ by  ROSAT.  This is  very bright but not  unrealistic for high  redshift galaxy clusters  (e.g., MS~1054-03, Donahue, Gioia, Luppino et al. 1997).  RX~J0911.4+0551 is  a new quadruply imaged QSO  with an  unusual image configuration. The lens configuration is complex, composed of one main lensing galaxy plus external shear possibly caused by a galaxy cluster at redshift between 0.6 and 0.8 and another possible group at $z>0.7$. Multi-object  spectroscopy is needed in  order  to confirm our cluster candidate/s and derive  its/their redshift and velocity dispersion. In addition, weak lensing analysis  of background galaxies might prove useful to  map the overall  lensing potential involved in this complex system."
    },
    "9803/astro-ph9803205_arXiv.txt": {
        "abstract": "This paper discusses the properties of scattering--dominated active galactic nuclei (AGN). We define these to be AGN for which the direct line-of-sight to the continuum source is obscured by Compton-thick material. The aim is to construct, for the first time, a model consistent with X-ray line luminosities, line ratios and various luminosity indicators. The \\ASCA\\ spectra of six such sources  show several X-ray lines that  can be reliably measured, mostly due to highly ionized magnesium, silicon sulphur and iron. These enable us to investigate the physical conditions of the scattering material. The sources show evidence of He-like and H-like iron lines that are likely to be produced in hot (T$\\sim 10^6$ K) photoionized gas. By measuring the EW of the lines, and by constructing a diagnostic line-ratio diagram, we demonstrate  that the silicon and magnesium lines are produced by the same gas emitting the highly ionized iron lines. The properties of this gas are rather different from the properties of  warm absorbers in type I AGN. Neutral 6.4 keV iron lines are also detected, originating in a different component which can be either Compton-thin or Compton-thick. The best measured iron lines suggest an enhancement of ${\\rm n_{Fe}/n_H}$ by a factor $\\sim 2$ compared to solar, in both the hot and cool Compton-thin components. We further show that in four of the sources, the Fe \\Ka(6.4~keV)/\\Hb\\ ($\\lambda 4861 \\AA$) line ratio is consistent with that predicted for typical narrow line region clouds, and the reddening corrected \\Hb\\ is known, provided the column density is larger than $\\sim 10^{22.5}$ \\cmii\\ , \\aox\\ is smaller than 1.3. For some sources, this is  a viable alternative to the commonly assumed Compton thick medium as the origin of the 6.4 keV iron line.  {\\it Subject headings:} galaxies:abundances - galaxies:Seyfert - galaxies:active - line:formation - X-ray:galaxies ",
        "introduction": "The optical and X-ray properties of type II AGN (i.e. those showing prominent narrow emission lines and very faint, if any, broad lines) have been discussed in numerous recent  papers that contain the analysis of their morphology, spectrum, geometry and relationship to type I (broad emission line) AGN. Various names, including Seyfert 2, narrow emission line galaxies (NELG), and narrow line X-ray galaxies (NLXGs) have been used to describe these objects. Obviously, there is some subjectivity in the classification of type II AGN leading to the assignment of a more than one ``type'' for some objects which have been studied by several authors. The geometry of the innermost region of such sources is a fundamental, yet still an open issue and the reader is referred to Antonucci (1993), and Mulchaey \\etal\\ (1994) for discussion and references regarding these questions.  X-ray observations  offer a unique view of type II AGN,  since X-rays can penetrate large column densities.  This has been a subject of much research, for example, see recent papers by Turner \\etal\\ (1997a,b,1998). These authors found some surprising similarities between the X-ray spectra of type I and type II AGN. In some type II sources, the equivalent width (EW) of the 6.4 keV line is similar to that observed in Seyfert 1s and the line profiles show broad, redshifted wings. In other type II sources, like NGC~1068, EW(Fe \\Ka) is an order of magnitude larger than in type I AGN, suggesting that the line is seen against a reduced continuum, presumably due to obscuration. Evidently, type II AGN fall into at least two X-ray  categories; those where the central source is directly observed below 10 keV, and those where it is not. Hereafter we refer to those type II AGN whose 0.5--10 keV spectra are dominated by scattered radiation, ``scattering--dominated AGN'', and they are the subject of this paper. We expect a good, but not necessarily a one-to-one correlation between such objects and those type II AGN who show highly polarized continuum and broad optical/UV emission lines, due to scattering (e.g. Tran 1995).  This paper investigates the properties of scattering--dominated AGN through detailed analysis of their 0.5--10 keV spectrum. We address the nature of the scattering medium  and try  to deduce its level of ionization, column density and covering fraction. We  also  investigate the metallicity of the gas and compare its properties to the ionized gas in Seyfert 1 galaxies, and to the narrow line region (NLR) gas. The analysis is aimed at a small number of scattering--dominated  AGN whose \\ASCA\\ spectra are of a sufficiently high quality to enable the measurement of at least 3 X-ray emission lines. It also suggests several new avenues for future study of such sources in preparation for the coming  {\\it AXAF}\\ and {\\it XMM}\\ missions. In \\S2 we discuss the predicted spectra of such sources. In \\S3 we compare predictions to a detailed analysis of the \\ASCA\\ spectra of six such galaxies. In \\S4 we discuss several implications of such a comparison, and implications for the state and the location of the scattering medium and its composition. ",
        "conclusions": "The following analysis is based on the line fluxes listed in Table 1 as well as on the analysis of the \\ASCA\\ spectrum of NGC~1068 (Netzer and Turner, 1997, Table 1). Obviously, the number of objects, and the number of measurable lines per object, are very small and the information content regarding the group properties is rather limited.  The most severe complication in analyzing the \\ASCA\\ spectra is  the likely contamination of the nuclear spectrum by extended, non-nuclear emission. This may be the result of hot gas in star forming regions, supernova remnants, or any other gas at T$\\simeq 10^7$~K. The limited \\ASCA\\ spatial resolution (with a half-power diameter $\\sim3$~arcmin), combined with the relative weakness of the scattered X-ray continuum, makes it almost impossible to separate the photoionized gas and hot plasma contributions. Indeed, some sources (e.g. NGC~1068) show clear indication of extended X-ray emission which is most likely due to starburst activity. This issue will not be resolved before {\\it AXAF} observations (with a half-power diameter $<$1~arcsec). Below we comment on the information obtained from the study of the Fe \\Ka\\ complex and address the potential use of diagnostic diagrams in the analysis of the spectrum of scattering--dominated AGN.  \\subsection{The \\Ka\\ complex and the iron abundance}  Except for NGC~1068, all our measurements of the 6--7 keV complex are somewhat ambiguous since we can not reliably resolve the highly ionized (6.7 and 6.96 keV) Fe \\Ka\\ lines from the neutral \\Kb\\ ($\\sim 7.1$ keV) line. Therefore, the analysis of the high ionization lines pertains to the {\\it combined intensity} of the H-like and He-like iron lines, which was obtained by subtracting the expected \\Kb\\ flux  (10\\% the flux of the 6.4 keV \\Ka\\ line) from the total. This makes the combined Fe{\\sc xxv-xxvi} line intensity in NGC~3488 consistent with zero because of the uncertainty on the \\Kb\\ flux (see Table 1) and the ones in Mkn~348 consistent with zero because of the large intrinsic error. The uncertainty in the relative strength of the 6.4 and 6.9 keV component is also affected by the uncertainty in the assumed, scattered broad 6.4 keV line.  The galaxies with measurable soft X-ray lines represent two different groups. In one object (NGC~1068), the highly ionized iron lines are comparable in strength to the 6.4 keV line. In three others (Circinus, Mkn~3 and NGC~4388) the combined intensity of the He-like and H-like iron lines is only about 10-15\\% the intensity of the low ionization component. Mkn~348 is possibly an intermediate case but the observational uncertainties are too large to tell. The intensity of the high ionization iron lines  in NGC~6240 is unknown, because of their  proximity  to the strong 7.1 keV absorption feature due to $2 \\times 10^{24}$ \\cmii\\ of neutral absorber. However, a comparison of the 6.4 keV intensity with the soft X-ray lines suggests that this source belongs to the same group as Mkn~3 and  Circinus. As for  the  lines of silicon, magnesium and sulphur, there are only three sources where reliable line ratios can be obtained and they all look quite similar (see below).  Regarding the 6.4 keV iron line, a comparison of the  measured line intensities with  the calculations shown in Figs. 2 and 3, indicate that in all sources where this line is much stronger than the 6.7--7.0 keV complex, the line can not originate in the same component producing the strong magnesium, silicon and sulphur lines. We therefore suggest that in those  sources, much of the 6.4 keV line emissivity is due to reprocessing in a large column of low-ionization gas. As for  NGC~1068, the 6.4 keV line in this source  may be due to warm (T$_e \\simeq 2 \\times 10^5$~K) photoionized gas (Marshall \\etal\\ 1993; Netzer \\& Turner 1997). This idea is in conflict with  the Iwasawa \\etal\\ (1997) model (see below). We conclude that in four (perhaps five) of the six sources, the reprocessing efficiency (\\Rf) of the neutral component is much greater (a factor of 5--10) than that of the ionized component.  Given the level of ionization of the gas producing the 6.4 keV iron line, and the shape of the ionizing continuum, we can estimate the iron abundance from the observed EW(Fe \\Ka) in two interesting limits. The first corresponds to Compton--thin gas and has been discussed by Krolik and Kallman (1987), Matt \\etal\\ (1996) and others. In this case, assuming negligible resonance line absorption and complete isotropy of the scattered continuum radiation (as appropriate for a case where the scattering medium is viewed at all possible angles), \\begin{equation} {\\rm EW(6.4~keV) \\simeq 3.17 [\\frac {1.11^{1-\\Gamma}}{2+\\Gamma}] [\\frac{F_Y}{0.3}] [\\frac{ n_{Fe}/n_H }{ 4 \\times 10^{-5} }] \\,\\, keV } \\,\\, , \\end{equation} where F$_{\\rm Y}$ is the fluorescence yield (of the order of 0.3 for low ionization iron). In deriving this expression we have adopted the small column density limit which allows us to neglect the absorption of the emitted \\Ka\\ photons on the way out. A much slower (logarithmic) dependence on ${\\rm n_{Fe}/n_H}$ is expected at very large columns.  The second case corresponds to Compton--thick gas, such as the walls of the hypothetical nuclear torus. This case has recently been calculated by Matt \\etal\\ (1996, 1997) for a range of metallicities and viewing angles. For a $\\Gamma=1.9$ continuum, larger than solar ${\\rm n_{Fe}/n_H}$ and $\\cos (i)=0.5$, where $i$ is the angle between the line of sight and the axis of the torus, Matt \\etal\\  estimate (see their Fig. 2), \\begin{equation} {\\rm EW(6.4~keV) \\simeq 1.8 (1+0.6 \\log [\\frac{ n_{Fe}/n_H }{ 4 \\times 10^{-5} }] ) \\,\\, keV .} \\end{equation} The range of observed angles  can change this value by up to a a factor of 1.5. Thus, for a typical $\\Gamma=1.9$ continuum, the Compton--thin and Compton--thick assumptions result in a factor 2 difference in the estimated iron abundance. The two cases  also predict very different 6-10 keV continuum shape since the Compton-thick gas produces a much stronger 7.1 keV absorption edge. Finally,  in  Compton-thick gas, some 10\\% of the 6.4 keV line intensity is in a broad red wing. Below we discuss the iron abundance under the two different scenarios. While the emphasis is on line intensities, we note  that our  continuum fits do not require the presence of a 7.1 keV absorption in any source except NGC~6240.  The following examples consider two gas components, one represents ``cold'' (small Ux) gas and the other ``hot'' (large Ux) gas. Each scatters the central  radiation and produces a scattered continuum. We refer to these as the ``two continuum components''. We first assume both components to be Compton-thin. In estimating the iron abundance from the observed EW, we note that the two continuum components are contributing at 6.4 keV and $\\Gamma$ is not directly measurable since scattering affects the observed continuum shape (\\S2.2). We use  the numerical calculations (Fig. 2) and assumed that the iron composition is identical in all components. We  also note  that resonance absorption is negligible, because iron is in a very low ionization state in the cold component and the column density is large in the hot component. Given these assumptions, we expect each component (hot or cold) to have a similar EW(Fe \\Ka)  {\\it relative to its own continuum}. Thus, in those sources  where the 6.7--6.96 keV iron lines are much weaker than the neutral 6.4 keV lines, most of the 6.4 keV continuum is due to reflection by the neutral component. Given the assumed continuum shape, this implies an iron over abundance by a factor 2--3 for Circinus and a factor of 1--1.5 for Mkn~3.   For NGC~6240 and NGC~4388, we estimate the EW relative to the scattered component by using our best fit model of these source(Fig. 6). According to the model, about half the observed 6.4 keV continuum is due to transmission and the other half due to scattering. We can thus estimate EW(6.4 keV) relative to the scattered component and deduce n$_{\\rm Fe}$/\\nh$\\simeq 1.5-2\\times$solar. Mkn~348 is so different in this respect that we cannot reliably estimate EW(Fe \\Ka) relative to the scattered continuum. The iron abundance in NGC~1068, assuming a Compton-thin gas, has been discussed by Marshall \\etal\\ (1993) and Netzer \\& Turner (1997), and found to be about three times solar. Thus, under the Compton-thin assumption, we have indications of  iron overabundance in 5 sources.  Regarding the Compton-thick case, the iron abundance inferred from the observed EW(\\Ka\\ 6.4 keV) line is about half the value deduced above,  i.e. consistent with solar for all sources. As we show below, the analysis of the highly ionized gas enables an independent check on the iron composition because the medium producing such lines is unlikely to be Compton-thick.   \\subsection{Diagnostic diagrams}  Diagnostic diagrams, involving various line ratios, have been used to separate Seyfert galaxies from galactic HII regions, and to search for the spectroscopic signature of LINERs (e.g. Baldwin, Phillips \\& Terlavich, 1981). Below we attempt to use the same method in the X-ray domain, in search for the unique signature of scattering--dominated AGN.  There are two major differences between our study and the investigation of the optical spectrum of LINERs and HII regions. First, the number of available X-ray lines  is very small and we can only measure, reliably, three line ratios and construct two such diagrams. Second, given the gas is photoionized, the X-ray line spectrum is dominated by recombination and not a single, purely collisionally excited line is strong enough to be used. Collisionally excited X-ray lines dominate the spectrum of hot plasmas with ratios vastly different from that expected in photoionized gas. Thus, there is no overlap in properties and line ratio diagrams can either be used for photoionized gas or for hot plasmas. In contrast, lower excitation photoionized nebulae contain a mixture of recombination and collisionally excited lines that provide very useful  diagnostics. Thus ratios like \\bOIIIb/\\Hb\\ have been used to derive the level of ionization of the gas and line ratios involving highly excited O$^{+2}$ transitions have been used to investigate the role of shock excited gas in the spectrum of LINERs and Seyfert galaxies (e.g. Ferland and Netzer 1983). This kind of analysis is not yet possible in the X-ray regime.  Fig. 7 shows a diagnostic diagram composed of the best observed line ratios in our sample, I(\\SiXIII)/I(\\FeXXV\\ + \\FeXXVI) versus I(\\MgXI)/I(\\SiXIII). The first ratio is a good indicator of regions of large ionization parameters, with electron temperature of the order of 10$^6$ K, and the second measures the conditions in lower ionization gas. Obviously, the excitation and ionization of \\MgXI\\ and \\SiXIII\\ is rather similar and future analysis, based on oxygen and neon lines, will be of greater use. Measurements of the two  ratios are available for  Circinus and NGC~1068 and an upper limit can be  obtained for Mrk~3 (see \\S2). The diagram shows the location of the three objects along with four theoretical curves, this time  assuming the L$_{\\rm E}\\propto {\\rm E}^{-0.5}$ continuum. Similar results, with appropriate scaling of Ux, are obtained for the L$_{\\rm E}\\propto {\\rm E}^{-0.9}$ continuum used in most other calculation. The curves are series of increasing Ux for various column density and gas composition. The  solid lines are standard composition models for three column densities, \\Ncol=10$^{22.4}$~\\cmii, 10$^{23}$~\\cmii\\ and 10$^{23.3}$~\\cmii. The dotted line is for \\Ncol=10$^{23}$~\\cmii\\ but with ${\\rm n_{Fe}}$/\\nh\\ three times larger.   Inspection of the line ratio diagram, and the observed spectra, suggest that: \\begin{enumerate} \\item\tThe observed line ratios cannot be simultaneously obtained in a single temperature collisionally ionized gas. Under such conditions, the temperature required to ionize Fe{\\sc xxv} and Fe{\\sc xxvi} is inconsistent with the observed strength of the silicon and magnesium lines. If all lines are produced in a single component, this gas must be photoionized. Indeed, \\ASCA\\ spectra of starburst galaxies (e.g. Ptak \\etal, 1997) generally show strong silicon and sulphur lines but little or no emission from neutral Fe \\Ka. \\item\tThe inferred ionization parameter, $U_X\\sim 1$, is 3--10 times larger than the ionization parameter of the highly ionized (warm absorber) gas in Seyfert 1 galaxies (George \\etal, 1998). This is in agreement with the finding of Turner \\etal\\ (1997a). We have examined the properties of this  gas and found a mean electron temperature of  about 10$^6$ K and no noticeable soft absorption features. Such gas, on the line of sight to a typical AGN continuum, with a column density not exceeding 10$^{23}$ \\cmii\\ (the column density of the great majority of warm absorbers in Seyfert 1 galaxies, see George \\etal\\ 1998) would escape detection by \\ASCA\\ type instruments. \\item\tThe diagnostic diagram cannot, by itself, be used to infer the Fe/Si abundance ratio since large column densities mimic the appearance of a small-column with large Fe/Si. As argued in \\S4.1, the analysis of the 6.4 keV lines suggest large ${\\rm n_{Fe}}$/\\nh\\ in all sources if the emitting medium is Compton-thin. If  the high ionization components have similar compositions, then according to the diagram, they must have relatively small column densities, perhaps similar to the lowest column shown in Fig. 7. \\item \tThe weakness of \\MgXII\\ and \\SiXIV\\ lines is somewhat surprising. The lines are predicted to be similar in strength to the lower ionization magnesium and silicon lines (Figs. 2 and 3) yet we could only obtain upper limits. The difficulty of detecting the \\SiXIV\\ line  may partly be explained by the notorious detector/mirror features in \\ASCA\\ around 2 keV. \\end{enumerate}  \\subsection{The scatterer location and the value of \\Rf}  So far we have focused on relative line intensities and line-to-continuum flux ratios. These  are useful in determining the ionization and composition but do not reveal the reprocessing efficiency, \\Rf, since EW(Fe 6.4 keV) is almost independent of the column density (Fig. 1). \\Rf\\ can not be obtained by comparing the absolute line flux with the intrinsic  luminosity since the latter is not known. However, there are several other luminosity indicators at longer wavelength, including the infrared flux and the \\bOIIIb\\ luminosity (e.g. Mulchaey \\etal, 1994, see extensive discussion in Turner \\etal, 1997b,1998), that can be used. Here we chose to use the reddening-corrected narrow \\Hb\\ line as our luminosity indicator. This  is  similar to the L(\\bOIIIb) method used in Turner \\etal\\ but enables a more direct comparison with the intrinsic ultraviolet luminosity.  Measurements of the narrow \\Hb\\ flux for the six sources, as well as for most known type II AGN, are readily available (Mulchaey \\etal, 1994; Polletta \\etal, 1996; and references therein). The above references contain also the measured \\Ha/\\Hb\\ line ratio which we use to correct for reddening and to obtain the intrinsic \\Hb\\ flux. In the following we assume  an intrinsic I(\\Ha)/I(\\Hb)=3.0 and a simple galactic type reddening with  A$_{\\rm V}$/E$_{\\rm B-V}$=3.1. Reddening corrected \\Hb\\ fluxes obtained this way are listed in Table 1.  A word of caution is in order. Applying a simple, screen-type reddening correction to the spectrum of Circinus and NGC~6240 is problematic since the observed Balmer decrement in both galaxies is very large (e.g. Fosbury and Wall 1979 for the case of NGC~6240). In addition, much of the \\Hb\\ flux in NGC~6240 is likely due to luminous star--forming regions in this galaxy. In both cases, and probably in many other narrow--line galaxies showing large Balmer decrements, the geometry is  rather complex with several clouds along each line--of--sight. We may be looking into  dusty environments for which a simple correction factor is inappropriate. This can  invalidate the  reddening-corrected \\Hb\\ fluxes used here.  The theoretical I(Fe \\Ka)/I(\\Hb) is easily obtained from the spectral energy distribution, the column density and the iron abundance. For low ionization, small Balmer optical depth gas, the number of \\Hb\\ photons is a known fraction (about 0.12  for Case B recombination) of the Lyman continuum photon flux and the number of Fe \\Ka\\ photons is a known fraction of the ionizing E$>7.1$ keV flux. The case of interest for this study is gas with very large Lyman optical depth yet relatively small hard X-ray  optical depth. Defining Q$_{7.1~keV}$ as the photon flux above 7.1 keV, and Q$_{13.6~eV}$ as the Lyman continuum photon flux, we get for this case \\begin{equation} {\\rm I(Fe \\Ka)/I(\\Hb)} \\simeq 1.5 \\times 10^3 \\exp(- \\tau_{\\rm 6.4 keV}) [\\frac{ {\\rm Q_{7.1~keV}}}{ {\\rm Q_{13.6~eV}}} ] [1-\\exp(-\\tau_{\\rm 7.1 keV})] \\,\\, , \\end{equation} where \\begin{equation} \\tau_{\\rm 7.1 keV} \\simeq 0.13 [\\frac{ {\\rm N_{col}} }{ 10^{23} } ] \\frac { ({\\rm n_{Fe}/n_H)} }{ 4 \\times 10^{-5} } \\,\\, , \\end{equation} is the 7.1 keV optical depth due to iron (assumed to be much smaller than unity), and $\\tau_{\\rm 6.4 keV}$ is the  absorption optical depth due to all metals, at 6.4 keV (of the same order as $\\tau_{\\rm 7.1 keV}$ for solar ${\\rm n_{Fe}/n_H}$). Thus at small $\\tau_{\\rm 7.1 keV}$, the line ratio increases with the iron abundance. Major complications arise due to  absorption of the E$>7.1$ keV photons by elements other than iron and by the non-negligible opacity at 6.4 keV, resulting in the destruction of emitted \\Ka\\ photons. This  makes the Fe \\Ka\\ emissivity sensitive to the covering fraction since the \\Ka\\ photons emitted by one cloud can be absorbed by another cloud. Having in mind the NLR conditions, we assume in the following  \\Cf=0.1.  Fig.~8 shows a series of calculated I(Fe \\Ka)/I(\\Hb) for N$_{\\rm E} \\propto {\\rm E}^{-1.9}$ continuum, solar metallicity, \\nh=10$^4$\\cc\\ applicable to the NLR, and various  \\aox. To enable a comparison with the expressions given above, we note that for those models with \\aox=1.3,  ${\\rm Q_{0.1-10~keV}}/{\\rm Q_{7.1~keV}} = 130$ and ${\\rm Q_{13.6~eV}}/{\\rm Q_{0.1-10~keV}}=52$.  The diagram shows that for \\Ux\\ $=10^{-3}$ and \\Ncol$>10^{21.7}$ \\cmii,   \\Hb\\  is already emitted at maximum efficiency while the Fe \\Ka\\ flux is proportional to the column density. For \\Ncol$>10^{23.5}$ \\cmii, both lines are emitted at maximum efficiency and their ratio reflects the spectral energy distribution, ${\\rm n_{Fe}/n_H}$ and the destruction of the Fe \\Ka\\ photons.  Inspection of Fig.~8  and  Table 1 suggests that in four sources, NGC~1068,  NGC~6240  Mrk~3 and Mkn~348, the I(Ka)/I(\\Hb) line ratio is below 0.2, which is consistent with  solar ${\\rm n_{Fe}/n_H}$  for \\Ncol$< 10^{23.5}$ \\cmii\\ for \\aox=1.1. Assuming an iron overabundance by 2--3 reduces the required column to below 10$^{23}$ \\cmii\\ for the same \\aox. Furthermore, in three of the four cases the required column can be substantially smaller than the above mentioned upper limit. The  column density is smaller if the UV bump is weaker than assumed and larger if \\aox\\ is larger than assumed. Thus in about half the sources the 6.4 keV iron line could originate in the NLR if the clouds in that region have column densities exceeding about 10$^{22.5}$ \\cmii. Circinus and NGC~4388 are different since  I(Fe \\Ka)/I(\\Hb) in those sources is an order of magnitude larger than in the other three. As shown in the diagram, this is unlikely to be due to a much larger column density. Either the X-ray source is very bright compared with the UV source (very small \\aox) or else there is an additional, large covering fraction, \\Ka\\ producing component which is neutral, very thick and inefficient \\Hb\\ emitter.  Current NLR models (see Ferguson \\etal\\ 1997 and references therein) do not make definite predictions regarding the size of the NLR clouds, since most observed narrow lines originate in  the highly ionized, H{\\sc ii} part of the clouds. An obvious complication is if the NLR gas is dusty (see Netzer and Laor 1993 and references therein). For example, it is conceivable that the NLR clouds are the illuminated faces of  dusty molecular clouds of significant column density. This would result in a reduced \\Hb\\ emissivity but the Fe \\Ka\\ line will hardly be affected. As already explained, the reddening correction for a dusty H{\\sc ii} region can differ substantially from the correction procedure applied here.  We have examined the much larger sample in Turner \\etal\\ (1997a) to estimate the intrinsic I(Fe \\Ka)/I(\\Hb) ratio in type II AGN. Out of 17 sources with reliable Fe \\Ka\\ and \\Hb\\ measurements (including the ones in this paper), 7 show  reddening corrected I(Fe \\Ka)/I(\\Hb)$<0.2$ which we consider consistent with origin in the NLR of these galaxies. The remaining sources show a larger ratio that requires \\Rf\\ in excess of what is expected from the NLR gas. The neutral Fe \\Ka\\ line in those is likely due to absorption by a larger column density, very neutral material that is either the walls of the central torus or large molecular clouds in the nucleus.   Given the likely origin of the 6.4 keV line, we can now comment on the nature and location of the gas producing the high ionization iron lines. Assuming the same ${\\rm n_{Fe}/n_H}$ in both components, we can derive  \\Rf(hot)/\\Rf(cold). This is found to be about 1 for NGC~1068 and about 0.1 for Mkn~3 and Circinus. The uncertainty  is about a factor 2 since, as explained, the scattering efficiency differs by about this factor when comparing Compton-thin and Compton-thick gas. We  further consider possible combinations of covering factor and note that for Compton-thin gas, \\Rf$\\propto$\\Ncol$\\times$\\Cf. This, combined with \\Ux(cold)/\\Ux(hot) (about 10$^{-3}$ with a large uncertainty,  see Fig. 1 and the parameters used earlier for the NLR gas), enables us to estimate several likely combinations of these quantities.   An interesting possibility involving the NLR idea, is that \\Ncol(hot)$\\sim 10^{-2}$\\Ncol(cold). This would imply \\Cf(hot)$\\simeq 10$\\Cf(cold), i.e. \\Cf(hot)$\\simeq$1. In this case, the hot and cold components coexist, spatially, and the large \\Cf(hot) does not allow a torus with a small opening angle. The typical NLR density is about $10^4$~\\cc, thus  \\nh(hot)$\\sim 10 {\\rm cm^{-3} }$. The physical thickness of the hot gas in this case is of order 10-100 pc, i.e. of the same order of the NLR size. We note, however, that the two components are not in pressure equilibrium since nkT$_{\\rm e}$(cold)$\\simeq 10$nkT$_{\\rm e}$(hot). Another possibility is that \\Cf$\\sim 0.1$ in both components and \\Ncol(cold)$\\simeq 10$\\Ncol(hot). This does not allow a co-spatial existance of the two components. Finally, the 6.4 keV line may be from the thick torus walls. The efficiency in this case is very large and suggests that the fraction of this wall visible to us is extremely small. It also suggests that the hot gas completely fills the opening in the torus.   Acknowledgments: It is a pleasure to acknowledge stimulating and useful discussions with our colleagues R. Mushotzky, K. Nandra, T. Yaqoob and T. Kallman. A very useful referee report helped us improve the presentation of this paper. This research is supported by the Universities Space Research Association (TJT, IMG) and by a special grant from the Israel Science Foundation (HN).  \\newpage"
    },
    "9803/astro-ph9803027_arXiv.txt": {
        "abstract": "New multi-epoch, mid-infrared (8-13\\,$\\micron$) spectrophotometric observations are presented for 30~late-type stars.  The observations were collected over a four year period (1994-1997), permitting an investigation of the mid-infrared spectral shape as a function of the pulsation cycle (typically 1-2 years).  The spectra of stars with little excess infrared emission and those with carbon-rich dust show the least spectral variability, while stars with evidence for dusty, oxygen-rich envelopes are most likely to show discernible variations in their spectral profile. Most significantly, a large fraction of variable stars with strong 9.7\\,$\\micron$ emission features show clear spectral profile changes which repeat from one cycle to the next.  The significant sharpening of the silicate feature near maximum light can not be fully explained by heating and cooling of the circumstellar dust shell during the pulsational cycle, suggesting that the dust optical properties themselves must also be varying.  In addition, the appearance of a narrow emission feature near the silicate peak for a few stars may require the production of especially ``pure'' silicate dust near maximum light.  The general narrowing of the silicate feature observed may reflect the evolution of the pre-existing dirty grains whose surface impurities have been evaporated off when the grain temperature rises preceding maximum light.  An improved theory of dust formation which can explain the observed changes in the grain properties around a single, pulsating star may lead to a definitive explanation for the diversity of silicate emission profiles observed amongst oxygen-rich, late-type stars. ",
        "introduction": "The mid-infrared (8-13\\,$\\micron$) spectra of late-type stars have been measured by many observers since the development of infrared detectors.  These red giants and supergiants are often surrounded by dusty envelopes which absorb stellar radiation and re-radiate the energy in the near- and mid-infrared.  The infrared spectra can be classified based on the chemical content of the circumstellar environment (oxygen- or carbon-rich) and on the optical thickness of the dusty envelope (e.g., Merrill \\& Stein 1976a,b).  Oxygen-rich circumstellar environments often produce spectra evincing a feature near 9.7\\,$\\micron$ resulting from the presence of silicate dust (Woolf \\& Ney 1969).  This feature appears in emission for optically thin envelopes or in absorption when large enough optical depths are encountered.  The emission spectra of dust surrounding carbon stars are nearly featureless, although often containing an 11.3\\,$\\micron$ feature attributed to SiC.  In addition, many of these red giants and supergiants are classified as long-period variables, pulsating with a typical period of 1-2 years.  The homogeneous set of survey measurements by the Infra-Red Astronomical Satellite (IRAS) in the mid-1980s allowed observers to classify silicate emission features based on various schemes (IRAS Science Team 1986; Little-Marenin \\& Little 1988, 1990; Goebel et al. 1989; Sloan \\& Price 1995).  The different shapes of the feature have been interpreted largely as due to differences in the chemical make-up of oxygen-rich dust.  Unfortunately the IRAS program did not conscientiously include observations of the mid-infrared spectra of long-period variable stars at different phases of their luminosity cycles, and there are only a few cases where such data have been retrieved from the IRAS Low Resolution Spectrometer (LRS) database. These observations have suggested silicate feature strength variations as a function of pulsational cycle, but have been hampered by limited temporal coverage (Little-Marenin, Stencel, \\& Staley 1996).  More recent results by Creech-Eakman et al. (1997) point towards evidence for variations in the silicate feature as a function of pulsational phase, but the comparison spectra were taken nearly a decade apart. Hence, the simple observational question of whether the mid-infrared spectra of LPVs change shape through the pulsational cycle has been left without a decisive answer.  A campaign of observations taken from 1994 to 1997 was designed to monitor the mid-infrared spectrum of nearly 30 late-type stars.  The observations,  sampling the spectrum of most stars multiple times within a pulsational cycle, used the same instrument and observing technique. The homogeneity of this data set is important for allowing reliable spectral comparisons, avoiding complicating issues such as different apertures and calibration methods.  This paper presents the full data set collected thus far and discusses the spectral variability (or lack of variability) of our sample stars. ",
        "conclusions": "The mid-infrared spectra of 30~late-type stars have been monitored in order to detect changes occurring on the pulsational time scale (typically 1-2~years) of long period variables (LPVs).  Stars which exhibited little or no bolometric variability (i.e. non-LPVs) generally showed no change in their spectral shape in the range 8-13\\,$\\micron$. Furthermore, most stars with no strong 9.7\\,$\\micron$ silicate feature, including carbon stars and oxygen-rich miras with broad, weak silicate features, showed no spectral shape change.  However, a few such stars in this category displayed either enhanced variability as a function of wavelength (IRC~+10216 and R~Leo) or a detectable change in the spectral slope correlated with pulsational phase (R~Cnc and W~Aql). The former effect has no clear explanation, while the latter effect can be explained by changes in the dust shell temperature ($\\Delta T_{\\rm{dust}}\\simle 200$\\,K).  The most significant result presented here is that nearly all of the observed sources with clear 9.7\\,$\\micron$ silicate features and definite bolometric variability showed strong evidence for changes in the silicate emission strength and spectral profile which are correlated with pulsational phase.  We conclude that silicate emission variation is a general property of long-period variables with optically thin silicate features.  The sharpening of the silicate feature near maximum light and its subsequent broadening can be explained by the heating and cooling of the dust envelope coupled with changing optical constants for the dust grains. The appearance of a spectrally narrow emission feature near the silicate peak of a few stars strongly indicates the existence of ``pure'' silicate dust grains near maximum light.  We hypothesize that the general narrowing of the dust emission spectra may arise from pre-existing dirty grains whose surface impurities have been evaporated off or whose amorphous molecular configuration has crystallized during the dust re-heating following minimum luminosity.  The solid-state resonance would naturally broaden after maximum light as impurities re-adsorb onto the cooled grain surface.  Such speculation awaits more detailed laboratory measurements of astrophysically-relevant grain types.  The observations presented here remind us that the dust formation process is still only partially understood.  Indeed, uncertainties in the optical constants for circumstellar dust are a primary obstacle in creating self-consistent multi-wavelength radiative transfer models incorporating interferometric observations of dusty objects (e.g., Monnier et al. 1997).  The changes observed in silicate optical properties as a function of pulsational phase are not predicted by any present dust formation theories, and more careful consideration is required of the effects of photospheric shocks propagating into the dust formation zone and of the changing temperature and density structure due to the pulsation.  Such models may then not only explain the changing optical properties of the dust around a single, pulsating object, but may also explain why different stars possess silicate emission with distinctly different spectral profiles."
    },
    "9803/astro-ph9803094_arXiv.txt": {
        "abstract": "The detection of rapid variability on a time\\-scale of hours in radio-quiet quasars (RQQSOs) could be a powerful discriminator between starburst, accretion disc and relativistic jet models of these sources. This paper contains an account of a dedicated search for rapid optical variability in RQQSOs. The technique used differential photometry between the RQQSO and stars in the same field of view of the CCD. The 23 RQQSOs that were observed all have high luminosities ($-27<M_\\mathrm{V}<-30$), and 22 of these sources are at redshifts $z>1$. The total amount of observation time was about 60 hours and these observations are part of an ongoing programme, started in September 1990, to search for rapid variability in RQQSOs. No evidence for short-term variability greater than about 0.1 magnitudes was detected in any of the 23 sources, however long-term variability was recorded for the radio-quiet quasar \\object{PG 2112$+$059}. The finding charts are included here because they identify the RQQSO and the reference stars used in the photometry, and hence are available for use by other observers.   The unusual properties of two RQQSOs that were not included in our source list are noted. X-ray results reveal that \\object{PG 1416$-$129} is variable on a timescale of days and that the remarkable source \\object{IRAS 13349$+$2438} varied by a factor of two on a timescale of a few hours. The latter source displayed blazar type behaviour in X-rays and implies that relativistic beaming may occur in at least some RQQSOs. Radio results also indicate the presence of jets in at least some RQQSOs. ",
        "introduction": "There is a general consensus that quasars belong to two different radio populations, radio-quiet quasars (RQQSOs) and radio-loud quasars. $R$ is usually defined as the ratio of the radio (6~\\mbox{cm}) to the optical (440~\\mbox{nm}) flux densities and the radio-quiet quasars have a value of $R<10$, while the radio-loud quasars have $R>10$ (Kellermann et al. 1989). It is found that $\\sim10~\\mbox{\\%}$ of quasars are in the radio-loud category. An additional distinction between active galactic nuclei (AGN) with strong and weak radio sources comes from the observation that radio loud objects essentially all occur in elliptical galaxies and RQQSOs appear to reside in galaxies that are dominated by exponential disks. However the RQQSOs that occur in elliptical host galaxies are in general more luminous than those that reside in disks (Taylor et al. 1996).   Little is known about the short-term variability of ra\\-dio-quiet quasars, because few studies have been carried out (Gopal-Krishna et al. 1993 and 1995; Jang \\& Miller 1995; Sagar et al. 1996). In contrast blazars display rapid variability in the wavelength range from radio to gamma rays. The blazar class encompasses both optically-violent\\-ly-variable (OVV) quasars and BL Lac objects and about one quarter of all radio-loud quasars are also in the blazar category (Webb et al. 1988; Pica et al. 1988).   There are many theoretical models which endeavour to explain the large and rapid variability exhibited by blazars and these are usually divided into extrinsic and intrinsic categories. One extrinsic mechanism is microlensing of emission knots in a relativistic jet when they pass behind planets in an intervening galaxy (McBreen \\& Metcalfe 1987; Gopal-Krishna \\& Subramanian 1991).  The rapid variability from superluminal-microlensing may be \\linebreak responsible for the variability observed in \\object{AO 0235$+$164} (Rabbette et el. 1996) and \\object{PKS 0537$-$441} (Romero et al. 1995). One family of intrinsic models is based on a rotating supermassive black hole which accretes matter from a surrounding accretion disc and ejects two oppositely directed jets. The shocked-jet model involves shocks which move with relativistic speeds along the jet (Qian et al. 1991; Marscher, 1980).  It is believed the shock propagates along the line of sight, through inhomogeneous, small-scale structures distributed along the jet.  These inhomogeneous structures are illuminated, or excited, by the moving relativistic shock, through the amplification of the magnetic field and the acceleration of electrons which causes the variability in polarization and in flux density that are observed over a wide range of frequencies (Hughes et al. 1986).  Another family of intrinsic models invokes numerous flares or hotspots in the accretion disk and the corona that is believed to surround the central engine (Wiita et al. 1992; Mangalam \\& Wiita 1993) and indeed a similar model has been proposed to explain X-ray variations in blazars (Abramowicz et al.  1991). The fact that RQQSOs generally lie on the far-infrared versus radio correlation (Sopp \\& Alexander 1991) suggest that star formation plays an important role in their radio emission. It has been suggested by Terlevich et al. (1992) that the low values of $R$ in RQQSOs can be explained without jets or accretion discs, by postulating a circumnuclear starburst within a dense, high-metallicity nuclear environment. In this model the optical/UV and bolometric luminosity arises from young stars; the variability comes from cooling instabilities in the shell of compact supernova remnants and supernova flashes. Variability on intranight timescales is however difficult to explain with this model because of the short timescales involved.  Furthermore radio-quiet and radio-loud quasars have very different radio power outputs but have similar spectral shapes in the radio region and suggest that a significant fraction of the RQQSOs may be capable of producing powerful radio emission (Barvainis et al. 1996).  Kellermann et al. (1994) found possible radio extensions up to about 300~\\mbox{kpc} in a few RQQSOs and assert that for at least these few cases, the emission is too large to be starburst related (Stein 1996).   Recently, some evidence suggesting rapid optical variability in the RQQSOs \\object{PG 0946$+$301} and \\object{PG 1444$+$407} was reported by Sagar et al. (1996). They also reported long-term variability for four RQQSOs. Jang \\& Miller (1995) reported intranight variability for one RQQSO out of a sample of nine sources. Brinkmann et al. (1996) obtained ASCA observations of the radio-quiet, infrared \\linebreak quasar \\object{IRAS 13349$+$2438} and detected substantial X-ray variability on a timescale of only a few hours.   The results of the photometric observations of a sample of mainly high luminosity and high redshift RQQSOs are presented. The observations and data reduction are given in Sect.~2. The results including tables listing the differential photometry and some light curves are presented in Sect.~3. The discussion and conclusions are given in Sects.~4 and 5. Sect.~4 also includes a discussion on two remarkable RQQSOs, \\object{PG 1416$-$129} and \\object{IRAS 13349$+$2438}. CCD images of the fields containing the radio-quiet quas\\-ars and reference stars used in the differential photometry are also included. A value of $\\mathrm{H}_\\mathrm{0}=50~\\mbox{km s}^{-1}~\\mbox{Mpc}^{-1}$ and $\\mathrm{q}_\\mathrm{0}=0.5$ has been adopted. ",
        "conclusions": "A long-term survey of a sample of high luminosity ($-27<M_\\mathrm{V}<-30$) and medium to high redshift ($0.466<z<4.7$) radio-quiet quasars was undertaken, in order to search for short and long term optical variability on timescales of hours to years. A large sample of 23 RQQSOs was observed, with a total observation time of about 60 hours spread over a period of several years. Long-term variability was detected in the RQQSO \\object{PG 2112$+$059} when it varied by 0.18 magnitudes in the V-band between 1992 and 1996. No rapid variability was observed in any of the sources in this sample of RQQSOs. The finding charts are included because they identify the RQQSO and the reference stars used in the photometry and hence are available to other observers.   There have been a few reports of rapid optical variability in a number of RQQSOs but these reports are not in conflict with the results presented here because such small variability would not have been detected in many of the sources monitored in this survey. The unusual properties of two sources are highlighted. These sources were not monitored in this survey but have recently been added to the list of sources for study. The remarkable source \\object{IRAS 13349$+$2438} combines some of the properties of blazars and radio quiet quasars and hence further observations may elucidate the nature of this hybrid source. The two unusual sources have $R$ values near the top of the range for RQQSOs and also have unusual radio spectra that may signify the presence of several source components. Further observations with VLA and VLBI should reveal new and enlightening views on the radio properties of these sources."
    },
    "9803/astro-ph9803230_arXiv.txt": {
        "abstract": "This paper presents \\ci, \\twcoft, and \\twcott\\ maps of the barred spiral galaxy M83 taken at the James Clerk Maxwell Telescope. Observations indicate a double peaked structure which is consistent with gas inflow along the bar collecting at the inner Lindblad resonance. This structure suggests that nuclear starbursts can occur even in galaxies where this inflow/collection occurs, in contrast to previous studies of barred spiral galaxies. However, the observations also suggest that the double peaked emission may be the result of a rotating molecular ring oriented nearly perpendicular to the main disk of the galaxy. The \\twcoft\\ data indicate the presence of warm gas in the nucleus that is not apparent in the lower $J$ CO observations, which suggests that \\twcooz\\ emission may not be a reliable tracer of molecular gas in starburst galaxies. The twelve \\ci/\\twcoft\\ line ratios in the inner 24$''\\times 24''$ are uniform at the 2$\\sigma$ level, which indicates that the \\twcoft\\ emission is originating in the same hot photon-dominated regions as the \\ci\\ emission. The \\twcoft/\\jtt\\ line ratios vary significantly within the nucleus with the higher line ratios occurring away from peaks of emission along an arc of active star forming regions. These high line ratios ($>$1) likely indicate optically thin gas created by the high temperatures caused by star forming regions in the nucleus of this starburst galaxy. ",
        "introduction": "The circumnuclear regions of galaxies are very often the setting for starbursts and other extraordinary events. Observations of the gas kinematics and distribution indicate that bars, resonances, gas inflow, and tidal shear play important roles in the formation and evolution of nuclear starbursts (\\eg\\ Handa \\etal\\ 1990; Kenney \\etal\\ 1992). Previous observations show that the molecular gas in the central regions of barred spiral galaxies often does not extend all the way into the centre of the nucleus. It accumulates some distance away from the centre, giving the emission a double peaked appearance with each peak occurring where the bar meets the nucleus (Kenney \\etal\\ 1992; Ishizuki \\etal\\ 1990). Dynamical models indicate that in the presence of a barred potential, gas will flow inward along the bar and slow its descent temporarily near inner Lindblad resonances (ILR, \\eg\\ Combes 1988; Shlosman, Frank \\& Begelman 1989). At these locations the gas may accumulate into larger complexes of molecular clouds.  In addition to complex dynamics, there is most likely a profusion of complicated photo-chemistry occurring within the nuclei of starburst galaxies. Interstellar clouds are believed to consist of smaller high density dark cores interspersed throughout a larger region of lower density. These high density regions are self-shielded from ultraviolet (UV) radiation that tends to dissociate molecular gas. The result is the population of the diffuse region by hydrogen atoms (\\ion{H}{1}), atomic and ionized carbon (C, C$^+$), and many other atomic and ionized species (\\eg\\ Morton \\etal\\ 1973), while the dark cores can contain molecular species such as H$_2$ and CO.  It is believed that atomic carbon can exist only in a small energy window, outside of which it will either be ionized or combined with oxygen to form CO. Inside the dense cores, most of the carbon combines to form CO, while outside the cores, the UV radiation acts to ionize atomic carbon. It is therefore expected that \\ion{C}{1} is the dominant species near the edges of dense, self-shielding cloud cores.  This simple model fails to explain the extended \\ci\\ emission observed in molecular clouds in our own Galaxy (Plume, Jaffe, \\& Keene 1994; Keene \\etal\\ 1985). One possible explanation is found in the clumpy structure of molecular clouds (\\eg\\ Stutzki \\& G\\\"usten 1990). This clumpiness would allow UV radiation to penetrate much deeper into the cloud allowing atomic carbon to exist at depths greater than would be allowed by a simple spherical model of molecular clouds (\\eg\\ Boisse 1990). Many alternative explanations have been proposed to explain the extended \\ci\\ emission. These ideas range from complicated chemical processes involving H$^+$ (Leung, Herbst, \\& Heubner 1984) to simpler ideas such as a C/O ratio greater than one (\\eg\\ Keene \\etal\\ 1985).  This paper presents \\ci\\ and CO maps of the barred spiral galaxy M83. Its low inclination angle ($i$ = 24\\arcdeg, Comte 1981) and close proximity ($D$ = 4.7 Mpc, Tully 1988) make it one of the best locations for studying the response of gas to a barred potential. It is believed to be undergoing a nuclear starburst (\\eg\\ Talbot \\etal\\ 1979), which may have been triggered by molecular gas inflow along the bar potential. This nuclear starburst would produce higher temperatures which should readily excite the higher $J$-transitions in the CO gas. Also, strong UV flux has been detected in the nucleus (Bohlin \\etal\\ 1983), which would help dissociate the CO into atomic carbon. By studying the CO and \\ci\\ data, we can understand better the dynamics of the gas and its role in fueling the nuclear starburst and also learn about the conditions conducive to the formation of atomic carbon. ",
        "conclusions": "}  This paper presents \\ci, \\twcoft, and \\twcott\\ maps of the nucleus of the barred spiral galaxy M83 taken at the JCMT. The main results are summarized below.  \\begin{enumerate}  \\item{We observe a double peaked structure in the molecular emission consistent with gas inflow along the bar collecting at the inner Lindblad resonance. The \\twcoft\\ emission suggests that some of the molecular gas has made it into the nucleus and is being heated by and possibly fueling the nuclear starburst. This result indicates that nuclear starbursts may occur even in galaxies which exhibit a double peaked emission structure, in contrast to the findings of Kenney \\etal\\ (1992).}  \\item{We observe different morphologies in the \\twcoft\\ channel maps than in the \\twcott\\ and \\twcooz\\ channel maps. These data suggest that \\twcooz\\ emission may not always be a good tracer of molecular gas in starburst galaxies, as the CO may be heated sufficiently to produce little emission in the \\joz\\ line. Thus, discretion should be applied in the interpretation of \\twcooz\\ emission as a tracer of molecular gas in starburst regions.}  \\item{The observations also suggest that the double peaked emission may be the result of a molecular ring out of the plane of the galaxy oriented nearly perpendicular to the main disk. This torus of cooler gas would need to contain a disk of hotter gas that fills its central void in order to explain the observed morphology.}  \\item{The \\ci\\ line strength indicates carbon column densities of $1 \\times 10^{18}$ cm$^{-2}$ while CO emission indicates CO column densities of $\\sim 3 \\times 10^{18}$ cm$^{-2}$ and H$_2$ column densities of $3 \\times 10^{22}$ cm$^{-2}$ at the peak of the emission. The $N$(C)/$N$(CO) ratio at this location is 0.33 $\\pm$ 0.10 which is similar to those found in other starburst galaxies.}  \\item{The twelve \\ci/\\twcoft\\ line ratios in the inner 24$''\\times 24''$ are uniform at the 2$\\sigma$ level and have an average value of 0.25 $\\pm$ 0.03, similar to those of other starburst galaxies. The uniformity of the line ratios suggests that both the high-excitation CO emission and atomic carbon form in photodissociation regions in the starburst nucleus.}  \\item{The \\twcoft/\\twcott\\ integrated intensity line ratios vary substantially over the central region of M83 with the highest ratios occurring towards the edges of the emission peaks. The \\twcoft/\\jtt\\ line ratios seem to be enhanced along an arc of active star forming regions. The high line ratios ($>$1) indicate that the higher $J$ transitions of CO are optically thin and are likely the result of the high temperatures and/or densities associated with star formation. }  \\end{enumerate}"
    },
    "9803/astro-ph9803006_arXiv.txt": {
        "abstract": "Using the IRAM interferometer we have mapped at high resolution ($2\\farcs2 \\times 1\\farcs2$) the $^{12}$CO(1--0) emission  in the nucleus of the doubled barred SABbc spiral M~100. Our synthesized map includes the zero spacing flux of the single--dish 30m map (Sempere \\& Garc\\'\\i a--Burillo, 1997, {\\bf paper I}). Molecular gas is distributed in a two spiral arm structure starting from the end points of the nuclear bar ($r=600$ pc) up to $r=1.2$ kpc,  and a  central source ($r\\sim$100 pc). The kinematics of the gas indicates the existence of a steep rotation curve (v$_{rot}$=180 km\\, s$^{-1}$ at $r\\sim 100$ pc) and strong streaming motions characteristic of a trailing spiral wave inside corotation.  Interpretation of the CO observations and their relation with stellar and gaseous tracers (K, optical, H$\\alpha$, H\\,I and radiocontinuum maps) are made in the light of a numerical model of the clouds hydrodynamics. Gas flow simulations analyse the gas response to a gravitational potential derived from the K-band plate, including the two nested bars. We develop two families of models: first, a single pattern speed solution shared by the outer bar+spiral and by the nuclear bar, and secondly, a two independent bars solution, where the nuclear bar is dynamically decoupled and rotates faster than the primary bar.  We found the best fit solution consisting of a fast pattern ($\\Omega_f$=160 kms$^{-1}$kpc$^{-1}$) for the nuclear bar (with corotation at R$^{F}_{COR}$=1.2 kpc) decoupled from the slow pattern of the outer bar+spiral ($\\Omega_f$=23 kms$^{-1}$kpc$^{-1}$) (with corotation at R$^{S}_{COR}$=8-9 kpc). As required by non-linear coupling of spirals (Tagger et al 1987), the corotation of the fast pattern falls in the ILR region of the slow pattern, allowing an efficient transfer of molecular gas towards the nuclear region. Solutions based on a single pattern hypothesis for the whole disk cannot fit the observed molecular gas response and fail to account for the relation between other stellar and gaseous tracers. In the two-bar solution,  the gas morphology and kinematics are strongly varying in the rotating frame of the slow large-scale bar, and fit the data periodically during a short fraction (about 20\\%) of the relative nuclear bar period of 46 Myr. ",
        "introduction": "The advent of high-sensitivity near-infrared imaging of galaxies has shown that a significant percentage of barred spirals host secondary bars in their nuclei. There could be two interpretations of the {\\it bars within bars} phenomenon, according to the relative pattern speeds of the two bars (Friedli and Martinet, 1993; Friedli and Benz, 1993 and 1995; Combes, 1994). The two patterns could be corotating if they are about parallel or perpendicular to each other. If the secondary inner bar is strongly misaligned with the primary outer bar, the two bars are likely to have distinct wave pattern speeds, as shown by numerical simulations. The decoupling of an inner faster pattern appears in self-consistent simulations with gas and stars thanks to the role of the dissipative component: as a result of gas inflow, under the action of the bar gravitational torques, mass accumulates onto the x$_2$ types of orbits which weakens the primary bar. The rotation period becomes much shorter in the nuclear regions due to mass concentration, which leads to the decoupling of a fast-rotating bar. Eventually, the nuclear bar destroys itself or it destroys the primary bar, modifying the overall disk potential. Evolution can then occur in much less than a Hubble time, and galaxies change their morphological type along the Hubble sequence. They change from barred to un-barred, and also they concentrate mass in the process, evolving slowly from late-types to early types.   The observation and modelling  of {\\it real} barred galaxies offers the opportunity to test theory predictions on galaxy evolution and it seems a necessary complement to numerical simulations of {\\it model} galaxies. The present work is intended to bring a combined observational and modelling effort in the study of the nearby barred spiral M100 (NGC4321). In this galaxy, classified as SABbc by de Vaucouleurs et al (1991), the hypothesis of a single mode common to the whole disk is dubious, both from observational and theoretical evidences.  \\bigskip  On the observational side, the nuclear region of M100 (up to r=3kpc) has been so far the subject of numerous studies. The pioneering work of Arsenault and collaborators (1988, 1989, 1990) established a connection between the ring-like H$\\alpha$ morphology of the nucleus and the existence of Inner Linblad Resonances. Further steps in sensitivity made appear, first, a four-armed structure (Cepa et al, 1990) and recently a fragmented two spiral arm structure in H$\\alpha$ (Knapen et al, 1996). The 6cm radio-continuum VLA map of Weiler et al (1981) shows also a two arm spiral pattern.  Near infrared images of the nucleus (Pierce 1986, Shaw et al 1995, Knapen et al 1995 (hereafter {\\bf K95}), Rauscher 1995) show the existence of either a secondary nuclear bar in K (nearly parallel to the main bar) together with a leading spiral structure, or an inner oval in the I band (with principal axes misaligned with respect to the outer bar). The synthesis aperture $^{12}$CO(1--0) maps (Canzian, 1992; Rand, 1995; Sakamoto et al 1995) indicate the existence of a two-arm molecular spiral structure connected to the K nuclear bar end points. The IRAM 30m map of {\\bf paper I} shows a strong concentration of CO emission towards the nuclear disk {\\bf ND}, a component clearly distinguishable from the main bar. A steep rotation curve gradient, unresolved by the 30m beam (12\\arcsec\\, in the 2--1 line), indicates a high mass concentration in the {\\bf ND}.  \\bigskip  On the modelling  side, Garc\\'\\i a-Burillo et al (1994) (hereafter  called {\\bf GB94}) and Sempere et al (1995) (hereafter {\\bf S95}) made numerical simulations of the cloud hydrodynamics to study the evolution of the molecular gas disk under the action of a realistic spiral+barred potential derived from a red band plate. The authors assume the whole disk to be fitted by a single well defined wave pattern characterized by $\\Omega_p$, shared by the primary bar and the spiral arms. However they lacked first, of the necessary spatial resolution and secondly, of a fair potential tracer to analyse the gas response in the inner 500 pc.   {\\bf K95} have made numerical simulations of the stellar and gas dynamics in M100, using a {\\it model} potential which departs markedly from the real mass distribution. Although they favour a one bar mode scenario their model fails to reproduce the molecular gas distribution observed by the interferometer.  \\bigskip  We present here a combined single-dish and interferometer data set fulfilling both high resolution (2.2\\arcsec$\\times$1.2\\arcsec) and sensitivity requirements. Contrary to the synthesis aperture maps so far published, we recover entirely the zero-spacing flux of the {\\bf ND}. The comparison between the different gaseous and stellar tracers of the {\\bf ND} is reexamined in this work. Particular attention is paid to the bias introduced by extinction in optical and even near-infrared images of the nucleus, and what might be the implications on the interpretation of the data. Observations are confronted to the result of new numerical simulations of the clouds hydrodynamics, based on a mass distribution directly derived from the infrared luminosity image of M100. We will focus on the feasibility of two independent patterns in the disk and how this scenario accounts better for the observations. ",
        "conclusions": "Simulations of the H$_2$ cloud hydrodynamics in the double barred system M100 have shown that the ensemble of observations (optical, infrared, HI and CO maps) are best explained by a two {\\it independent bars} scenario. The primary stellar bar (of 4.5 kpc radius) and the outer spiral structure share a common pattern speed of $\\Omega_s$=23\\,kms$^{-1}$kpc$^{-1}$ which places corotation at R$_{COR}^S$=8-9 kpc, i.e. beyond the bar end-points though well inside the optical disk. Although the nuclear stellar bar is mostly aligned with the primary bar (within 20\\deg) it has been shown to lead a fast pattern rotating at $\\Omega_f$=160\\,kms$^{-1}$kpc$^{-1}$, having corotation at R$_{COR}^{F}$=1.2 kpc radius. Both modes are dynamically decoupled and they show overlapping of their major resonances: corotation of the fast mode falls well within the ILR region of the slow mode.  The present model explains the efficient gas transport towards the nucleus, suggested by the interferometer observations, as a consequence of secular evolution driven by the stellar bar. Molecular gas crosses the ILR region of the slow pattern, spiraling inwards and forming a trailing spiral structure and an ultracompact source encircled by the ILR of the fast pattern (R$_{iILR}^{F}$=2.5$\\arcsec$). Alternative solutions are unable to account for the CO observations. In particular, in the slow pattern solution gas is stopped at the ILR barrier and forms a nuclear ring outside the {\\bf ND} extent. No central gas condensation is formed either.  The fast pattern solution proposed by K95 ($\\Omega_p$=70\\,kms$^{-1}$kpc$^{-1}$) worsens the fit for the outer bar+spiral structure found by {\\bf GB94}. In addition, two independent methods based on the morphology of the residual velocity field for the gas ({\\bf S95}) and the identification of spurs in optical pictures (e.g. Elmegreen et al 1992) confirm the value of R$_{COR}^S$ reported above.  We conclude that the gas response derived from the CO interferometer map, and the relation between the different stellar (K image) and gaseous tracers of the {\\bf ND} (H$\\alpha$) are best explained by the two pattern model. In particular, it explains the high CO concentration in the central part. This gas concentration could be eventually the cause of the nuclear bar destruction in this fastly evolving galaxy (see Norman et al 1996).  {\\it Acknowledgements}. This work has been partially supported by the Spanish CICYT under grant number PB96-0104. We thank J. Knapen for providing us with the H\\,I, H$\\alpha$ and K-band images used in this paper."
    },
    "9803/astro-ph9803140_arXiv.txt": {
        "abstract": "The Westerbork Northern Sky Survey (WENSS) has been used to select a sample of Gigahertz Peaked Spectrum (GPS) radio sources at flux densities one to two orders of magnitude lower than bright GPS sources investigated in earlier studies. Sources with inverted spectra at frequencies above $325$ MHz have been observed with the WSRT\\footnote{The Westerbork Synthesis Radio telescope (WSRT) is operated by the Netherlands Foundation for Research in Astronomy with financial support from the Netherlands Organisation for Scientific Research (NWO).} at 1.4 and 5 GHz and with the VLA\\footnote{The Very Large Array (VLA) is operated by the U.S. National Radio Astronomy Observatory which is operated by the Associated Universities, Inc. under cooperative agreement with the National Science Foundation.} at 8.6 and 15 GHz to select genuine GPS sources. This has resulted in a sample of 47 GPS sources with peak frequencies ranging from $\\sim$500 MHz to $>$15 GHz, and peak flux densities ranging from $\\sim$40 to $\\sim$900 mJy. Counts of GPS sources in our sample as a function of flux density have been compared with counts of large scale sources from WENSS scaled to 2 GHz, the typical peak frequency of our GPS sources. The counts can be made similar if the number of large scale sources at 2 GHz is divided by 250, and their flux densities increase by a factor of 10. On the scenario that all GPS sources evolve into large scale radio sources, these results show that the lifetime of a typical GPS source is $\\sim 250$ times shorter than a typical large scale radio source, and that the source luminosity must decrease by a factor of $\\sim 10$ in evolving from GPS to large scale radio source. However, we note that the redshift distributions of GPS and large scale radio sources are different and that this hampers a direct and straightforward interpretation of the source counts. Further modeling of radio source evolution combined with cosmological evolution of the radio luminosity function for large sources is required. ",
        "introduction": "Gigahertz Peaked Spectrum (GPS) radio sources are a class of extragalactic radio source characterized by a spectral peak near 1 Gigahertz in frequency (e.g. Spoelstra et al. 1985) The spectral peak in these compact luminous objects is believed to be due to synchrotron self absorption caused by the high density of the synchrotron emitting electrons in the radio source. GPS sources are interesting objects, both as Active Galactic Nuclei (AGN) and as cosmological probes. It has been suggested that they are young radio sources ($<10^4$ yr) which evolve into large radio sources (Fanti et al. 1995, Readhead et al. 1996, O'Dea and Baum 1997), and studying them would then provide us with important information on the early stages of radio source evolution. Alternatively GPS sources may be compact because a particularly dense environment prevents them from growing larger (e.g. O'Dea et al. 1991).  Important information about the nature of GPS radio sources comes from the properties of their optical counterparts. The galaxies appear to be a homogeneous class of giant ellipticals with old stellar populations (Snellen et al. 1996a, 1996b, O'Dea et al. 1996) and are thus useful probes of galaxy evolution with little or no contamination from the active nucleus in the optical. GPS quasars have a different redshift distribution to their galaxy counterparts ($2<z<4$, O'Dea et al 1991) and their radio morphologies are also quite unlike those of GPS galaxies. The relationship between GPS quasars and galaxies, if any, remains uncertain (Stanghellini et al. 1996).  Previous work on GPS sources has concentrated on the radio-bright members of the class, with $S_{5GHz}>1$ Jy (Fanti et al. 1990, O'Dea et al. 1991, Stanghellini et al. 1996, de Vries et al. 1997). We are carrying out investigations of GPS sources at fainter flux density levels, in order to compare their properties with their radio bright counterparts. This enables us to investigate the properties of GPS sources as a function of radio luminosity, redshift, and rest frame peak frequency. The selection of a sample at intermediate flux densities was described in Snellen et al. (1995a). This paper describes and discusses the selection of an even fainter sample from the Westerbork Northern Sky Survey (WENSS, Rengelink et al. 1997).   \\section {Selection of GPS Sources} \\subsection{The Westerbork Northern Sky Survey}  The Westerbork Northern Sky Survey (WENSS) is being carried out at 325 and 609 MHz (92 and 49 cm) with the Westerbork Synthesis Radio Telescope (WSRT). At 325 MHz, WENSS covers the complete sky north of $30^\\circ$ to a limiting flux density of approximately 18 mJy ($5 \\sigma$). At 609 MHz, about a quarter of this area, concentrated at high galactic latitudes, has been surveyed to a limiting flux density of approximately 15 mJy ($5 \\sigma$). The systematic errors in flux density in WENSS were found to be $\\sim 5\\%$ (Rengelink et al. 1997). The survey was conducted in mosaicing mode which is very efficient in terms of observing time. In this mode, the telescope cycles through 80 evenly spaced field centres, during each of a number of $12^h$ syntheses with different spacings of array elements. The visibilities are sufficiently well sampled for all 80 fields that it is possible to reconstruct the brightness distribution in an area of the sky, $\\sim$100 square degrees, which is many times larger than the primary beam of the WSRT. Individual fields are referred to as {\\it mosaics}, and have a resolution (FWHM of the restoring beam) of $54'' \\times 54''$ cosec $\\delta$ at 325 MHz and $28''\\times 28''$ cosec $\\delta$ at 609 MHz. From the combined mosaics, maps are made with a uniform sensitivity and regular size, which are called {\\it frames}. The 325 MHz frames are $6^\\circ \\times 6^\\circ $ in size and positioned on a regular $5^\\circ \\times 5^\\circ $ grid over the sky, which coincides with the position grid of the new Palomar Observatory Sky Survey (POSS II, Reid et al. 1991) plates. A detailed description of WENSS is given by Rengelink et al. (1997)  \\subsection{Selection of a Sample of Candidate GPS Sources.}  A deep low frequency radio survey such as WENSS is crucial for selecting a sample of faint GPS sources. It is the inverted spectrum at low frequencies which distinguishes them from other types of radio sources. Figure \\ref{surveys} shows the major large-sky radio surveys in the northern sky with theoretical spectra of homogeneous synchrotron self absorbed radio sources (eg. Moffet 1975) superimposed, which have spectral peak frequencies of 1 GHz. Samples of GPS sources can be constructed using WENSS flux density measurements in the optically thick part of their spectra which are ten times fainter than samples selected using the Texas Survey (Douglas 1996).  \\begin{figure}[!t] \\centerline{ \\psfig{figure=figure1.ps,width=8cm}} \\caption{\\label{surveys} Overview of the major radio surveys in the northern sky: the Greenbank Surveys (Condon and Broderick 1985, Gregory and Condon 1991), the Texas Survey (Douglas et al. 1996), and the Cambridge 3C, 4C, and 6C surveys. The curves represent the spectra of a homogeneous synchrotron self absorbed radio source, with a peak frequency of 1 GHz and peak flux density of 300 mJy (lower curve) and 3000 mJy (upper curve). Samples of GPS sources can be constructed using WENSS flux density measurements in the optically thick part of their spectra which are more than an order of magnitude fainter than samples selected using the Texas Survey.} \\end{figure}  When we selected our sample, only a small part of the WENSS region had been observed and the data reduced to the point of providing source lists. The 325 MHz WENSS data used to select GPS sources are from two regions of the sky; one at $15^{\\rm h} < \\alpha < 20^{\\rm h}$ and $58^\\circ< \\delta < 75^\\circ$, which is called the {\\it mini-survey} region (Rengelink et al. 1997), and the other at $4^{\\rm h}00^{\\rm m} < \\alpha < 8^{\\rm h}30^{\\rm m}$ and $58^\\circ< \\delta < 75^\\circ$, where $\\alpha$ is right ascension and $\\delta$ is declination. These were the first two regions observed, reduced and analysed for WENSS. The mini-survey region, which is roughly centered on the North Ecliptic Pole, was chosen as the first area for analysis because it coincides  with the NEP-VLA survey at 1.5 GHz (Kollgaard et al. 1994), the deep 7C North Ecliptic Cap survey (Lacy et al. 1995, Visser et al. 1995), the deepest part of the ROSAT All Sky survey (Bower et al., 1996) and the IRAS survey (Hacking and Houck, 1987). The high declination of the two regions is very convenient for VLBI experiments, because their locations are circumpolar for almost all the major EVN and VLBA radio telescopes. At the time of selection WENSS 609 MHz data was available for only about one third of the mini-survey region. The regions for which both 325 and 609 MHz source lists were available cover 119 square degrees of the sky. The regions for which only 325 MHz data were available cover 216 square degrees in the mini-survey region and 306 square degrees in the other region.  These source lists were correlated with those from the Greenbank 5 GHz (6 cm) survey (Gregory and Condon 1991, Gregory et al. 1996), which has a limiting flux density of 25 mJy ($5 \\sigma$). For the faintest sources the new Greenbank source list (Gregory et al. 1996) was used, which is based on more data. Candidate GPS sources were selected on the basis of a positive spectral index $\\alpha$ between 325 MHz and 5 GHz, where the spectral index is defined by $S \\sim \\nu^{\\alpha}$. If 609 MHz data was also available, an ``inverted'' spectrum between 325 MHz and 609 MHz was used as the selection criterion. This in fact increased the sensitivity of the selection process to GPS sources with low peak frequencies ($< 1 $ GHz). Note that in general for a GPS source, the 325-609 MHz spectral index will be more positive than the 325-5000 MHz spectral index for a spectral peak in the 1 GHz range. Hence, using the 325-609 MHz selection criterion will not miss any GPS sources which would have been found using the 325-5000 MHz selection criterion, it will only add extra sources with lower peak frequencies.  In total, 117 inverted spectrum sources were selected; 37 using the 325-609 MHz selection and 82 using the 325-5000 MHz selection. They are listed in table \\ref{canGPS}. Columns 1, and 2 give the name, right ascension and declination (B1950) (obtained from the VLA observations), columns 3, 4 and 5 the 325 MHz, 609 MHz and 5 GHz flux densities, and columns 6 and 7 give the 325-609 MHz and 325-5000 MHz spectral indices. The uncertainties in the 325-5000 MHz spectral indices range from 0.03 to 0.05 (for the faintest objects), and the uncertainties in the 325-609 MHz spectral index range from 0.10 to 0.40.  \\subsection{Additional Radio Observations.}  An apparently inverted or peaked spectrum could be caused by variability at any or all of the selection frequencies, due to the fact that the 325, 609 and 5000 MHz surveys were observed at different epochs. To select the genuine GPS sources, additional quasi-simultaneous observations at other frequencies are required to eliminate flat spectrum, variable radio sources. The 5 GHz Greenbank survey was made in 1987, while the 325 MHz and 609 MHz  data were taken in 1993. Furthermore, high frequency data is needed to confirm their turnover, and measure the (steep) spectrum in the optically thin part of their spectra. Therefore VLA observations were taken at 8.4 and 15 GHz, and WSRT observations at 1.4 and 5 GHz. Later, after the selection process, data at 1.4 GHz from the NRAO VLA Sky Survey (NVSS, Condon et al. 1996) became available and were used to supplement our spectra.  \\subsubsection{WSRT Observations at 1.4 and 5 GHz}  The WSRT was used to observe the candidate GPS sources at 1.4 and 5 GHz. The 1.4 GHz observations were performed on 20 February and 10 March 1994 using 8 bands of 5 MHz between 1377.5 and 1423.5 MHz, providing a total bandwidth of 40 MHz. The sources were all observed for about 100 seconds at two to three different hour angles. This resulted in a noise level of typically 1 mJy/beam and a resolution of $15''\\times 15''cosec \\ \\delta$. The results are shown in column 8 of table \\ref{canGPS}.  In order to improve the 5 GHz Greenbank flux density measurements, observations were carried out with the WSRT  at 4.87 GHz on May 15 1994 using a bandwidth of 80 MHz, at a time when the WSRT was participating a VLBI session. Unfortunately only three telescopes were equipped with 5 GHz receivers.  Only sources between 4 and 8 hours right ascension were observed, and the uncertainty in the measured flux densities is large ($\\sim 15$\\%). The resulting flux densities are listed in column 13 in table \\ref{canGPS}.   \\begin{table*} \\renewcommand{\\arraystretch}{0.90} \\setlength{\\tabcolsep}{1mm} \\begin{tabular}{|c|rrrrrr|rrr|rr|rrr|c|rrr|} \\hline Source&\\multicolumn{3}{c}{R.A.(1950)}& \\multicolumn{3}{c|}{Decl.(1950)}&$S_{325}$&$S_{609}$&$S^{gb}_{5.0}$&$\\alpha ^{325}_{609}$&$\\alpha ^{325}_{5000}$& $S^{wsrt}_{1.4}$&$S^{VLA}_{8.6}$&$S^{VLA}_{14.9}$&GPS&$S^{nvss}_{1.4}$& $S^{wsrt}_{5.0}$&$S^{merlin}_{5.0}$\\\\ & h &  m &   s  &$^{\\circ}$&$'$&$''$&{\\tiny (mJy)}&{\\tiny(mJy)}& {\\tiny(mJy)}&&    &{\\tiny(mJy)}&{\\tiny(mJy)}&{\\tiny(mJy)}& &{\\tiny(mJy)}&{\\tiny(mJy)}&{\\tiny(mJy)}\\\\ \\hline B0400+6042 & 4 &  0 &  7.22 & 60 & 42 & 29.0   & 81 &  & 100&    &+0.08& 180  & 85& 36 &+& 166& 73& 85           \\\\ B0402+6442 & 4 &  2 & 56.73 & 64 & 42 & 52.4   & 48 &  &  69&    &+0.13&  63  & 56& 45 & &  70& 32&              \\\\ B0406+7413 & 4 &  6 & 37.39 & 74 & 13 & 22.5   & 63 &  &  66&    &+0.02&  56  & 54& 36 & &  60& 49&              \\\\ B0418+6724 & 4 & 18 & 50.86 & 67 & 24 &  7.2   & 47 &  &  95&    &+0.26&  54  & 96& 79 & &  50& 46&              \\\\ B0436+6152 & 4 & 36 & 15.80 & 61 & 52 & 10.0   & 70 &  & 127&    &+0.22& 208  &108& 60 &+& 238&101&122           \\\\ B0441+5757 & 4 & 41 & 53.90 & 57 & 57 & 21.7   & 60 &  & 115&    &+0.24&  95  &128& 96 &+&  91& 91&101           \\\\ B0456+7124 & 4 & 56 &  0.15 & 71 & 24 & 10.1   & 88 &  & 148&    &+0.19& 121  &213&216 & & 116&181&              \\\\ B0507+6840 & 5 & 7 &   4.48 & 68 & 40 & 44.7   & 27 &  &  33&    &+0.07&  30  & 41& 29 & &  31& 28&              \\\\ B0513+7129 & 5 & 13 & 38.82 & 71 & 29 & 55.1   &121 &  & 181&    &+0.15& 236  & 96& 60 &+& 244&123&131           \\\\ B0515+6129 & 5 & 15 & 19.62 & 61 & 28 & 59.3   & 40 &  &  45&    &+0.04&  26  & 62& 41 & &  28& 43&              \\\\ B0518+6004 & 5 & 18 & 42.78 & 60 &  4 & 55.9   & 88 &  & 113&    &+0.09&  67  & 92&102 & & 101& 56&              \\\\ B0531+6121 & 5 & 31 & 55.43 & 61 & 21 & 31.0   & 18 &  &  38&    &+0.27&  19  & 44& 23 &+&  22& 33& 23           \\\\ B0535+6743 & 5 & 35 & 57.15 & 67 & 43 & 49.8   & 83 &  & 182&    &+0.29&  97  &235&136 &+&148&177 &168          \\\\ B0536+5822 & 5 & 36 &  8.26 & 58 & 22 &  3.8   & 26 &  &  27&    &+0.01&  36  & 54& 39 & & 31& 37 &             \\\\ B0537+6444 & 5 & 37 & 15.12 & 64 & 45 &  3.7   & 18 &  &  32&    &+0.21&  29  & 17&  9 &+& 28& 17 & 17          \\\\ B0538+7131 & 5 & 38 & 38.30 & 71 & 31 & 20.9   & 19 &  &  77&    &+0.51&  45  & 68& 29 &+& 48& 90 & 73          \\\\ B0539+6200 & 5 & 39 & 54.51 & 62 &  0 &  2.2   & 47 &  & 104&    &+0.29& 123  & 80& 66 &+&126&112 & 99          \\\\ B0542+7358 & 5 & 42 & 50.98?& 73 & 58 & 32.5   & 29 &  &  29&    &+0.00&  46  & 33& 24 & & 44& 34 &             \\\\ B0543+6523 & 5 & 43 & 40.36 & 65 & 23 & 24.6   & 26 &  &  43&    &+0.18&  65  & 47& 27 &+& 72& 49 & 43          \\\\ B0544+5847 & 5 & 44 &  3.18 & 58 & 46 & 55.8   & 33 &  &  42&    &+0.09&  67  & 43& 22 &+& 60& 48 & 34          \\\\ B0552+6017 & 5 & 52 & 35.07 & 60 & 17 & 30.1   & 15 &  &  26&    &+0.28&  44  & 11&$<3$&+& 47& 11 & 13          \\\\ B0556+6622 & 5 & 56 & 13.00 & 66 & 22 & 57.7   & 22 &  &  25&    &+0.05&  12  & 30& 20 & & 24& 27 &             \\\\ B0557+5717 & 5 & 57 & 31.76 & 57 & 17 & 19.7   & 19 &  &  29&    &+0.16&  63  & 28& 14 &+& 69& 36 & 30          \\\\ B0601+7242 & 6 &  1 & 57.83 & 72 & 42 & 54.6   & 16 &  &  26&    &+0.18&  14  & 18&  8 & & 14& 37 &             \\\\ B0601+5753 & 6 &  1 & 22.05 & 57 & 53 & 31.8   & 19 &  & 162&    &+0.83& 141  &149&138 &+&   &207 &192          \\\\ B0605+7218 & 6 &  5 &  6.15 & 72 & 18 & 51.1   & 24 &  & 116&    &+0.51&  80  & 39& 60 & & 60& 65 &             \\\\ B0607+7335 & 6 &  7 & 33.86 & 73 & 35 & 53.0   & 20 &  &  54&    &+0.36&  73  & 86& 54 & & 53& 87 &             \\\\ B0607+7107 & 6 &  7 & 54.42 & 71 &  8 & 14.1   & 21 &  &  25&    &+0.10&  24  & 21& 18 & & 27& 42 &             \\\\ B0609+7259 & 6 &  9 & 14.93 & 72 & 59 & 50.7   & 20 &  &  25&    &+0.08&  16  & 16& 11 & & 22& 23 &             \\\\ B0738+7043 & 7 & 38 & 37.86 & 70 & 43 &  9.2   & 47 &  &  84&    &+0.19&  34  & 18&106 & & 73&100 &             \\\\ B0741+7213 & 7 & 41 & 30.37 & 72 & 12 & 58.5   & 65 &  & 106&    &+0.18&  96  & 54& 55 & & 99& 75 &             \\\\ B0748+6343 & 7 & 48 & 27.42 & 63 & 43 & 31.7   & 24 &  &  54&    &+0.42&  39  & 69& 85 &+& 44& 87 &130          \\\\ B0752+6355 & 7 & 52 & 21.41 & 63 & 55 & 59.5   & 15 &  & 254&    &+1.04& 133  &298&282 &+&196&376 &303          \\\\ B0755+6354 & 7 & 55 & 20.66 & 63 & 54 & 25.4   & 25 &  &  38&    &+0.15&  26  & 17& 18 &+& 26& 21 & 19         \\\\ B0756+6647 & 7 & 56 & 47.89 & 66 & 47 & 27.8   & 44 &  & 164&    &+0.48&  90  & 80& 77 &+&104&109 & 97          \\\\ B0758+5929 & 7 & 58 & 13.00 & 59 & 29 & 56.9   & 97 &  & 178&    &+0.22& 203  &127&101 &+&207&185 &163          \\\\ B0759+6557 & 7 & 59 & 12.95 & 65 & 57 & 45.9   & 15 &  &  27&    &+0.22&  40  & 14&  8 &+& 48& 26 & 21          \\\\ B0800+6754 & 8 &  0 &  6.92 & 67 & 54 & 31.9   & 78 &  &  78&    &+0.00&  30  & 33& 32 & & 39& 38 &             \\\\ B0802+7323 & 8 &  2 & 32.27 & 73 & 23 & 53.3   &318 &  & 321&    &+0.00& 277  &387&445 & &307&437 &             \\\\ B0808+6518 & 8 &  8 &  5.14 & 65 & 18 & 10.8   & 35 &  &  42&    &+0.07&  52  & 17& 21 & & 31& 32 &             \\\\ B0810+6440 & 8 & 10 &  7.59 & 64 & 40 & 29.1   & 96 &  & 193&    &+0.26&  92  &132&179 & & 90&133 &             \\\\ B0820+7403 & 8 & 20 & 43.56 & 74 &  2 & 53.8   &104 &  & 108&    &+0.01&  81  & 65& 63 & &102& 97 &             \\\\ B0824+6446 & 8 & 24 & 31.62 & 64 & 46 & 27.4   & 24 &  &  35&    &+0.14&  30  & 11& 13 & & 44& 29 &             \\\\ B0826+7045 & 8 & 26 & 52.55 & 70 & 45 & 44.1   & 34 &  & 109&    &+0.43&  73  & 63& 56 &+& 79& 99 & 92            \\\\ B0827+6231 & 8 & 27 &  3.37 & 62 & 31 & 45.9   & 28 &  &  28&    &+0.00&  32  & 25& 33 & & 34& 32 &             \\\\ B0828+5756 & 8 & 28 & 33.49 & 57 & 56 &  5.6   & 29 &  &  32&    &+0.04&  37  & 27& 26 & & 53& 40 &             \\\\ B0828+7307 & 8 & 28 & 49.08 & 73 &  6 & 58.7   & 81 &  & 104&    &+0.09&  68  & 58& 68 & &102& 86 &             \\\\ B0830+5813 & 8 & 30 & 12.71 & 58 & 13 & 38.8   & 39 &  &  65&    &+0.29&  65  & 31& 23 &+& 59& 43 & 38             \\\\ B0830+6300 & 8 & 30 & 37.77 & 63 &  0 &  8.1   & 26 &  &  36&    &+0.12&  54  & 50& 54 & & 62& 65 &             \\\\ B0830+6845 & 8 & 30 & 59.69 & 68 & 45 & 33.1   & 53 &  &  74&    &+0.12&  83  &136&124 & & 86&159 &             \\\\ B1525+6801 &15 & 25 & 21.12 & 68 &  1 & 48.9   & 90 &  & 103&    &+0.05& 153  & 54& 29 &+&161&    & 91          \\\\ B1529+6741 &15 & 29 & 17.85 & 67 & 41 & 58.6   & 47 &  &  55&    &+0.06&  19  & 32& 26 & & 35&    &             \\\\ B1529+6829 &15 & 29 & 45.15 & 68 & 29 &  9.3   & 51 &  &  51&    &+0.00&  29  & 21& 23 & & 27&    &             \\\\ B1536+6202 &15 & 36 & 54.56 & 62 &  2 & 56.4   & 16 &  &  30&    &+0.23&  15  & 34& 24 & & 22&    &             \\\\ B1538+5920 &15 & 38 & 27.46 & 59 & 20 & 39.0   & 29 &  &  73&    &+0.34&  47  & 36& 23 &+& 45&    & 45          \\\\ B1539+6156 &15 & 39 & 32.28 & 61 & 56 &  1.9   & 34 &  &  40&    &+0.06&  15  & 54& 36 & & 33&    &             \\\\ B1542+6139 &15 & 42 &  5.03 & 61 & 39 & 20.9   & 60 &  & 129&    &+0.28&  86  &114&119 & & 90&    &             \\\\ B1542+6631 &15 & 42 & 54.19 & 66 & 31 & 18.2   & 54 &  &  86&    &+0.17&  44  & 81& 84 & & 51&    &             \\\\ B1550+5815 &15 & 50 & 55.59 & 58 & 15 & 37.5   & 86 &  & 362&    &+0.53& 157  &214&212 &+&   &    &237          \\\\ B1551+6822 &15 & 51 & 53.07 & 68 & 22 & 38.7   & 23 &  &  34&    &+0.14&  49  & 26& 10 &+& 55&    & 27          \\\\\\hline \\end{tabular} \\caption{\\label{canGPS} The sample of candidate GPS sources. Column 1 gives the B1950 source name, column 2 the VLA position, column 3, 4 and 5 the flux densities from WENSS at 325 and 609 MHz and of the Greenbank Survey at 5 GHz. Column 6 and 7 give the 325-609 MHz and the 325-5000 MHz spectral indices, column 8, 9 and 10 the flux densities from the WSRT at 1.4 GHz, and from the VLA at 8.4 and 15 GHz. A cross in column 11 indicates whether the source was selected in the final sample. Column 12, 13 and 14 give the NVSS 1.4 GHz, the WSRT 5 GHz and the MERLIN 5 GHz flux densities.} \\end{table*}  \\addtocounter{table}{-1}  \\begin{table*} \\renewcommand{\\arraystretch}{0.90} \\setlength{\\tabcolsep}{1mm} \\begin{tabular}{|c|rrrrrr|rrr|rr|rrr|c|rrr|} \\hline Source&\\multicolumn{3}{c}{R.A.(1950)}& \\multicolumn{3}{c|}{Decl.(1950)}&$S_{325}$&$S_{609}$&$S^{gb}_{5.0}$&$\\alpha ^{325}_{609}$&$\\alpha ^{325}_{5000}$& $S^{wsrt}_{1.4}$&$S^{VLA}_{8.6}$&$S^{VLA}_{14.9}$&GPS&$S^{nvss}_{1.4}$& $S^{wsrt}_{5.0}$&$S^{merlin}_{5.0}$\\\\ & h &  m &   s  &$^{\\circ}$&$'$&$''$&{\\tiny (mJy)}&{\\tiny(mJy)}& {\\tiny(mJy)}&&    &{\\tiny(mJy)}&{\\tiny(mJy)}&{\\tiny(mJy)}& &{\\tiny(mJy)}&{\\tiny(mJy)}&{\\tiny(mJy)}\\\\ \\hline B1557+6220 &15 & 57 &  8.43 & 62 & 20 &  7.4   & 23 &   &   37 &     &+0.17& 40   & 12&  4 &+& 42& & 18           \\\\ B1559+5715 &15 & 59 &  5.07 & 57 & 15 & 19.2   & 27 &   &   47 &     &+0.20& 55   & 36& 39 & & 68& &               \\\\ B1600+5714 &16 &  0 &  8.62 & 57 & 14 & 18.8   & 29 &   &   45 &     &+0.16& 24   & 12& 76 & & 41& &               \\\\ B1600+7131 &16 &  0 & 57.00 & 71 & 31 & 40.2   & 26 &   &  103 &     &+0.50&311   & 37& 15 &+&308& & 85            \\\\ B1604+5939 &16 &  4 & 56.32 & 59 & 39 & 43.7   & 71 &   &  110 &     &+0.16&111   &130&112 & &   & &               \\\\ B1607+6026 &16 &  7 & 30.32 & 60 & 26 & 34.6   & 79 &   &   79 &     &+0.00& 16   & 31& 22 & & 34& &               \\\\ B1608+6540 &16 &  8 & 50.55 & 65 & 40 & 15.1   & 66 &   &   90 &     &+0.11& 63   & 40& 78 & & 68& &               \\\\ B1616+6428 &16 & 16 & 26.42 & 64 & 28 &  7.7   & 19 & 37&   62 &+1.06&+0.43& 64   & 58& 58 & & 61& &               \\\\ B1620+6406 &16 & 20 & 46.19 & 64 &  6 & 12.6   & 27 & 30&   25 &+0.16&$-$0.03& 41   &  6&$<3$&+& 43& & 11            \\\\ B1622+6630 &16 & 22 & 50.52 & 66 & 30 & 52.6   & 23 & 61&  517 &+1.55&+1.14&178   &230&176 &+&159& &230            \\\\ B1623+6859 &16 & 23 & 36.01 & 68 & 59 & 46.0   & 32 & 32& $<25$&+0.00&     & 11   &  3&$<3$& & 18& &               \\\\ B1624+6622 &16 & 24 &  7.26 & 66 & 22 &  6.7   & 38 & 42& $<25$&+0.16&     & 33   &  4&$<3$& & 26& &               \\\\ B1633+6506 &16 & 33 &  7.49 & 65 &  6 & 52.4   & 48 & 66&  114 &+0.51&+0.30&111   & 97& 90 & & 95& &               \\\\ B1639+6711 &16 & 39 & 10.76 & 67 & 11 & 47.2   & 34 & 61&   30 &+0.93&$-$0.05& 54   & 27& 19 &+& 76& & 40            \\\\ B1642+6701 &16 & 42 & 16.42 & 67 &  1 & 22.6   &124 & 126&  65 &+0.02&$-$0.24&121   & 43& 24 &+&126& & 56            \\\\ B1645+6738 &16 & 45 & 38.00 & 67 & 38 &  0.8   & 22 & 22& $<25$&+0.00&     & 28   &  8&  5 & & 28& &               \\\\ B1647+6225 &16 & 47 & 31.26 & 62 & 25 & 49.7   & 31 & 59&   33 &+1.02&+0.02& 69   & 13&  3 &+& 56& & 17            \\\\ B1655+6446 &16 & 55 & 21.09 & 64 & 46 & 21.2   & 23 & 52&   34 &+1.30&+0.14& 68   & 16&  9 &+& 61& & 23            \\\\ B1657+5826 &16 & 57 & 15.96 & 58 & 26 & 31.5   & 56 & 63&   17 &+0.19&$-$0.44& 44   & 18& 13 &+& 48& & 23            \\\\ B1711+6031 &17 & 11 & 39.99 & 60 & 31 & 45.3   & 15 & 29& $<25$&+1.05&     & 17   &  3&$<3$& & 19& &               \\\\ B1712+6727 &17 & 12 & 50.73 & 67 & 27 & 10.2   & 28 & 30& $<25$&+0.11&     & 28   & 23& 13 & & 31& &               \\\\ B1714+5819 &17 & 14 & 56.27 & 58 & 19 & 16.2   & 21 & 35& $<25$&+0.81&     & 25   &  6&  3 & & 28& &               \\\\ B1718+6024 &17 & 18 & 18.75 & 60 & 24 & 11.7   & 19 & 29& $<25$&+0.67&     & 20   &  3&$<3$& & 18& &               \\\\ B1730+6027 &17 & 30 & 15.71 & 60 & 27 & 24.9   & 41 & 83&   34 &+1.12&$-$0.13&178   &103& 89 & &161& &               \\\\ B1746+6921 &17 & 46 & 53.21 & 69 & 21 & 33.5   & 65 & 96&  144 &+0.62&+0.31&161   &127&100 &+&154& &139            \\\\ B1749+6919 &17 & 49 & 31.62 & 69 & 19 & 41.3   & 21 & 31&   18 &+0.62&$-$0.06& 27   & 10&  7 & & 32& &               \\\\ B1755+6905 &17 & 55 & 42.48 & 69 &  5 & 48.4   & 16 & 72&   77 &+2.40&+0.28& 84   & 51& 48 & & 78& &               \\\\ B1807+5959 &18 &  7 & 17.36 & 59 & 59 & 26.5   & 16 & 43&   30 &+1.52&+0.28& 47   & 37& 22 &+& 42& & 38            \\\\ B1807+6742 &18 &  7 & 23.43 & 67 & 42 & 22.3   & 29 & 52&$<25$ &+0.93&     & 47   & 12&  8 &+& 43& & 20            \\\\ B1808+6813 &18 &  8 & 25.41 & 68 & 13 & 36.1   & 33 & 37&   24 &+0.18&+0.04& 42   & 11&  8 &+& 42& & 19              \\\\ B1818+6445 &18 & 18 & 24.79 & 64 & 45 & 17.1   &  35 & 38 &$<25$&+0.13     &  &  24 & 27 &$<3$& & 27& &               \\\\ B1818+6249 &18 & 18 & 50.02 & 62 & 49 & 56.2 &  41 & 50 &$<25$&+0.32     &  &  34 &  8 &  6&  & 33& &              \\\\ B1819+6707 &18 & 19 & 48.42 & 67 &  7 & 20.8 & 265 &330 & 154 &+0.35   &$-0.20$  & 297 & 93 & 68&+ &311& &142           \\\\ B1821+6251 &18 & 21 & 20.03 & 62 & 51 & 52.6 &  30 & 38 &  31 &+0.38   &+0.01 &  32 & 28 & 19&  & 29& &              \\\\ B1827+6432 &18 & 27 & 55.40 & 64 & 32 & 13.7 & 152 &228 & 262 &+0.65   &+0.20 & 204 &135 &120&  &216& &              \\\\ B1829+6419 &18 & 29 & 16.62 & 64 & 19 & 23.1 &  62 & 95 &$<25$&+0.68     &  &  80 &  6 &  4&  & 75& &              \\\\ B1834+6319 &18 & 34 & 48.26 & 63 & 19 & 49.6 &  37 & 46 &$<25$&+0.35     &  &  36 &  5 &$<3$& & 36& &               \\\\ B1838+6239 &18 & 38 & 12.00 & 62 & 39 & 56.2 &  15 & 33 &$<25$&+1.25     &  &  54 &  7 & 6&   & 36& &             \\\\ B1841+6715 &18 & 41 &  7.21 & 67 & 15 & 51.2 &  36 & 94 & 163 &+1.53   &+0.55  & 142 & 98 &68& + &178& &125          \\\\ B1841+6343 &18 & 41 & 18.25 & 63 & 43 & 56.3 &  15 & 29 &$<25$&+1.05  &    &  41 & 10 &  6  &   & 36& &           \\\\ B1843+6305 &18 & 43 &  6.16 & 63 &  5 & 42.8 &  15 & 41 &  52 &+1.67   &+0.45  &  59 & 27 & 16  & + & 81& & 40        \\\\ B1850+6447 &18 & 59 & 27,78 & 64 & 47 & 31.6 &  49 & 70 &$<25$&+0.57  &  &  52 &  6 & $<3$&   & 55& &             \\\\ B1916+6817 &19 & 16 & 37.66 & 68 & 17 & 51.6 &  23 & 39 &$<25$&+0.84  &  &  20 &  6 &  4  &   & 26& &           \\\\ B1919+6912 &19 & 19 & 57.99 & 69 & 12 & 26.2 &  21 & 28 &  18 &+0.46  &$-0.06$  &  16 & 16 & 14 15&   & & &           \\\\ B1926+6111 &19 & 26 & 49.66 & 61 & 11 & 20.9 & 404 &    & 613 &  &+0.15 & 718 & 85 &678  &   & &535&           \\\\ B1934+7111 &19 & 34 & 41.81 & 71 & 11 & 10.4 &  92 &    & 108 &  &+0.06 & 142 &119 &102  &   &179 & &           \\\\ B1938+5824 &19 & 38 & 50.57 & 58 & 24 & 49.9 &  24 &    &  29 &  &+0.07 &  34 & 30 & 25  &   &21 & &           \\\\ B1942+7214 &19 & 42 &  2.24 & 72 & 14 & 31.9 &  81 &    & 158 &  &+0.24 & 233 &147 &110  & + &233 & &183        \\\\ B1944+6007 &19 & 44 & 21.42 & 60 &  7 & 40.5 &  20 &    &  79 &  &+0.50 &  12 & 62 & 32  &   &17 & &           \\\\ B1945+6024 &19 & 45 & 24.83 & 60 & 24 & 12.6 &  25 &    &  80 &  &+0.43 &  55 &125 &188  & + & 55& & 84        \\\\ B1946+7048 &19 & 46 & 12.02 & 70 & 48 & 21.6 & 234 &    & 643 &  &+0.37 & 887 &389 &268  & + & 953& &574        \\\\ B1951+6915 &19 & 51 & 34.02 & 69 & 15 &  6.3 &  23 &    &  32 &  &+0.12 &  24 & 40 & 35  &   & 33& &           \\\\ B1951+6453 &19 & 51 & 42.52 & 64 & 53 & 56.1 & 111 &    & 103 &  &+0.00 &  89 & 83 & 59  &   & 88& &           \\\\ B1954+6146 &19 & 54 & 11.69 & 61 & 45 & 58.1 &  66 &    & 132 &  &+0.25 &  66 &182 &152  & + & 61& &153        \\\\ B1958+6158 &19 & 58 & 45.58 & 61 & 58 & 27.1 &  52 &    & 140 &  &+0.36 & 111 & 96 & 84  & + & 129& &136        \\\\ B2006+5916 &20 & 06 & 52.14 & 59 & 16 & 43.5 &  30 &    &  37 &  &+0.08 &  36 & 21 & 18  &   & 30& &           \\\\ B2011+7156 &20 & 11 & 22.95 & 71 & 56 &  9.4 &  48 &    & 134 &  &+0.38 & 118 &103 &102  &   & & &           \\\\\\hline \\end{tabular} \\caption{{\\it Continued...}} \\end{table*}   \\subsubsection{VLA Observations at 8.4 and 15 GHz}  The candidate GPS sources were observed with the VLA in B-configuration at 8.4 and 15 GHz on 23 July 1994. At both frequencies, the objects were observed in a standard way using a bandwidth of $2\\times 25$ MHz. The phases were calibrated using standard nearby VLA phase calibrators. Total integration times were typically 100 seconds at both frequencies, resulting in noise levels of 0.2 and 1.0 mJy/beam respectively. Systematic errors in flux density of VLA observations at these frequencies are typically about $3\\%$ (eg. Carilli et al. 1991). The data were reduced using AIPS in a standard manner, including several iterations of phase self-calibration. The synthesized beams have half widths of $1.5''\\times 0.8''$ and $0.8''\\times 0.5 ''$ at 8.4 and 15 GHz respectively. Several candidate GPS sources had already been observed at 8.4 GHz on February 26 1994 and April 3 1994 during the Cosmic Lens All Sky Survey (CLASS) program (eg. Myers et al. 1995); these were not re-observed by us at 8.4 GHz. The CLASS 8.4 GHz observations were made using the VLA in A configuration in a standard way, also with a bandwidth  of $2\\times 25$ MHz and an average integration time of 30 seconds. The resolution of the CLASS observations was $\\sim 0.2''$, and the noise level $\\sim 0.4$ mJy/beam.  The results of the VLA observations are listed in columns 9 and 10 in table \\ref{canGPS}. All of the sources were unresolved, except for B1608+6540, which was found to be a quadruple gravitational lens (Snellen et al. 1995b, Myers et al. 1995, Fassnacht et al. 1996)   \\subsubsection{The NRAO VLA Sky Survey at 1.4 GHz}  Observations for the 1.4 GHz NRAO VLA Sky Survey (NVSS, Condon et al. 1996) began in September 1993 and are planned to cover the sky north of declination $-40^{\\circ}$ (82\\% of the celestial sphere). Data in our regions of interest were  taken on 1 November 1993 for the region $4^h00^m<R.A.<8^h00^m$, and on 2 April 1995 for the region between $15^h00^m<R.A.<20^h00^m$. The noise level in an image is typically 0.5 mJy/beam and the resolution is $45''$.  \\subsubsection{MERLIN Observations at 5 GHz}  The final sample of genuine GPS sources, as selected in section 2.4, was observed with MERLIN at 5 GHz on 15 and 16 May 1995 during our global VLBI measurements. The sources were observed in three to four ``snapshots'' of 13 minutes each, resulting in a noise level of typically 0.3 mJy/beam, and a resolution of 0.04$''$. All the sources were unresolved. The results are listed in column 14 of table \\ref{canGPS}. Note that these observations were obtained after, and therefore not used for, the final selection.  \\subsection{Selection of the Genuine GPS Sources}  The genuine GPS sources were selected using the 325 MHz and, if available, the 609 MHz WENSS data, 1.4 GHz WSRT data, 5 GHz Greenbank data and 8.4 and 15 GHz VLA data. The WSRT 5 GHz data were not used for selection because they were only available for a part of the sample. The NVSS data was not used for selection, being not available at the time of source selection. However, both the 5 GHz WSRT and 1.4 GHz NVSS data were used for variability studies (see section \\ref{varsec}).  The selection criteria were as follows: \\begin{itemize} \\item[1.] The spectrum must decrease monotonically below the frequency with the highest flux density, taking into account an assumed uncertainty of 10\\% in flux density. \\item[2.] The spectrum must decrease monotonically above the frequency with the highest flux density,taking into account an assumed uncertainty of 10\\% in flux density. \\item[3.] The Full Width Half Maximum (FWHM) defined by the logarithm of the spectrum must be less than 2 decades in frequency. A spectral index of 0.5 is assumed below 325 MHz, and a spectral index of -0.5 above 15 GHz. \\item[4.] The Greenbank 5 GHz flux density must be greater than 20 mJy. This allowed imaging of the source with global VLBI at 5 GHz without recourse to phase referencing. Note that if no Greenbank flux density was available (noise level is about 5 mJy), the flux density was estimated by interpolating the 1.4 and 8.4 GHz flux density points. \\end{itemize}  The resulting sample of 47 sources is listed in table \\ref{GPS}. One of the sources, B1807+5959, did not obey the criterion of decreasing flux density above the peak frequency, because the 5 GHz Greenbank flux density flux point is too low. However it was kept in the sample because the fall off in flux density at both low and high frequencies suggests that the low flux density point at 5 GHz is due to the different epoch of the Greenbank observations. Additional observations showed this to be true.   The spectra of the selected GPS sources were fitted with the following function \\begin{equation} \\label{eq1} S(\\nu ) = S_{max} / (1-e^{-1}) \\times \\left( \\frac{\\nu }{\\nu _{max}} \\right)^k \\times (1  - e^{- \\left( \\frac{\\nu }{\\nu _{max}}\\right) ^{l - k}}) \\end{equation} where $k$ is the optically thick spectral index, $l$ the optically thin spectral index, and $S_{max}$ and $\\nu _{max}$ respectively the peak flux density and peak frequency. This equation, which represents a homogeneous synchrotron self absorbed radio source for $k = 2.5$ (eg. Moffet 1975), fits the spectral peak well in most cases, however it did not always fit the flux density points at the lowest and highest frequency frequency adequately, and therefore was only used to determine the peak flux density, peak frequency and Full Width Half Maximum of the spectra. The optically thick and thin spectral indices have been determined from the two lowest and the two highest frequency data points respectively. The fitted spectra are shown in figure \\ref{spectra}. Table \\ref{GPS} gives the characteristics of the GPS sources: column 1 gives the source name, columns 2 and 3 the peak frequency and peak flux density, columns 4 and 5 the optically thick and optically thin spectral indices, and column 6 the FWHM of the spectrum in logarithmic units.  Figure \\ref{spectra} shows that three of the 47 sources initially selected probably do not have genuine GPS spectra, namely B0531+6121, B0748+6343 and B0755+6354. In these cases differences between MERLIN, Greenbank and WSRT  5 GHz flux density measurements suggest that the measured spectra are contaminated by flux density variability and it is not clear whether the spectra do indeed exhibit a peak. Although we have included them in the sample, we omit them from the analysis below. Note that no sign of a turnover is seen in B1945+6024, and that there are some sources in the sample which do not have a ``clean'' peaked spectrum, like B1954+6146 and B0535+6743. To obtain a better determination of the spectral peak of B1954+6146, the 325 MHz flux density data point is not used to fit the spectrum.   \\begin{table} \\begin{center} \\setlength{\\tabcolsep}{1.05mm} \\renewcommand{\\arraystretch}{0.9} \\begin{tabular}{|ccrrrc|}\\hline GPS& $\\nu _{max}$ &  $S_{max}$ & $\\ \\ \\ \\alpha _{thick} $ &$\\ \\ \\ \\alpha _{thin}$ & FWHM\\\\ Source& (GHz) & (mJy) & & & log(freq)\\\\ \\hline B0400+6042&      1.0&         184&     0.52&    -1.48&      0.7\\\\ B0436+6152&      1.0&         237&     0.79&    -1.01&      0.7\\\\ B0441+5757&      6.4&         109&     0.30&    -0.50&      1.9\\\\ B0513+7129&      1.5&         242&     0.47&    -0.81&      0.9\\\\ B0531+6121&      5.9&          36&     0.09&    -1.12&      1.0\\\\ B0535+6743&      5.7&         192&     0.27&    -0.94&      1.1\\\\ B0537+6444&      2.3&          29&     0.31&    -1.10&      1.7\\\\ B0538+7131&      4.2&          85&     0.61&    -1.47&      0.7\\\\ B0539+6200&      1.9&         129&     0.67&    -0.33&      0.9\\\\ B0543+6523&      1.2&          69&     0.66&    -0.96&      1.0\\\\ B0544+5847&      1.4&          63&     0.45&    -1.16&      0.9\\\\ B0552+6017&      1.0&          50&     0.91&    -1.16&      0.5\\\\ B0557+5717&      1.1&          69&     0.85&    -1.20&      0.6\\\\ B0601+5753&      4.4&         187&     1.37&    -0.13&      0.9\\\\ B0748+6343&      6.6&          92&     0.37&     0.36&      1.1\\\\ B0752+6355&      6.4&         314&     1.79&    -0.10&      1.4\\\\ B0755+6354&      4.2&          28&     0.03&     0.10&      2.0\\\\ B0756+6647&      3.4&         127&     0.54&    -0.07&      0.8\\\\ B0758+5929&      2.0&         215&     0.51&    -0.40&      1.0\\\\ B0759+6557&      1.7&          46&     1.01&    -0.97&      0.6\\\\ B0826+7045&      3.5&         105&     0.79&    -0.20&      0.8\\\\ B0830+5813&      1.6&          65&     0.32&    -0.51&      1.1\\\\ B1525+6801&      1.8&         163&     0.38&    -1.07&      1.0\\\\ B1538+5920&      3.5&          64&     0.32&    -0.77&      1.0\\\\ B1550+5815&      4.6&         293&     0.41&    -0.02&      0.9\\\\ B1551+6822&      1.5&          52&     0.56&    -1.65&      0.8\\\\ B1557+6220&      2.3&          49&     0.40&    -1.89&      0.9\\\\ B1600+7131&      1.7&         346&     1.70&    -1.56&      0.3\\\\ B1620+6406&      2.2&          47&     0.17&    -1.56&      1.0\\\\ B1622+6630&      4.0&         363&     1.55&    -0.46&      0.5\\\\ B1639+6711&      1.0&          68&     0.92&    -0.61&      0.8\\\\ B1642+6701&      1.3&         124&     0.02&    -1.01&      2.0\\\\ B1647+6225&      0.9&          71&     1.03&    -2.53&      0.5\\\\ B1655+6446&      1.0&          69&     1.29&    -0.99&      0.5\\\\ B1657+5826&      0.5&          64&     0.19&    -0.56&      0.7\\\\ B1746+6921&      2.2&         164&     0.63&    -0.41&      1.1\\\\ B1807+5959&      1.0&          47&     1.56&    -0.90&      1.4\\\\ B1807+6742&      0.8&          54&     0.94&    -0.70&      0.6\\\\ B1808+6813&      1.3&          42&     0.20&    -0.55&      1.7\\\\ B1819+6707&      0.8&         338&     0.35&    -0.54&      1.0\\\\ B1841+6715&      2.1&         178&     1.53&    -0.63&      0.8\\\\ B1843+6305&      1.9&          75&     1.53&    -0.90&      0.7\\\\ B1942+7214&      1.4&         233&     0.72&    -0.50&      1.0\\\\ B1945+6024&    $>15$&      $>188$&     0.54&     0.70&        -\\\\ B1946+7048&      1.8&         929&     0.91&    -0.64&      0.6\\\\ B1954+6146&      8.4&         169&     0.00&    -0.31&      1.4\\\\ B1958+6158&      3.3&         142&     0.52&    -0.23&      0.9\\\\ \\hline \\end{tabular} \\end{center} \\caption{\\label{GPS}The resulting sample of GPS sources.} \\end{table}  \\begin{figure*} \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen0400+6042.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0436+6152.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0441+5757.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0513+7129.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen0531+6121.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0535+6743.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0537+6444.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0538+7131.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen0539+6200.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0543+6523.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0544+5847.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0552+6017.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen0557+5717.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0601+5953.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0748+6343.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0752+6355.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen0755+6354.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0756+6647.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0758+5929.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0759+6557.ps,width=5.1cm} } \\caption{\\label{spectra} Radio spectra of individual sources. Crosses indicate MERLIN data, diamonds indicate at 1.4 GHz NVSS data and at 5 GHz WSRT data.} \\end{figure*}  \\addtocounter{figure}{-1}  \\begin{figure*} \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen0826+7045.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen0830+5813.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1525+6801.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1538+5920.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen1550+5815.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1551+6822.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1557+6220.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1600+7131.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen1620+6406.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1622+6630.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1639+6711.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1642+6701.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen1647+6225.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1655+6446.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1657+5826.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1746+6921.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen1807+5959.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1807+6742.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1808+6813.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1819+6707.ps,width=5.1cm} } \\caption{{\\it Continued...}} \\end{figure*} \\addtocounter{figure}{-1}  \\begin{figure*} \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen1841+6715.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1843+6305.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1942+7214.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1945+6024.ps,width=5.1cm} } \\vspace{-0.5cm} \\hbox{\\hspace{-0.8cm} \\psfig{figure=snellen1946+7048.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1954+6146.ps,width=5.1cm}\\hspace{-0.8cm} \\psfig{figure=snellen1958+6158.ps,width=5.1cm} } \\caption{{\\it Continued...}} \\end{figure*} ",
        "conclusions": "A sample of GPS sources has been selected from the Westerbork Northern Sky Survey, with flux densities one to two orders of magnitude lower than bright GPS sources investigated in earlier studies. Sources with inverted spectra at frequencies $>325$ MHz have been observed with the WSRT at 1.4 and 5 GHz and with the VLA at 8.6 and 15 GHz to select genuine GPS sources. This has resulted in a sample of 47 GPS sources with peak frequencies ranging from $\\sim$500 MHz to $>$15 GHz, and peak flux densities ranging from $\\sim$40 to $\\sim$900 mJy.  Five GPS sources in our sample show extended emission or nearby components in the NVSS maps at 1.4 GHz. However it is not clear if these components are related to the GPS sources.  About 30\\% of the objects show flux density differences greater than 20\\% between the Greenbank and MERLIN 5 GHz measurements, with the Greenbank data points all higher than the MERLIN observations. We believe this is due to variability, and that the lack of sources with reverse variability (the MERLIN flux density greater than the Greenbank flux density) is due to a selection effect caused by the ``old'' epoch (1987) of the Greenbank observations.  GPS source counts are comparable to 1/250 of the 2 GHz source counts for large scale radio sources, if the latter sources were to have 10 times their measured flux densities. Unfortunately, apparent differences in redshift distributions between GPS and large scale radio sources hamper a direct and straightforward interpretation of the source counts. Potentially, the comparison of GPS source counts with that of large scale radio sources can provide clues about the age of GPS sources and their luminosity evolution. If it is assumed that the redshift distributions are the same for GPS and large size radio sources, the source counts indicate that GPS sources have to decrease in luminosity by a factor of $\\sim 10$ if they all evolve into large scale radio sources."
    },
    "9803/astro-ph9803283_arXiv.txt": {
        "abstract": "We present the results of a numerical code that combines multi-zone chemical evolution with 1-D hydrodynamics to follow in detail the evolution and radial behaviour of gas and stars during the formation of elliptical galaxies. We use the model to explore the links between the \\ev\\ and formation of elliptical galaxies and QSO activity. The  knowledge of the radial gas flows in the galaxy allows us to trace metallicity gradients, and, in particular, the formation of a high-metallicity core in ellipticals. The high-metallicity core is formed soon enough to explain the metal abundances inferred in high-redshift quasars. The star formation rate and the subsequent feedback regulate the episodes of wind, outflow, and cooling flow, thus affecting the recycling of the gas and the chemical enrichment of the intergalactic medium. The \\ev\\ of the galaxy shows several stages, some of which are characterized by a complex flow pattern, with inflow in some regions and outflow in other regions. All models, however, exhibit during their late \\ev\\ a \\gw\\ at the outer boundary and, during their early \\ev, an inflow towards the \\gal\\ nucleus. The characteristics of the inner inflow could explain the bolometric luminosity of a quasar lodged at the galaxy centre as well as the evolution of the optical luminosity of quasars. ",
        "introduction": "Imaging studies of the faint extensions around QSOs indicate that \\el\\ \\gals\\ are the host \\gals\\ of the radio-loud and the brightest QSOs (Smith et al. 1986; Hutchings, Janson \\& Neff 1989; Hutchings et al. 1994; Hutchings \\& Morris 1995; McLeod \\& Rieke 1995; Aretxaga, Boyle \\& Terlevich 1995; Disney et al. 1995; Bahcall, Kirkakos \\& Schneider 1996; Ronnback et al. 1996; Taylor et al. 1996). As a matter of fact, at high redshifts ($z  > 2$), the only galactic systems available to harbour QSOs are the spheroids, since the disks are formed much later. In addition, the epoch of completion of the large spheroids ($z  > 2$) coincides with the peak in the QSO activity ($2 < z < 3$) (Schmidt et al. 1991), suggesting a relation between the QSO phenomenon and the \\for\\ of large \\el s. In fact, in recent years there has been increasing evidence linking QSO activity with galaxy formation. The high metal content of high redshift QSOs, the high dust content (several $10^8$ \\msun\\ of dust) of distant QSOs (Andreani, Franca \\& Cristiani 1993; Isaak et al. 1994; McMahon et al. 1994; Omont et al. 1996) plus the possible relation between the galaxy luminosity function (LF) and the QSO LF (Terlevich \\& Boyle 1993, hereafter TB93; Boyle \\& Terlevich 1998, hereafter BT98) provides tantalising evidence of this link. In any case, the fact that even the highest redshift QSO has strong metal lines in its spectrum requires that the broad line region (BLR) gas has been enriched by a stellar population formed before $z\\sim 5$. Work by TB93 and Hamann \\& Ferland (1993) (hereafter HF93) highlighted the importance of metal production in the early evolution of a galaxy. Since QSOs are seen up to redshifts of $z\\sim 5$, the supersolar metallicities required by the BLR models should be reached by $\\sim 1$~Gyr since the beginning of the galaxy formation epoch. This evolutionary time scale is an important constraint in chemical enrichment models.  Two fundamental relations involving intrinsic parameters of elliptical galaxies show remarkable little dispersion and point towards an early formation of ellipticals. They are the colour-luminosity relation and the ``fundamental plane'' relating the total luminosity of an elliptical to its central velocity dispersion and surface brightness.  The tightness of the colour-luminosity relation provides strong evidence that most of the present stellar population was formed at $z > 2$ (Bower, Lucey and Ellis 1992). The narrowness of the ``fundamental plane'' gives additional support to that conclusion (Renzini and Ciotti 1993) and in addition indicates that the properties of the core (velocity dispersion, Mg$_2$ strength) are intimately linked to the galaxy global ones (D$_n$, luminosity). There is also good evidence that the stellar population in massive ellipticals is metal rich with respect to the Sun, and shows large radial metallicity gradients (e.g., Worthey, Faber and Gonzalez, 1992; Davies, Sadler and Peletier 1993). These properties of the \\el\\ \\gals\\ may also be related to the fact that they probably harbour QSOs in their centres at some stage of their \\ev.  One-zone chemical evolution models have been used (Hamann \\& Ferland 1992, HF93, Padovani \\& Matteucci 1993, Matteucci \\& Padovani 1993) to investigate the chemical history and the fueling of QSOs. However, during the early evolution of the elliptical galaxy, it is expected several episodes of gas outflow and inflow which cannot be followed by the one-zone model. In addition, the one-zone chemical evolution models that attempt to explain the high metal content in high redshift QSOs tend to overproduce metals (averaged over the entire galaxy) and predict an excessively high luminosity for the parent galaxy of the QSO. For example, HF93 model M4 reproduces the rapid metal production needed but it is overluminous. In this model, an elliptical of $10^{11}$ \\msun\\ has a peak bolometric luminosity of $\\sim 2 \\times10^{13}$ \\lsun\\ at an age of $\\sim 0.1$ Gyr. But  an elliptical with $10^{11}$ \\msun\\ is at present only a sub-$L^*$ galaxy (with a blue luminosity of $\\sim 0.3 L^*$, for $M_B^*=-21$ and $[M/L_B]=10$), so that for the most luminous ($M_B \\approx -24$) ellipticals in the nearby Universe HF93 model M4 predicts luminosities of up to $10^{15}$ \\lsun\\ during its formation. These luminosities are higher than the QSO luminosities!!  Note that only the core of the galaxy has to be metal rich in accordance with the observed metallicity gradients in nearby galaxies. In the starburst model for QSOs (TB93 and references therein), the QSOs are the young cores of massive ellipticals forming most of the dominant metal-rich population in a short starburst. The core mass, which participates in the starburst, comprises only a small fraction ($\\sim 5$ \\%)  of the total galaxy mass. Also in the standard supermassive black hole model for QSOs, only $0.5-1$ per cent of the \\gal\\ mass goes into the black hole (Haehnelt \\& Rees 1993). The excessive production of energy and metals in the one-zone model arises, therefore, from its inability to resolve the core of the \\gal.  The QSO LF undergoes strong \\ev\\ between $z=2$ and the present epoch, with the redshift dependence of the LF being well-described by a constant comoving space density  and a pure power-law luminosity \\ev\\ $L(z) \\propto (1+z)^k$ (Boyle et al. 1988, 1991, BT98), with $k$ in the range $3.1<k<3.6$. This luminosity \\ev\\ is linked to the mass flow into the galactic nucleus hosting the QSO, both in the supermassive black hole scenario and in the starburst model for AGN. Since the luminosity is expected to be proportional to the mass accretion rate into the nucleus, in order to make predictions about the luminosity \\ev\\ of the QSO one needs the \\ev\\ of the central gas inflow. Again, this piece of information is not provided by the one-zone model.  In view of the limitations of one-zone chemical models, we have developed a multi-zone chemo-dynamical model that self-consistently combines chemical evolution and numerical hydrodynamics. In this paper, we use this code to investigate a number of topics related to \\for\\ and \\ev\\ of \\el\\ \\gals\\ and the link young \\el s-QSOs: 1) the formation of a high-metallicity core in ellipticals; 2) the radial metallicity gradients in ellipticals; 3) the chemical enrichment of the intracluster medium by ellipticals; 4) the evolution of the luminosity of young cores of \\el s/QSOs. The approach of the present work is first to develop a realistic sequence of models which reproduces the main properties of \\el\\ \\gals. We then investigate the \\ev\\ of the \\gal\\ core and of the gas inflow into the nucleus and, within the scenario in which the galaxy nucleus hosts a QSO, we compare the predictions of the model with QSO observations.  The chemo-dynamical evolution code is described in Section 2. Section 3 presents the evolution of the interstellar medium (ISM) and of the \\sfor. Section 4 considers the chemical enrichment of the intracluster medium by \\gw s from elliptical galaxies. Section 5 is devoted to \\ab s  and metallicity gradients of the stellar population in ellipticals. The predictions of our models for the \\ev\\ of the QSO LF are presented in Section 6. Some  concluding remarks are given in Section 7. ",
        "conclusions": "In this work, we explored the relation between young \\el\\ \\gals\\ and QSOs within a chemo-dynamical model for \\ev\\ of \\gals. In this model, we perform a multi-zone modelling of  the chemical enrichment of the gas and stars of a  \\gal\\ taking into account the gas flow obtained self-consistently from 1-D hydrodynamical calculations. We were particularly interested in the \\cf\\ towards the centre of the \\gal, which could feed a QSO hosted in the galactic nucleus.  From a minimal set of assumptions, based on standard one-zone \\chev\\ models, our model reproduces the main observational features of \\el\\ \\gals: 1) the central metallicities of massive \\gals\\ are supersolar; 2) the ratio [Mg/Fe] is supersolar in the core of the \\gal\\ as well as over the whole \\gal; 3) the \\gal\\ shows sizable metallicity gradients; 4) systems with larger masses tend to have larger metallicities, thus reproducing the mass-metallicity relation of \\el\\ \\gals; 5) elliptical galaxies can be the main responsibles for the ICM iron content: the ICM iron mass per unit luminosity of cluster galaxies ($\\sim 10^{-2}$ \\msun/\\lsun) is reproduced.  One very important time scale for the enrichment of the ICM is the time $t_{w,e}$ of the end of early wind phase, during which the gas content in the \\gal\\ is reduced by a factor of ten. $t_{w,e}$ is longer than $\\sim 1.5$ Gyr for normal or bright \\gals. For the iron enrichment, the relevant time scale is even longer, because the Fe-rich gas front arrives at the galaxy edge  later  than $t_{w,e}$.  From the success of the model in reproducing \\el\\ \\gals, a number of standard assumptions of the classic one-zone chemical \\ev\\ models remain valid when the hydrodynamics is considered: $10^8$ yr \\sfor\\ time scale, Salpeter IMF, stellar yields, $A_{SN\\,I}=0.1$. However, the linear \\sfor\\ law usually adopted by one-zone models seems to be excluded, since the resulting model would be unrealistic, not reproducing the properties of \\el\\ \\gals, as seen from the drawbacks of model 2(0):  too low central metallicities; too high [Mg/Fe] ratio; the nucleus mass is too small to explain the luminosity of a QSO.  We should take note of two possible discrepancies between the results of our models and the observations: 1) supersolar metallicities are predicted for the hot gas in the galactic halo, while ASCA measurements imply subsolar abundances; 2) we predict a trend of [Mg/Fe] decreasing with galactic mass, which is the opposite of the trend inferred from determinations of metal indices. If these contradictions are real, they could indicate that some of the standard assumptions of the \\chev\\ modelling do not apply. We should be aware, however, of the uncertainties in deriving the abundances in the hot galactic halos (besides the dilution by the ICM, which lowers the metallicity in the halo), and of the difficulties in obtaining [Mg/Fe] trends from the observed variation of Mg and Fe line strengths with the galaxy mass.  In addition to the fact that the model satisfy a number of observational constraints on \\el s, it makes definite predictions about the relation between the \\ev\\ of \\el s and that of QSOs.  A high-metallicity core is rapidly built-up. For the gas, solar metallicities are reached in $10^8$ yr for oxygen, and $3\\times 10^8$ yr for iron. For the stars, the magnesium abundance becomes solar at $2\\times 10^8$ yr and the iron at $8\\times 10^8$ yr. In this way, the high abundances derived for high redshift QSOs are reached in a reasonably short time scale. One important application of these results is that the enrichment time scales predicted by our chemo-dynamical model provide a chemical clock which could constrain cosmological scenarios. For instance the $\\sim 1$ Gyr time scale for metal production implies an age at least larger than 1 Gyr for the Universe at $z \\sim 5$, since even the more distant QSOs exhibit metals in their spectra. In addition, the short time scales for enrichment both of gas and stars make plausible the starburst model for AGN, which requires a high metallicity core formed early in the \\ev\\ of the \\gal.  It is important to note that the luminosity and metallicity of QSOs identified with the central region of the young elliptical are explained with no need for all the galaxy having a global starburst coordinated with the central starburst. In this way, extremely high luminosities ($\\sim 10^{15}$ \\lsun) are avoided for the proto-QSO. These extreme luminosities, predicted by one-zone models of formation of \\el\\ \\gals\\ (e.g. model M4 of HF93), have never been detected in high redshift observations. Note that probably only the inner regions of the young \\gal, due to their higher surface brightness, would have \\sfor\\ detectable at high redshifts. As a matter of fact, deep spectroscopy observations have revealed a population of star-forming \\gals\\ at redshift $3 \\la z \\la 3.5$, the Lyman break \\gals\\ (LBGs), discovered on the basis of a Lyman limit break superposed on their UV continuum (Steidel et al. 1996; Giavalisco, Steidel \\& Macchetto 1996). The LBGs show many characteristics expected for primeval \\gals, and, assuming a Salpeter IMF, their inferred SFRs are in the range $4-90 h_{50}^{-2}$ \\msun\\ yr$^{-1}$, inside a typical half-light radius of $1.8 h_{50}^{-1}$ kpc ($q_0=0.5$). These levels of \\sfor\\ are remarkably close to the typical range of $20-50$ \\msun\\ yr$^{-1}$ for the SFR inside a radius of 1 kpc in the fiducial model. In addition, the LBGs are expected to be the high-redshift counterparts of the present-day spheroidal component of luminous \\gals, since their co-moving density is roughly comparable to that of present-day bright ($L \\geq L^*$ )\\gals\\ and the widths of the interstellar absorption lines in their spectra imply circular velocities of $170-300$ km s$^{-1}$, typical of the potential well depth of luminous \\el s. Therefore, for $z \\la 3.5$, the observations of ongoing \\sfor\\ in \\gals\\ seems to rule out the high luminosities predicted by the one-zone models for the progenitors of present-day $\\sim L^*$ \\gals.  The luminosities of the one-zone model could still be consistent with the observations of LBGs, if there is dust absorbing the blue and UV light and re-emitting it in the far-IR/sub-mm. However, comparisons between the observed colours of LBGs and those predicted by spectral synthesis models (Pettini et al. 1998) indicate only modest dust attenuation (an extinction at 1500 \\AA\\ between lower and upper limits of $\\sim 2$ and $\\sim 6$), and, therefore, dust absorption is unable to hide the high luminosities of the one-zone model. In addition, in view of the evidence that the LGBs are the progenitors of the present-day luminous ($\\sim L^*$ or brighter) \\gals, rather than sub-units being assembled into a larger system, it seems that the LBGs are allowing us to witness one stage in a single event of formation of a massive \\gal, with the global star formation proceeding, however, at a milder rate than in the one-zone model. In the end, the main reason why the one-zone model overpredicts the luminosity is that it overproduces metals. The model M4 of HF93 predicts $\\sim 10$ times the solar metallicity over the whole \\gal. The observations, however, allow this extremely high metallicity only at the very nucleus of the \\gal, at most. Over an effective radius, the metallicity of a giant \\el\\ is roughly solar. These metallicities are correctly predicted by our multi-zone model (see the values of $\\langle$[Mg/H]$\\rangle_{10}$ and $\\langle$[Fe/H]$\\rangle_{10}$ in Table 3). Moreover, assuming that the observed negative metallicity gradients continue beyond the effective radius, the mass-averaged metallicity would be subsolar over the whole \\gal. Therefore, scaling the metals produced in the one-zone approximation to amounts consistent with the observations, brings down the predicted luminosities by about one order of magnitude.  All our models predict a massive central (through the inner 100 pc) \\cf\\ during the first 1-2 Gyr of the \\gal\\ \\ev. Within the scenario in which the luminosity of the galactic nucleus is fed by the central inflow, for a reasonable epoch of formation of the spheroidal systems, the epoch of building-up of the nucleus by the central inflow coincides with the maximum in QSO activity ($2<z<3$). Also the decrease of the central inflow rate for $t>1$ Gyr is consistent with the decline of luminosity inferred for $z<2$ from the \\ev\\ of QSO LF with redshift.  One of our models exhibits  recurrent late central \\cf\\ episodes, which are brief (a few $10^7$ yr) and involve decreasing amounts of mass. The gas inflow into the inner 100 pc is regulated by episodes of star formation leading naturally to several short episodes of central inflow, thus giving support  to the episodic scenario for evolution of the LF of QSOs.  The model also explains the luminosities of QSOs. The central \\cf\\ rates explain bolometric luminosities of up to $10^{47}$ erg s$^{-1}$, for an efficiency $f$ of mass-energy conversion of 0.1. However, the bolometric luminosities of the brightest QSOs ($\\approx 10^{48}$ erg s$^{-1}$) cannot be explained by a continuous deposition of the central inflow. Rather, the highest QSO luminosities require a discrete deposition, in which the gas is  accumulated during $\\approx 1$ Gyr and then consumed by a central engine in few $10^7$ yr. Accordingly, we have made some predictions on the QSO LF at $z \\ga 1$ based on a simple discontinuous model for QSO activity, in which there is two short gas consumption events during the first central inflow episode. We scaled the QSO LF to the present day elliptical LF, assuming that all \\el s have harboured a QSO during their \\ev. Both the starburst and the supermassive black hole models predict the right shape of the QSO LF, but the nuclear starburst systematically underestimates the density  number of QSOs. In addition, our model reproduces the \\ev\\ of the LF between $z=1.25$ and $z=2.9$. In our models, the mass deposited by the first central inflow represents 0.15-0.5 \\% of the present day luminous mass of the \\gal. Assuming that this mass goes to the formation of a central object (star cluster or black hole), the model correctly predicts for the Dark Massive Objects (DMOs) in the nuclei of  \\gals, both the masses and the DMO-to-galaxy mass ratios.  Another conclusion derived from our models is that the hosts of high-redshift AGN should be relatively mature objects. The calculated \\ev\\ of the inner \\cf\\ and energetic considerations imply that the gas deposited by the inflow in the nucleus should accumulate during $\\sim 1$ Gyr before triggering a short-lived AGN activity event. On the other hand, except for extremely large \\gals\\ (present-day $M_B=-24$), the first, massive \\cf, responsible for maintaining the AGN activity, lasts for $2-3$ Gyr. In this scenario, if the high redshift \\gal\\ is to display strong AGN activity, its probable age would range from $\\sim 1$ to $\\sim 3$ Gyrs. The minimum age of $\\sim 1$ Gyr for QSO hosts derived above is consistent with the  $\\sim 1$ Gyr time scale for metal enrichment, needed to explain the strong metal lines observed even in $z\\sim 5$ QSOs.  One further prediction of our models is that, at the present, only the most massive objects should be host to powerful AGN, but at high redshift powerful AGN activity is expected even for hosts of smaller mass. Interestingly enough, this is what seems to be observed, for radio \\gals\\ at least. Models 1 and 2 are examples of sub-$L^*$ \\gals\\ with strong AGN activity at high redshift. Note that $\\zeta$ increases with decreasing \\gal\\ mass for masses below $M_G \\approx 10^{12}$ \\msun. Since this parameter describes the relative importance of the first \\cf\\ episode occuring when the \\gal\\ is less than $2-3$ old, this means that for the lower mass systems, the efficiency of building-up the nucleus is higher than in larger systems, and that they have an early \\cf\\ massive enough to sustain a strong AGN activity. However, this activity is limited only to the 2-3 first Gyr of the \\gal\\ and, therefore, smaller systems with strong nuclear activity are to be found only at high redshift. Even in the case of the episodic scenario of model 2(1/3), the late nuclear activity is much weaker at low redshift (in this model, the two late central inflow episodes deposit only 1.8 \\% and 1 \\% of the mass of the first inflow episode). On the other hand, models 10, 20 and 50 exhibit a present-day massive central \\cf\\ which could trigger AGN activity. In these massive systems, the late cooling flow may accumulate into the nucleus an amount of mass comparable to that of  the early \\cf\\  (model 20 is typical of this case, the late \\cf\\ deposits $1.65\\times 10^9$ \\msun, i.e. 47 \\% of the mass of the first \\cf). Only these objects, therefore, would harbour, intense AGN activity at low redshift. Support to this picture is given by recent analysis of host \\gals\\ of powerful nearby ($z \\la 0.3$) AGN, belonging to three samples --- radio \\gals\\, radio-loud quasars, and radio-quiet  quasars (Taylor et al. 1996). For all three classes of AGN, the host \\gals\\ are large (half-light radius $r_{1/2} \\geq 10$ kpc) and luminous (K-band luminosity $L_K\\geq L^*$). (Note that the less massive model exhibiting a present-day central is model 10, with $M_B=-21.8$, or $L_B=2 L^*$.)  Finally, we should note that, although pure luminosity \\ev\\ seems to reproduce the \\ev\\ of the QSO LF, in our models the individual QSOs do not dim over cosmological time scales, but rather are short-lived (i.e., $t_{on}=1-3\\times 10^7$ yr). In order to comply with the energetic requirements of the QSOs, their activity must occur in short episodes of massive consumption of mass accumulated in the galactic nucleus during a much longer span of time ($\\sim 1$ Gyr). This sort of episodic activity displayed by our model for QSOs seems to be a rule among the AGN in general, since for other class of AGN, the radio galaxies, the radio LF also seems to follow pure luminosity \\ev, and yet the radio sources themselves seem to have lifetimes of only a few $10^7$ yr. As a matter of fact, a comparison between the properties of the radio-loud population and the present model would be very useful to clarify the relation of the QSO phenomenon and the early \\ev\\ of \\el\\ \\gals, since radio observations allow us to explore the $z>2$ domain without the need of uncertain corrections for dust absorption and lensing bias that hamper the optical tecniques."
    },
    "9803/astro-ph9803297_arXiv.txt": {
        "abstract": "We discuss the chemical evolution of dwarf irregular and blue compact galaxies in light of recent data, new stellar yields and chemical evolution models. We examine the abundance data for evidence of \\hii\\ region self-enrichment effects, which would lead to correlations in the scatter of helium, nitrogen, and oxygen abundances around their mean trends.  The observed helium abundance trends show no such correlations, though the nitrogen--oxygen trend does show strong evidence for real scatter beyond observational error. We construct simple models for the chemical evolution of these galaxies, using the most recent yields of \\he4, C, N and O in intermediate- and high-mass stars. The effects of galactic outflows, which can arise both {}from bulk heating and evaporation of the ISM, and from the partial escape of enriched supernova ejecta are included. In agreement with other studies, we find that supernova-enriched outflows  can roughly reproduce the observed He, C, N, and O trends; however, in models that fit N versus O, the slopes $\\Delta Y/\\Delta$O and $\\Delta Y/\\Delta$N consistently fall more than $2\\sigma$ below the fit to observations. We discuss the role of the models and their uncertainties in the extrapolation of primordial helium from the data. We also explore the model dependence arising nucleosynthesis uncertainties associated with nitrogen yields in intermediate mass stars, the fate of $8-11 \\msol$ stars, and massive star winds. ",
        "introduction": "Helium-4 plays a central role in big bang nucleosynthesis.  The calculation of the primordial helium abundance is robust, being only weakly sensitive to the cosmic baryon-to-photon ratio, $\\eta$. However, the helium abundance is quite sensitive to, and provides a strong constraint on the physics of the early universe. The primordial abundance of helium (with mass fraction \\yp) is best determined via observations of \\hii\\ regions in the most metal-poor galaxies---dwarf irregulars and  blue compact galaxies (hereafter, BCGs). For these systems, the helium evolution is derived empirically:  the data show that the helium abundance increases with metallicity indicators such as oxygen and nitrogen. Following Peimbert \\& Torres-Peimbert, \\pcite{ptp}, the primordial abundance of helium is inferred from an extrapolation of the observed trend to zero metallicity.  As emphasized by Fields \\& Olive\\pcite{fo}; Fields, Kainulainen, Olive, \\& Thomas \\pcite{fkot}; and Olive, Steigman, \\& Skillman (\\cite{oss}; hereafter OSS97), the empirical extrapolation of primordial helium sidesteps any reliance on detailed results of chemical evolution models, especially since the extrapolation to zero metallicity from the data is small. Indeed, until the measurements of D/H in quasar absorption systems can be confirmed to represent a uniform primordial value, the \\he4 abundance is crucial when used in conjunction with \\li7 as a test of big bang nucleosynthesis theory.  (This remains true so long as $\\yp \\la 0.24$; at high \\yp, the insensitivity of \\yp\\ to $\\eta$ makes \\he4 a very poor discriminator for primordial nucleosynthesis.) Thus, given the importance of \\he4 as observed in BCGs, it is clearly of interest to compare the predictions of chemical evolution with the observations of these systems. A less than successful comparison of theory with observation could point to the subtleties in chemical evolution modeling of even these (apparently) simple systems.  Several groups have modeled the chemical evolution of BCGs (Matteucci \\& Chiosi \\cite{mc}; Matteucci \\& Tosi \\cite{mt}; Gilmore \\& Wyse \\cite{gw}; Pilyugin \\cite{pil}; Clayton \\& Tantelaki \\cite{cp}; Marconi, Matteucci \\& Tosi \\cite{mmt}; Carigi, Col\\'{\\i}n, Peimbert, \\& Sarimento \\cite{ccps}). The high star formation activity in BCGs indicates that star formation must occur in stochastic bursts, unless the systems are very young; all models included this behavior. In addition, these galaxies probably drive outflows of material, perhaps with efficiencies high enough to significantly alter the chemical and/or dynamical evolution. For example, the gas fraction--metallicity relations of BCGs may not be compatible with closed box models, suggesting the importance of gas outflow via supernova-driven outflows (Lequeux, Rayo, Serrano, Peimbert, \\& Torres-Peimbert \\cite{leq}; Matteucci \\& Chiosi \\cite{mc}; Carigi, Col\\'{\\i}n, Peimbert, \\& Sarimento \\cite{ccps})---though the difficulties of determining an accurate gas--to--baryon fraction lead to significant uncertainties in these arguments. In any case, X-ray observations of diffuse, hot gas clearly support the existence of outflows in some BCGs (Heckman et al.\\ \\cite{heck}; Della Ceca, Griffiths, Heckman, \\& Mackenty \\cite{dc_etal}; Della Ceca, Griffiths,  \\& Heckman \\cite{dcgh}), though Skillman \\pcite{evan96} and Skillman \\& Bender \\pcite{sb} present arguments against a dominant role of outflows in {\\em all} dwarf galaxies. Independent of the question of outflows, another key result from these studies showed that the usual instantaneous {\\em mixing} approximation is inappropriate. Self-enrichment of \\hii\\ regions in a burst phase (Kunth \\& Sargent \\cite{ksgt}; Kunth, Lequeux, Sargent, \\& Viallefond \\cite{klsv}; but see Pettini \\& Lipman \\cite{pl}) can lead to significant scatter in abundance trends.  Most recently, it has been argued that the due to the intense bursts of star formation and energy release in the evolution of these systems, their stability requires of the presence of significant amounts of dark matter (Bradamante, Matteucci, \\& D'Ercole \\cite{bmd}).  This may have an impact on the effect of supernovae driven winds on the elements abundances. In a somewhat different approach, Mac Low \\& Ferrara \\pcite{mlf} take the presence of dark matter as a starting point, and examine the requirements for BCG outflows as a function of supernova rate and galaxy mass. Their models of the the dark matter potential, and the detailed gas dynamics of outflows, suggest that the loss of the entire ISM---``blowaway''---is only possible for smaller galaxies ($\\la 10^6 \\msol$). However, Mac Low \\& Ferrara also find that the preferential escape of the energetic and metal-enriched supernova ejecta is much easier, occurring at significant levels for galaxies with baryonic masses up to $\\sim 10^8 \\msol$.  While there has been great progress in laying out the basic features of BCG chemical evolution, some important issues remain unresolved, including some bearing on derivation of primordial helium. For example, there remains some question as to the theoretical justification for the linear fit of helium versus metals, particularly nitrogen (e.g., Mathews, Boyd, \\& Fuller \\cite{mbf}; Balbes, Boyd, \\& Mathews \\cite{bbm}). Closely related to this is the dependence of the inferred \\yp ~on the metallicity tracer. Several studies (e.g., OSS97; Olive, Steigman, \\& Walker \\cite{osw}) have found that the primordial helium abundance derived from a linear fit to $Y-$N consistently differs  by about $1\\sigma$ {}from that derived via $Y-$O. While this difference is as yet too mild to be a problem, it does suggest need to re-examine the appropriateness of the linear fits in the context of detailed chemical evolution models. Recent new observations such as the growing set of carbon observations are ripe for more theoretical attention.  In this paper, we examine the chemical evolution of BCGs, with particular attention to these issues related to primordial helium, and the uncertainties in the chemical evolution modeling. In particular, we note the interplay between galactic outflow prescriptions and nucleosynthesis uncertainties. Different parameterizations for the nucleosynthesis of nitrogen and their role \\yp(N) versus \\yp(O) discrepancy are also considered. We note the model-dependence of $\\Delta Y/\\Delta Z$ and other slopes, and examine in detail a suggestion by Fields \\pcite{fields} that helium production processes due to low-mass supernovae or high-mass stellar winds may account for the large observed values of the slopes. We conclude that the empirical fitting procedures used to obtain \\yp\\ do find support from our models. However, we also note that models using the most recent and detailed nucleosynthesis yields do no reproduce the observed $Y-$N,O trends, even in the presence of outflows. Possible solutions to this problem are discussed. ",
        "conclusions": "We have considered the chemical evolution of BCGs, through an analysis of the observed HeNO abundances, and by studying chemical evolution models.  Our basic conclusions are as follows.  The scatter in the $Y-$N and $Y-$O data is entirely consistent with the observational errors.  In particular, these data do not show the correlations (in the N--O plane) that one would expect if the scatter were due to self-enrichment by massive star ejecta.  On the other hand, it appears that the N--O scatter real. We expect that at least some of the observed scatter is due to self-enrichment, as is also suggested by the observed correlation of high C/O systems with high N/O.  Our study of chemical evolution models of BCGs concludes that  the favored models for these systems (i.e., including outflows), coupled with recent and detailed stellar yield calculations (which are metallicity dependent), are not successful in reproducing simultaneously all of the abundance trends. We do find that in general, the models do predict linear slopes in He versus N and O, as has long been assumed in phenomenological fits to the data. The large slopes are, however, hard to reproduce quantitatively---even in models with supernova-enriched outflows. The only models which {\\em can} produce large slopes, while (possibly) maintaining agreement with N--O data, are those using the (less detailed) stellar yields of Maeder \\pcite{mae} in combination with a log-normal IMF.  Consequently, we suspect that the root of the problem lies in the nucleosynthesis inputs, particularly the He yields.  We emphasize that any means of improving the helium slopes must be tested in models which fit {\\em all} of HeNO, and preferably C as well.   Another shortcoming of the models regards nitrogen. We find that the models with the detailed yields typically do not predict a strong enough secondary character for N versus O.  As a result, they do not find that the $\\Delta Y/\\Delta $N slopes change much with metallicity.  In particular, the models do not reproduce the differences found in the $Y-$N slopes (and intercepts) between the full and low-metallicity data sets.  Despite these unresolved issues, our study of BCG models does allow us to draw some conclusions. Most models with detailed yields predicted $Y-$N,O relations that were close to linear; this gives some theoretical justification for the adoption linear regressions of the data (particularly for $Y$--O). Also, we confirm that the nature of the star formation rate---bursting versus smooth---does strongly affect the detailed evolutionary history, and an ensemble of different histories leads to significant scatter in N--O.  Fortunately, however, the N--O evolution in the smooth star forming history forms an upper envelope to the N--O trends in the bursting models. Thus, if one is not interested in the scatter, one can adopt a smooth star forming model as long as one is careful to fit appropriate envelope of the data. We have also shown that carbon data provides an important additional constraint on BCG chemical evolution.  In particular, C and N together can diagnose the tradeoff between C and N controlled by hot bottom burning. Additional C observations would be very useful in further constraining BCG models.  Finally, aspects of our results have implications for systems other than BGCs. The difficulty for our models to reproduce the helium slopes probably carries over to our own Galaxy.  Although there is some evidence that the Galactic $\\Delta Y/\\Delta Z$ may be lower than in BCGs, it nevertheless seems to exceed the values ($\\sim 1$) consistently predicted by all of our models without strong enriched winds. A good solar neighborhood $\\Delta Y/\\Delta Z$ would help illuminate whether the helium slope problem is due to yields, or due to environmental differences between the Galaxy and BCGs. Finally, we note the similarity between the low metallicity N--O evolution in BCGs, QSO absorbers (c.f. Lu, Sargent, \\& Barlow \\cite{lsb}), and other external galaxies (van Zee, Salzer, \\& Haynes \\cite{vzsh}). The observed trends are compatible with each other, at least to a first approximation.  This suggests that the N--O evolution of these objects is similar, which may then imply that the enriched outflow of BCGs is small.  If so, this is additional evidence that the heart of the problem of the low helium slopes lies in the stellar yields themselves."
    },
    "9803/astro-ph9803318_arXiv.txt": {
        "abstract": "We use high-resolution hydrodynamic simulations to investigate the density profile of hot gas in clusters of galaxies, adopting a variant of cold dark matter cosmologies and employing a cosmological N-body/smoothed particle hydrodynamics code to follow the evolution of dark matter and gas. In addition to gravitational interactions, gas pressure, and shock heating, we include bremsstrahlung cooling in the computation. Dynamical time, two-body relaxation time, and cooling time in the simulations are examined to demonstrate that the results are free from artificial relaxation effects and that the time step is short enough to accurately follow the evolution of the system. In the simulation with nominal resolution of $66h^{-1}\\,{\\rm kpc}$ the computed cluster appears normal, but in a higher (by a factor 2) resolution run, cooling is so efficient that the final gas density profile shows a steep rise toward the cluster center that is not observed in real clusters. Also, the X-ray luminosity of $7\\times10^{45}{\\rm ergs\\,s^{-1}}$ far exceeds that for any cluster of the computed temperature. The most reasonable explanation for this discrepancy is that there are some physical processes still missing in the simulations that actually mitigate the cooling effect and play a crucial role in the thermal and dynamical evolution of the gas near the center. Among the promising candidate processes are heat conduction and heat input from supernovae. We discuss the extent to which these processes can alter the evolution of gas. ",
        "introduction": "Hot X-ray--emitting gas in clusters of galaxies contains a variety of information relevant to many fields in astrophysics. Density and temperature profiles of the gas give among the most reliable estimates of cluster mass, which are of unparalleled importance to cosmology (e.g., \\cite{bah95}). They will also reflect physical processes that have played a crucial role in the thermal and dynamical evolution of the gas. In addition to gravitation, hydrodynamics, and shock heating, such processes may include radiative cooling, heat conduction, and feedback from star formation.  One of the central questions about the cluster gas structure involves the presence of a core. Observations have revealed that the gas density profile has a distinct core, inside of which density is nearly constant. Gravitational N-body simulations have demonstrated that gravity alone cannot produce such a core in an object formed as a result of hierarchical structure formation; halos formed in high-resolution N-body simulations have density profiles with significant slope toward the center up to the resolution limit (\\cite{nav96}, 1997; \\cite{fuk97}; \\cite{moo98}). For example, Navarro, Frenk, and White (1996, 1997) found that the density profiles of halos with a wide range of masses can be fitted by \\begin{equation} \\label{eqn:nfw} \\rho(r) = \\frac{\\rho_c}{(r/r_s)(1+r/r_s)^2} \\end{equation} (hereafter referred to as the NFW profile), with $\\rho_c$ and $r_s$ being the fitting parameters. In the NFW profile, gas temperature would approach zero in a cluster's central regions ($T \\propto r$) were it to have the same profile as dark matter. While this model does produce a convergent X-ray luminosity, the temperature structure is different from what is observed. It follows that processes other than gravity are responsible for the formation of the cores. Makino, Sasaki and Suto (1998) pointed out that the gas distribution develops a core in a dark matter halo having a central cusp (e.g., the NFW profile) if the gas is isothermal and in hydrostatic equilibrium. Even if the gas obeys a less stringent limit of constant entropy ($T/n^{2/3} \\to {\\rm const}$), it would avoid a central cusp and have an apparently constant-density, constant-temperature core. But it is not evident what physical mechanisms could enforce either isothermality or isentropy.  Many authors have studied the evolution of clusters using simulations that include hydrodynamics (\\cite{evr90}; \\cite{tho92}; \\cite{kat93}; \\cite{bry94a}, 1994b; \\cite{kan94}; \\cite{nav95}; \\cite{bar96}; \\cite{bry97}, 1998; \\cite{eke98}; \\cite{pen98}; \\cite{yos98}). In general these simulations have succeeded in producing clusters that have cores similar to those observed. However, the situation is far from satisfactory for at least two reasons. First, artificial two-body relaxation may affect the dynamics, especially in the central region. Indeed, Steinmetz and White (1997) showed that this effect gives rise to artificial energy transfer from dark matter to gas. Both spatial and mass resolutions are only marginally adequate in most published works. Second, almost none of the simulations include radiative cooling of gas. Although cooling is probably unimportant in the outer part of a cluster (\\cite{sar86}), it may affect the dynamics in the central region.  This work is an investigation of cluster gas using hydrodynamic simulations that attempt to address these limitations. First, we minimize two-body relaxation effects for a given computational cost by employing a multiresolution technique that enables us to improve resolution only inside the clusters where we really need high resolution. Second, we include cooling due to bremsstrahlung, which dominates cooling of gas at above $\\gtrsim 10^7{\\rm K}$. We ignore line cooling, but this is not a serious problem because we focus our attention on X-ray--emitting gas. A practical reason for ignoring line cooling is that doing so enables us to avoid very short timescales in moderate temperature ($10^4$ -- $10^6 {\\rm K}$), high density regions. Allowing for line cooling would strengthen the conclusions of this paper.  We organize the rest of the paper as follows. In \\S\\ 2 we describe the method and parameters of the simulations. We present the results in \\S\\ 3. As we will see, the most important result is that cooling and increased resolution give rise to a density profile of the gas that rises steeply toward the center and consequently produces excessive X-ray luminosity.  In \\S\\ 4 we discuss the implications of our results in connection with physical processes that are still missing in the simulations. In \\S\\ 5 we give our conclusions. ",
        "conclusions": "\\label{sect:conclusion}  Finally let us briefly summarize our conclusions. \\begin{enumerate} \\item Our multiresolution simulation has followed with reasonable accuracy the evolution of gas and dark matter in a typical cluster under the influence of bremsstrahlung cooling as well as gravity and hydrodynamics. \\item The high resolution simulation resulted in a gas density profile steeply rising toward the center, with consequent very high X-ray luminosity; however, these properties are not observed. \\item Heat conduction and SN heating are among the processes that may account for the discrepancy. Had we allowed for their likely importance in the real world we might have been able to recover the observed gas density profile. In a future work we will examine their effects by directly incorporating them into the simulation. \\end{enumerate}"
    },
    "9803/astro-ph9803154_arXiv.txt": {
        "abstract": "Recent calculations indicate that in the outer parts of neutron stars nuclei are rod-like or slab-like, rather than roughly spherical. We consider the elastic properties of these phases, and argue that they behave as liquid crystals, rather than rigid solids. We estimate elastic constants and discuss implications of our results for neutron star behavior. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803224_arXiv.txt": {
        "abstract": "We examine all possible stationary, optically thick, geometrically thin accretion disc models relevant for active galactic nuclei (AGN) and identify the physical regimes in which they are stable against the thermal--viscous hydrogen ionization instability. Self--gravity and irradiation effects are included. We find that most if not all AGN discs are unstable. Observed AGN therefore represent the outburst state, although some or all quasars could constitute a steady population having markedly higher fuelling rates than other AGN. It has important  implications for the AGN mass supply and for the presence of supermassive black holes in nearby spirals. ",
        "introduction": "Accretion discs are a nearly ubiquitous feature of close binary systems, and their presence is widely invoked in models of active galactic nuclei (AGN). A major feature of the discs in binaries is the thermal--viscous instability driven by hydrogen ionization (Meyer \\& Meyer--Hofmeister 1982, Smak 1982). It is now commonly accepted that this instability drives outbursts in cataclysmic variables and soft X--ray transients. Lin \\& Shields (1986) showed by a local stability analysis that this instability can also operate in accretion discs thought to be present around supermassive black holes in AGN. They concluded that these discs were unstable at radii ($ \\approx 10^{15} - 10^{16}$ cm), where the surface temperature is several thousand degrees. The expected characteristic time scale for this instability is $10^4 - 10^7$ years.  Because of its generic nature, the ionization instability plays a dominant role in characterizing the observed behaviour of the host systems. In the binary context, attempts to understand the precise conditions (mass of accreting object, accretion rate) under which it occurs have been at least partially successful (e.g. Smak, 1982, van Paradijs 1996; King, Kolb, \\& Burderi, 1996; King, Kolb \\& Szuszkiewicz, 1997, and references therein). These studies show that self--irradiation of the disc by the central X--ray source has a determining effect on the disc stability if the accreting object is compact, as in soft X--ray transients (see below). Delineating the stable and unstable disc regions is equally important for AGN. If the instability is present in AGN discs the suppression of central accretion in the quiescent state means that we can identify only the outburst states of unstable systems as AGN. Two important consequences follow: (1) quiescent AGN must appear as quite normal galaxies, and (2) the average mass fuelling rate in many if not all AGN is much lower than implied by their current luminosities. This in turn limits the masses that their central black holes are expected to reach.   The idea of intermittent activity in AGN was already suggested by Shields \\& Wheeler (1978). They noticed that the fuelling problem could  be solved if active nuclei store mass during quiescence and this mass then feeds  the hole for a shorter period of intense activity. The thermal--viscous hydrogen  ionization instability found to operate in AGN accretion discs (Lin \\& Shields, 1986) is capable of triggering such  behaviour. Clarke \\& Shields (1989), Mineshige \\& Shields (1990), Cannizzo \\& Reiff, 1992, Cannizzo (1992) studied the full range of black hole masses and accretion rates in order to determine the observational consequence of the instability for the AGN population. Siemiginowska \\& Elvis (1997) attempted to reproduce the observed luminosity function, assuming that this mechanism operates in all AGN.  Our aim here is to decide if the ionization instability still operates in AGN when irradiation effects are included. As we have seen irradiation is central to the discussion of disc stability in soft X--ray transients. Further, irradiation is often thought to dominate the disc emission (e.g. Collin--Souffrin, 1994). For both reasons it is vital to include it in any attempt to decide the disc stability. The actual form of the instability when irradiation is included is outside the scope of our paper. Siemiginowska, Czerny \\& Kostyunin (1996) have performed studies for particular black hole masses and accretion rates, with assumed forms of irradiation.    There is a simple criterion for the instability to appear: the disc must contain regions with effective temperature $T_{\\rm eff}$ close to the value $T_{\\rm H}$ at which hydrogen is ionized. In practice $T_{\\rm H}$ depends on the density and may be quite different in different environments; we shall consider a range of values in this paper. However the criterion is not easy to use in this form, as one does not in general know the radial distribution of the accretion rate, and thus the run of $T_{\\rm eff}$, in a time--varying disc. Accordingly one usually uses the criterion in an indirect form: a disc with a given constant accretion rate $\\dot M$ is self--consistently steady only if $T_{\\rm eff} > T_{\\rm H}$ throughout it. If the criterion fails we may expect outbursts, although the precise nature of these will depend for example on the detailed behaviour of the disc viscosity.  This version of the stability criterion is easy to apply. Since $T_{\\rm eff}$ always decreases with disc radius $R$ in a steady disc, the condition is most stringent at the outer disc radius $R_{\\rm out}$, so we need apply it only there. If the disc's only source of energy is local viscous dissipation we have \\begin{equation} \\left[ T_{\\rm eff}(R) \\right] ^4 = {3GM\\dot{M}\\over 8\\pi R^3\\sigma}f , \\label{eqa} \\end{equation} (e.g. Frank, King \\& Raine (1992); all symbols are explained after equations (2--9)), and the criterion is simply $T_{\\rm eff}(R_{\\rm out}) > T_{\\rm H}$. In a binary system we can estimate $R_{\\rm out}$ with reasonable accuracy as 70\\% of the Roche lobe radius of the accreting star, and the problem is now well determined. Using this approach, Smak (1982) successfully divided outbursting cataclysmic variables (dwarf novae) from the persistent systems (novalikes). The extension to low--mass X--ray binaries is complicated by the fact that the dominant heat source for the disc is not local viscous dissipation (equation  (\\ref{eqa})), but irradiation by the central X--rays. The instability is similarly suppressed if the disc surface temperature given by irradiation exceeds $T_{\\rm H}$ (Tuchman, Mineshige \\& Wheeler, 1990). Provided that due account is taken of this, one can again successfully divide the outbursting systems (soft X--ray transients) from the persistent systems (van Paradijs, 1996; King, Kolb \\& Szuszkiewicz, 1997). The key feature, as in the unirradiated case, is that the edge temperature of the disc can be simply expressed in terms of $\\dot M$ and $R_{\\rm out}$, without any need to solve for the full internal disc structure. In both the CV and LMXB cases there are important consequences for the study of the binary evolution (e.g. King, Kolb \\& Burderi, 1996), which gives a connection between $\\dot M, M$ and $R_{\\rm out}$.  The extension of this approach to AGN is more complicated; here the outer edge of the disc is no longer determined by the simple Roche lobe condition which holds in binaries, but by the requirement that the disc becomes locally self--gravitating (see equation  (\\ref{13}) below). This condition requires a knowledge of the disc density at the outer edge, so we are now required to solve the full global structure of the steady disc to find $R_{\\rm out}$. Thus we examine all possible, stationary, optically thick, geometrically thin disc models relevant for AGN. If these correspond to stable states, AGN discs may be globally steady, and require fuelling at the currently inferred central accretion rates. If not, they will be the outburst states, and the required fuelling rates will be lower than the current central accretion rate. ",
        "conclusions": "Our aim in this study was to investigate whether the thermal--viscous ionization instability operates in AGN  in the presence of irradiation. We have studied stationary, optically thick, geometrically thin discs, in the range of accretion rates and central black hole masses for which these models are self--consistent. It is worth mentioning here that advection dominated optically thin discs can in principle coexist in some particular regions of the parameter space, but which type of the solution will be actually chosen in nature is still an open question. We used a very simple analytic criterion to determine the stability of each model; if the disc is hot enough for  hydrogen to be completely ionized everywhere all way out till its self--gravity radius the ionization  instability cannot operate. We identify such hot regions in the relevant parts  of $\\dot{M}$ -- $M$ plane and show them as  grey (for $\\alpha$--discs) and hatched (for $\\beta$--discs) areas in Figure 2. Unlike other authors (Clarke \\& Shields, 1989; Mineshige \\& Shields, 1990, Cannizzo,  1992) we consider only the upper stable branch of the whole cycle, where our method is appropriate. A major advantage of our approach is that  we do not need a complicated discussion of the  limit cycle. This method  proved successful in  similar studies of accretion discs in X--ray binaries. We gave careful consideration to the opacities used in  our calculations. There are only small differences between results using opacities from Mazzitelli (1989) and Cox \\& Stewart (1970). However differences appear when using  simple fitting  formulae such as  (10) instead of (11) (see Figure 1): it is important to check carefully that a particular fit found in the literature is appropriate for the range of temperatures and densities used in a given  problem.  Another result of our study is that for $\\alpha \\gta 0.003$ the region between region $a$ and $c$ differs from the standard Shakura--Sunyaev region $b$. We denote it $b^*$. It is radiation pressure dominated, but the main source of opacity is true absorption. We have confirmed the existence of this region in numerical calculations of global disc structure performed using the Cox \\& Stewart (1970) opacity tables. It is interesting that $b^*$ is stable against disc instabilities  triggered by  radiation pressure (Pringle, 1976): while irradiated  it might significantly change its properties.  In Figure 3 we compare our results with those based on detailed studies of the outburst cycle over the parameter space considered by various authors. The dotted lines are from Mineshige \\& Shields (1990), dotted--dashed from Clarke \\& Shields (1989), long--dashed from Cannizzo (1992) and the bold lines from  this paper. The short--dashed line gives the Eddington limit. Our results for non--irradiated disc are in good agreement with those obtained previously. Our main result, quite contrary to the case of close binaries, is that irradiation does  not change the  borders between unstable and stable (partially or completely ionized) regions. In other words, irradiation by a central point source is unable to stabilize the whole disc out to its self--gravity radius. An important reason for this is that one of the effects  of such irradiation is to move the self--gravity radius even farther out from the central black hole. The irradiated disc structure for low--luminosity, low--mass objects differs from that of the  equivalent discs without irradiation (regions C$^+$ in Figure 2). Thus the actual appearance of the ionization instability might well  be affected. This can be studied only by detailed calculations of thermal limit cycles in the presence of irradiation.  We see from Figure 2 that in general AGN discs will be subject to the ionization instability, even if they are irradiated by a central point source. For typical AGN luminosities, corresponding to central accretion rates $\\lta 10^{-2}M_{\\odot}$~yr$^{-1}$, we see from Figure 2 that it is inconsistent to assume that the disc is stable. Since central accretion (and thus e.g. X--ray emission) is suppressed in the quiescent state, all observed AGN must presumably be identified as such only in their outburst states (which last $\\gta 10^{3}$ yr). Thus AGN currently observed to have central accretion rates below the stability limits $\\sim 10^{-1} - 10^{-2}M_{\\odot}$ yr$^{-1}$ shown in Figure 2 must actually have considerably lower fuelling rates. Even rather brighter observed AGN need not be steady systems, but may simply represent the outburst states of unstable disc with fuelling rates below the stability limits.  As pointed out in the Introduction, if most AGN discs are unstable, then in quiescence these systems must be indistinguishable from normal galaxies. Moreover the mass fuelling rates needed to power AGN must be much lower than implied by their current luminosities. If the duty cycle for the outburst can be made short enough ($\\lta 10^{-2}$), no fuelling rates greater than about $10^{-2}\\ M_{\\odot}$~yr$^{-1}$ would be needed in AGN. This would also remove the problem that the remnant black holes are predicted to have excessively high masses if accretion is continuous (Cavaliere \\& Padovani 1988).   Alternatively, since most quasars have observed central accretion rates above the stability limits in Figure 2, some or all of them could have steady discs. This group would then form a separate class with much higher fuelling rates $\\dot M \\sim 0.1 - 1\\ M_{\\odot}$~yr$^{-1}$. It is not easy to decide between these possibilities by looking at detailed properties of the individual systems, as outbursting discs rapidly take on a quasi--steady surface density profile (cf Cannizzo, 1993: this property is well known in the context of cataclysmic variables, where the persistent systems -- novalike variables -- look like dwarf novae in permanent outburst). A complicating feature is that many of the objects with high steady fuelling rates would be subject to the radiation--pressure (Lightman--Eardley) instability.  We conclude that many (if not all) AGN represent the outburst state of a thermal--viscous disc instability. We should then consider candidates for the quiescent state. It is tempting to suggest that this may comprise most or all ``normal'' spirals. Galaxies such as our own could therefore harbour moderately massive ($10^6 -10^8M_{\\odot}$) black holes in their nuclei.   {\\bf Acknowledgements} We thank Ulrich Kolb for valuable discussions. This work is supported by a PPARC Rolling Grant for theoretical astrophysics to the University of Leicester. ARK gratefully acknowledges the support of a PPARC Senior Fellowship.  \\clearpage"
    },
    "9803/astro-ph9803131_arXiv.txt": {
        "abstract": "The Solar Diameter Monitor measured the duration of solar meridian transits during the 6 years 1981 to 1987, spanning the declining half of solar cycle 21. We have combined these photoelectric measurements with models of the solar limb-darkening function, deriving a mean value for the solar near-equatorial radius of 695.508 $\\pm$ .026 Mm. Annual averages of the radius are identical within the measurement error of $\\pm$ .037 Mm. ",
        "introduction": "The Sun is the only star for which reasonably precise values of the mass, surface radius and luminosity are known. The solar mass $\\Msun$ is known from planetary motion, with accuracy limited only by the uncertainty in the gravitational constant $G$. The solar radius can in principle be obtained from direct optical measurement of the solar angular diameter, given the very accurate determinations of the mean distance between the Earth and the Sun. In solar modeling, the value $\\Rsun = 695.99 \\Mm$ (Allen 1973) has been commonly used. The models are calibrated to this photospheric radius, in the present paper defined by the point in the atmosphere where the temperature equals the effective temperature, by adjusting some measure of the convective efficacy, such as the mixing length.  Recent accurate observations of solar f-mode frequencies from the SOI/MDI instrument on the SOHO satellite (e.g. Kosovichev {\\etal} 1997) have raised some doubts over this value of $\\Rsun$. The frequencies of these modes are predominantly determined by $G \\Msun/\\Rsun^3$. By comparing the observed frequencies with frequencies of solar models calibrated to $\\Rsun = 695.99 \\Mm$ Schou {\\etal} (1997) and Antia (1998) concluded that the actual solar radius was smaller by about $0.3 \\Mm$ than the assumed radius of the model. Other aspects of the modeling of the solar f modes may affect their frequencies at this level (e.g. Campbell \\& Roberts 1989; Murawski \\& Roberts 1993; Ghosh, Antia \\& Chitre 1995). Thus it is obviously important to obtain independent verification of the proposed correction to the solar radius.  There are indeed significant uncertainties associated with the currently adopted radius value. These are related to the problem of the definition of the solar limb adopted in the radius determinations, and the reduction of the measured value to the photosphere. It is not clear how the value quoted by Allen (1973) was obtained. However, it appears that the more recent determinations, which are generally consistent with Allen, in most cases refer to the inflection point of the solar limb intensity. According to solar atmospheric models this corresponds to a height of about $0.3 \\Mm$ above the photosphere, thus perhaps accounting for the radius correction inferred from the f-mode frequencies.  The uncertainty in the precise definition of the measured values of the solar radius highlights the need to combine the observations with careful modeling of the quantity that is observed. Here we consider a long series of observations obtained with the High Altitude Observatory's Solar Diameter Monitor (Brown {\\etal} 1982). This is based on a definition of the solar limb which minimizes the effect of seeing (Hill, Stebbins \\& Oleson 1975). By combining daily data obtained over more than 6 years, extending between solar maximum and solar minimum, the possible effects of solar activity can be checked. The analysis of the data is carried out by means of a model of the solar limb intensity, following as closely as possible the actual procedure used in the reduction of the data and testing for the effects of seeing. In this way we have eliminated several of the uncertainties affecting earlier determinations to arrive at what we believe to be an accurate measure of the solar photospheric radius. ",
        "conclusions": "We adopted the modified IAU (1976) value of 1.4959787066 $\\times 10^5$ Mm (Astronomical Almanac, 1997) for the astronomical unit, and adjusted this value by -4.678 Mm to account for the mean displacement between the telescope's noontime location and the Earth's center, and by +0.449 Mm for the displacement of the Sun's center relative to the barycenter of the Earth-Sun system. This distance, combined with $D_0$ from Eq. (2), yields the Sun's apparent radius. Applying the model corrections described in the last section, we obtain $$ \\Rsun = (695.5260 \\pm 0.0065 ) \\Mm\\qquad \\hbox{\\rm for Model 1} $$ $$ \\Rsun = (695.4892 \\pm 0.0065 ) \\Mm\\qquad \\hbox{\\rm for Model 2} $$  We estimate the modeling errors to be $1/\\sqrt 2$ of the difference between these estimates, or about 0.020 Mm. Based on the uncertainties in the geometric corrections that were made to the measured radius, we estimate the systematic errors in the measured value to be 0.015 Mm, or about twice as large as the random errors. Averaging our results for Models 1 and 2, and adding the various error sources in quadrature, we arrive at our final estimate of $$ \\Rsun = (695.508 \\pm 0.026 ) \\Mm $$  The inferred solar photospheric radius is smaller by about $0.5 \\Mm$ than the normally used value of $695.99 \\Mm$ (Allen 1976). A review of recent observations was given by Schou {\\etal} (1997), concluding that these were consistent with an angular diameter of $1919.26\\arcsec \\pm 0.2\\arcsec$, corresponding to Allen's value of $\\Rsun$. This is also consistent with the observed value obtained here (cf. eq. 2). However, it appears that the observations considered by Schou {\\etal} refer to the inflection point of intensity (or, in one case, to an FFTD determination) and hence do not contain the correction to photospheric radius. Such a correction, taking into account the observational characteristics, is an essential part of the radius determination.  Some confirmation of the reliability of the modeling comes from the comparison in Fig. 2 of computed and observed slopes of the limb position as function of the scan widths. Nevertheless, it is striking that, as indicated by the difference between Models 1 and 2, the major uncertainty in $\\Rsun$ appears to come from the modeling. Indeed, it is evident that the real solar atmosphere is substantially more complicated than the one-dimensional model resulting from the ATLAS code or the mean model obtained from the hydrodynamical simulations. A more accurate determination of the radius correction can probably be obtained from a detailed calculation of the limb intensity, taking into account the inhomogeneous nature of the relevant layers, on the basis of the simulations. Such an investigation is beyond the scope of the present paper, however.  We find no significant variation in the observed diameter during the observation period (cf. Fig. 2); annual averages of the radius for the years 1981 to 1987 all agree within their measurement errors of $\\pm .037$ Mm. These limits are substantially smaller than diameter changes reported previously for the same interval of time (e.g. Ulrich \\& Bertello 1995, Laclare {\\it et al.} 1996), but are in agreement with measurements by Wittman (1997). On the other hand, the limb-position slope shows fairly substantial variations. We also note that during solar maximum, the daily slope values tended to be highly variable as well as small in magnitude; this suggests that the long-term variation may result from localized activity-dependent features such as faculae. It is plausible that the previously inferred variations in solar diameter with solar activity is in fact a reflection of such variations in the limb-darkening slope.  It is interesting that the value of $\\Rsun$ obtained here is somewhat smaller than that inferred from the solar f-mode frequencies, indicating additional contributions to the differences between the observed and model values of these frequencies. This issue, and the effects of the reduction of the model radius on the helioseismically determined structure of the solar interior will be considered elsewhere. We note, however, that Antia (1998) and Schou {\\etal} (1997) found significant effects on the helioseismically inferred sound speed from corresponding radius changes."
    },
    "9803/astro-ph9803307_arXiv.txt": {
        "abstract": "I review the status of observational determinations of central masses in nearby galactic nuclei.  Results from a variety of techniques are summarized, including ground-based and space-based optical spectroscopy, radio VLBI measurements of luminous water vapor masers, and variability monitoring studies of active galactic nuclei.  I will also discuss recent X-ray observations that indicate relativistic motions arising from the accretion disks of active nuclei.  The existing evidence suggests that supermassive black holes are an integral component of galactic structure, at least in elliptical and bulge-dominated galaxies.  The black hole mass appears to be correlated with the mass of the spheroidal component of the host galaxy.  This finding may have important implications for many astrophysical issues. ",
        "introduction": "The discovery of quasars in the early 1960's quickly spurred the idea that these amazingly powerful sources derive their energy from accretion of matter onto a compact, extremely massive object, most likely a supermassive black hole (SMBH; Zel'dovich \\& Novikov 1964; Salpeter 1964; Lynden-Bell 1969) with $M\\,\\approx\\,10^6-10^9$ \\solmass.  Since then this model has provided a highly useful framework for the study of quasars, or more generally, of the active galactic nucleus (AGN) phenomenon (Rees 1984; Blandford \\& Rees 1992).  Yet, despite its success, there is little empirical basis for believing that this model is correct.  As pointed out by Kormendy \\& Richstone (1995, hereafter KR), our confidence that SMBHs must power AGNs largely rests on the implausibility of alternative explanations.  To be sure, a number of characteristics of AGNs indicate that the central engine must be tiny and that relativistic motions are present.  These include rapid X-ray variability, VLBI radio cores, and superluminal motion.  However, solid evidence for the existence of SMBHs in the centers of galaxies has, until quite recently, been lacking.  As demonstrated by Soltan (1982), simple considerations of the quasar number counts and standard assumptions about the efficiency of energy generation by accretion allows one to estimate the mean mass density of SMBHs in the universe.  The updated analysis of Chokshi \\& Turner (1992) finds $\\rho_{\\bullet}\\,\\approx\\,2 \\times 10^5 \\epsilon_{0.1}^{-1}$ \\solmass\\ Mpc$^{-3}$ for a radiative efficiency of $\\epsilon\\,=\\,0.1\\epsilon_{0.1}$.  Comparison of $\\rho_{\\bullet}$ with the $B$-band galaxy luminosity density of 1.4\\e{8}$h$ \\solum\\ Mpc$^{-3}$ (Lin \\etal 1996), where the Hubble constant $H_{\\rm 0}$ = 100$h$ \\kms\\ Mpc$^{-1}$, implies an average SMBH mass per unit stellar luminosity of $\\sim$1.4\\e{-3}$\\epsilon_{0.1}^{-1} h^{-1}$ \\solmass/\\solum.  A typical bright galaxy with $L_B^*\\,\\approx\\,10^{10} h^{-2}$ \\solum\\ potentially harbors a SMBH with a mass \\gax $10^7\\epsilon_{0.1}^{-1}h^{-3}$ \\solmass.  These very general arguments lead one to conclude that ``dead'' quasars ought to be lurking in the centers of many nearby luminous galaxies.  The hunt for SMBHs has been frustrated by two principal limitations.  The more obvious of these can be easily appreciated by nothing that the ``sphere of influence'' of the hole extends to $r_{\\rm h}\\,\\simeq\\,G M_{\\bullet}/ \\sigma^2$ (Peebles 1972; Bahcall \\& Wolf 1976), where $G$ is the gravitational constant and $\\sigma$ is the velocity dispersion of the stars in the bulge, or, for a distance of $D$, $\\sim$1\\asec ($M_{\\bullet}/2\\times10^8$ \\solmass)($\\sigma$/200 \\kms)$^{-2}$($D$/5 Mpc). Typical ground-based observations are therefore severely hampered by atmospheric seeing, and only the heftiest dark masses in the closest galaxies can be detected.  The situation in the last few years has improved dramatically with the advent of the {\\it Hubble Space Telescope (HST)} and radio VLBI techniques.  The more subtle complication involves the actual modeling of the stellar kinematics data, and in this area much progress has also been made recently as well.  Here I will highlight some of the observational efforts during the past two decades in searching for SMBHs, concentrating on the recent advances. Since this contribution is the only one that discusses nuclear BHs aside from that in the Milky Way (Ozernoy, these proceedings) and in NGC 4258 (Miyoshi, these proceedings), I will attempt to be as comprehensive as possible, although no claim to completeness is made, as this is a vast subject and progress is being made at a dizzying pace.  To fill in the gaps, I refer the reader to several other recent review papers, each of which has a slightly different emphasis (KR; Rees 1998; Richstone 1998; Ford \\etal 1998; van der Marel 1998). ",
        "conclusions": ""
    },
    "9803/astro-ph9803288_arXiv.txt": {
        "abstract": "We consider the formation of low--mass X--ray binaries containing accreting neutron stars via the helium--star supernova channel.  The predicted relative number of short--period transients provides a sensitive test of the input physics in this process.  We investigate the effect of varying mean kick velocities, orbital angular momentum loss efficiencies, and common envelope ejection efficiencies on the subpopulation of short--period systems, both transient and persistent. Guided by the thermal--viscous disk instability model in irradiation--dominated disks, we posit that short--period transients have donors close to the end of core--hydrogen burning. We find that with increasing mean kick velocity the overall short-period fraction, $s$, grows, while the fraction, $r$, of systems with evolved donors among short-period systems drops.  This effect, acting in opposite directions on these two fractions, allows us to constrain models of LMXB formation through comparison with observational estimates of $s$ and $r$.  Without fine tuning or extreme assumptions about evolutionary parameters, consistency between models and current observations is achieved for a regime of intermediate average kick magnitudes of about 100--200\\,km\\,s$^{-1}$, provided that (i)~orbital braking for systems with donor masses in the range $1-1.5\\,\\msun$ is weak, i.e., much less effective than a simple extrapolation of standard magnetic braking beyond $1.0\\,\\msun$ would suggest, and (ii)~the efficiency of common envelope ejection is low. ",
        "introduction": "The structure and properties of accretion disks around non-magnetic compact objects in close binaries are primarily determined by the rate at which matter is supplied by the Roche--lobe filling donor star. If this rate is smaller than a critical value $\\dmc$, the disk is subject to a thermal--viscous instability and undergoes a limit cycle evolution alternating between a hot and a cool state, the outburst and quiescent phases. This model has been successfully applied to dwarf nova outbursts in cataclysmic variables (CVs), where the accretor is a white dwarf (WD) (for reviews see, e.g., Cannizzo\\markcite{C93} 1993; Osaki\\markcite{O96} 1996), and also to X--ray transient outbursts in low-mass X-ray binaries (LMXBs), where the accretor is a neutron star or a black hole (for reviews see, e.g., Lasota\\markcite{L96} 1996; Wheeler\\markcite{W98} 1998).  In the case of LMXBs, where the accretion efficiency and hence the X-ray luminosity from the compact object is higher than in CVs, there is strong evidence that self--irradiation is important (see, e.g., van~Paradijs \\& McClintock\\markcite{vP94} 1994; King \\& Ritter\\markcite{KR98} 1998).  X--ray irradiation from the central accretor raises the disk temperature for a given mass transfer rate, thereby suppressing the instability.  As a consequence, $\\dmc$ is significantly smaller in LMXBs than in CVs, consistent with observations (van~Paradijs\\markcite{vP96} 1996). The strength of the irradiation in LMXBs depends on the nature of the compact object. In the neutron--star case, the irradiating source is equivalent to a point source at the center of the disk, while, in the black--hole case, the lack of a hard stellar surface implies that the irradiating source is only the innermost disk, which is weaker by a factor about equal to the relative disk thickness, the other system parameters being constant (King, Kolb, \\& Szuskiewicz\\markcite{KK97} 1997).  Disk irradiation has important consequences for short--period LMXBs. These systems evolve towards shorter orbital periods $P$ under the influence of orbital angular momentum losses $\\dot J$ (henceforth ``j--driven'' systems), caused by a magnetic stellar wind from the donor star (magnetic braking) or by gravitational radiation. Assuming that $\\dot{J}_{\\rm MB}$ from magnetic braking (MB) is independent of the nature of the accretor, the same orbital braking formalism must apply for both CVs and LMXBs.  For a given MB law, i.e., for a given dependence of $\\dot J$ on binary parameters, the mass transfer rate (proportional to the fractional angular momentum loss rate) is smaller in black--hole binaries than in neutron--star binaries, as the total angular momentum increases with primary mass. This, together with the higher values for $\\dmc$, leads to transient behavior for all j--driven black--hole binaries, as observed. The converse appears to be true for neutron--star systems. The MB transfer rate may be up to two orders of magnitude larger than $\\dmc$ (King, Kolb, \\& Burderi\\markcite{KK96} 1996), suggesting that there are no or only very few neutron star transients.  However, it is clear from observations that the fraction of transients among short--period neutron--star LMXBs is non--negligible. Five out of 23 LMXBs with a confirmed or possible neutron star primary and with periods in the range $3\\,{\\rm h} < P < 20\\,{\\rm h}$ are classified as transient (Ritter \\& Kolb\\markcite{R98} 1998).  Although it is difficult to estimate the {\\em intrinsic} value of this fraction, it may be even larger than $5/23$ (see \\S\\,5.1 below).    A possible resolution of this apparent contradiction between the disk-instability model and the observations has been suggested by King, Kolb, \\& Burderi\\markcite{KK96} (1996) and King \\& Kolb\\markcite{K97} (1997, hereafter KK97). The mass transfer rates in j--driven LMXBs\\footnote{Hereafter, we will refer to neutron--star LMXBs simply as LMXBs, unless otherwise stated.} with somewhat evolved donor stars (close to the end of core--hydrogen burning), are significantly smaller than in systems with unevolved donors, therefore favoring transient behavior. Then the large observed transient fraction demands that the contribution of these evolved systems to the total population is correspondingly large. Under the assumption of spherically symmetric supernovae (SN) and for strong magnetic braking at high ($\\gtrsim 1.2\\,{\\rm M}_\\odot$) donor star masses, KK97 showed that LMXBs forming via the standard evolutionary channel involving a helium star supernova (Sutantyo\\markcite{S75} 1975; van den Heuvel\\markcite{vdH83} 1983) do indeed have this property. However, for weaker magnetic braking (constrained by rotational velocity data for F stars), Kalogera \\& Webbink\\markcite{K98} (1998; hereafter KW98) showed that j--driven LMXBs form only if supernovae are asymmetric. KK97 also showed, in a qualitative way, that asymmetric SNe with on average large kick velocities imparted to the neutron stars would inevitably lead to a large number of unevolved systems in the population, and hence a small number of transient systems among the short-period LMXB population, contrary to observations.  This qualitative conclusion by KK97 appears to be in contradiction to numerous pieces of evidence for the existence of rather substantial kick velocities (Kaspi et al.\\ \\markcite{Kp96}1996; Hansen \\& Phinney\\markcite{H97} 1997; Lorimer, Bailes, \\& Harrison\\markcite{L97} 1997; Fryer \\& Kalogera\\markcite{F97} 1997; Fryer, Burrows, \\& Benz\\markcite{F98} 1998), so the need for a quantitative study of the problem arises. The formation of LMXBs including the effect of natal neutron--star kicks has been studied in detail by KW98.  Here, we use these detailed synthesis models to address the issue of the transient fraction among LMXBs and investigate its dependence on different evolutionary parameters. The increase of this fraction in the model populations is accompanied by an increase of the fraction of systems with donors that have evolved beyond the main sequence, i.e., of long--period systems driven by nuclear expansion of the secondary (``n--driven'' systems).  Our goal is twofold: first, to disentangle the differential dependences on model parameters of the two observable quantities, the predicted fraction of ``j--driven'' systems among all LMXBs and the predicted transient fraction among ``j--driven'' LMXBs. Second, to constrain quantitatively the input parameters by requiring that both predicted fractions are consistent with current observations. These two observational constraints operate in opposite directions, and this allows us to derive limits on the mean magnitude of natal kicks imparted to neutron stars, the strength of magnetic braking, and the efficiency with which orbital energy is consumed during the common envelope phase prior to the supernova.  As the observational sample increases in the future, the model calculations presented here can be used to tighten these constraints still further.  In \\S\\,2, we review the evolutionary picture of j--driven LMXBs and derive a simplified criterion for transient behavior. In \\S\\,3, we describe the different models for LMXB formation considered here, while the results and basic effects of varying model parameters are presented in \\S\\,4. In \\S\\,5, we first (\\S\\,5.1) describe the observed sample and derive the observational constraints, and then (\\S\\,5.2) evaluate the models based on these constraints. We conclude in \\S\\,6 with a discussion of our results and of possible extensions of the present study. ",
        "conclusions": "We applied the population synthesis techniques developed by KW98 to investigate the influence of varying mean supernova--kick magnitudes, orbital angular momentum loss strengths, and common envelope efficiencies on the population of transient and persistent short--period LMXBs that form via the helium star supernova channel. A main premise of our study is the identification of short--period transients with systems where the donor is significantly nuclear--evolved, close to the end of core--hydrogen burning. We tested our models against two properties inferred from the observed sample: that a significant fraction ($s \\ga 50\\%$) of nascent LMXBs are short--period, and that a significant fraction of these ($r \\ga 20\\%$) have donors close to the end of core--hydrogen burning.   We found that any combination of model parameters that results in a large fraction of short-period systems with evolved donors, hence high values of $r$, at the same time results in the formation of many systems with donors that have evolved beyond the main sequence, and hence leads to low values of $s$. It is exactly these two counteracting effects that allow us to constrain the three model parameters.  With an increasing mean kick magnitude $s$ grows, while $r$ drops. Model predictions for $s$ and $r$ are consistent with observational estimates in an intermediate regime of moderate mean kick magnitudes, $\\vmean \\simeq 100-200$~km/s, if (i)~the orbital braking for systems with donor masses $1\\la M_d \\la 1.5\\,\\msun$ is weak, i.e., much less effective than a simple extrapolation of magnetic braking beyond $1\\,\\msun$ would suggest, and (ii)~the efficiency of common envelope ejection is low ($\\alpha_{\\rm CE} \\la 0.5$).  Consistency with observational estimates could also be achieved in the absence of kicks or for a very large mean kick velocity, but {\\em only} in combination with a fine-tuned common-envelope efficiency and/or fairly extreme assumptions on the efficiency of magnetic braking at masses $\\ga 1\\,\\msun$ (very high for small kicks, almost negligible for large kicks).  However, in view of the various uncertainties that enter our study we cannot unambiguously rule out these models.  Overall, we have shown that the model predictions are sufficiently sensitive to the variation of the three input parameters to place constraints on them when comparing to observations.  Given accurate input from the observational sample, the constraints can be reasonably tight (for example, had the adopted values of $s$ and $r$ been known within a factor of 2, we could unambiguously exclude a single peak of kick magnitudes at $\\lesssim 50$ or $\\gtrsim 300$\\,km\\,s$^{-1}$). Nevertheless, significant uncertainties exist that have their origin mainly in two areas:  secular evolution of j--driven LMXBs and observational selection effects, which prevent us from accurately inferring the intrinsic properties of the LMXB population from the observed sample.  In the first area, a systematic study of j--driven LMXB evolution with detailed stellar models is needed to relate transient behavior quantitatively to the degree of evolution of the donor for different orbital angular momentum loss rates. This will also allow one to examine in detail the transition regime between j--driven and n--driven systems and to quantify the relative lifetimes of persistent and transient systems.  The efficiency of magnetic braking affects both the lifetime of the systems and the critical degree of evolution for transient behavior.  An increasing braking strength restricts transient behavior to donors closer to the end of core--hydrogen burning. At the same time, it leads to higher mass transfer rates in persistent systems, and hence shorter relative lifetimes. The first effect increases the critical degree of evolution, $f_0$, whereas the second reduces the intrinsic transient fraction, $r_0$, derived from the observed sample. Therefore, the consistency criterion, $\\rf > r_0$, moves along a typical curve $\\rf$ (Fig.\\,\\ref{r1} and \\ref{r2}), so that the choice of successful or unsuccessful models is not, to zeroth order, affected by the neglect of these dependences. We note that another long--standing problem in the evolution of LMXBs, the fate of j--driven systems once they approach to orbital periods of 3\\,h, and the apparent lack of systems with $P \\la 3$\\,h (see KK97 for a discussion), does not affect our study as long as any comparison is restricted to systems with periods longer than 3\\,h.  The Skumanich--type magnetic braking formulation adopted here is by no means the only possible parametrization of the orbital angular momentum losses. Constraints on the real functional form could come from further analysis of stellar rotational rates in open clusters of different age. Some studies (e.g., Krishnamurthi et al.\\markcite{K97} 1997, and references therein) suggest that $\\dot J$ becomes a less steep function of $\\omega$ above a certain critical angular velocity. Repeating our model calculations with such more general magnetic braking laws seems worthwhile only once the critical degree of evolution, $f_{\\rm 0}$, for transient behavior can be estimated more quantitatively.  Progress in the second area of uncertainty requires a systematic study of observational selection effects on X--ray binaries (given the X-ray instruments used so far).  This is hampered by the difficulty of calculating the spectral distribution of the emergent X--ray flux in a system with given period and transfer rate.  The assessment of the relative completeness of the different subgroups involved (black--hole, neutron--star, transient, persistent, short--period, and long--period systems) will certainly gain reliability when future observations increase the known sample and the number of systems with determined binary and transient parameters.  The completeness is further affected by a possible dependence of the transient outburst recurrence time on orbital period. Giant donor systems have larger disks, which take longer to reestablish the critical pre--outburst surface density. If the recurrence time is systematically longer in long--period systems, the short--period fraction could be significantly smaller than estimated. This in turn would make models with even smaller mean kick velocities acceptable.  Despite these uncertainties, the preference for moderate mean kick velocities of $\\simeq 100-200$~km/s inferred from our study is likely to persist even for a more detailed treatment of secular evolution and a better understanding of selection effects.  This preference seems to be in conflict with the fairly large {\\em mean} natal kicks traditionally deduced from pulsar proper motions (e.g., $\\vmean \\simeq 500$~km/s is favored by Lorimer et al.\\ 1997). (Note that our study cannot constrain the fraction of very high ($\\gtrsim 500$\\,km\\,s$^{-1}$) kick magnitudes, because this only affects the absolute LMXB birth-rate normalization.) While this might point to an underlying physical difference in the way supernova explosions of type II and type Ib proceed, there is also the possibility that natal pulsar velocities are significantly smaller than this estimate. Indeed, it has been pointed out that a relatively wide variety of qualitatively different distributions are consistent with the observed pulsar velocity distribution (Hansen \\& Phinney 1997; Cordes \\& Chernoff\\markcite{C98} 1998; Fryer et al.\\ 1998; Hartmann\\markcite{H98} 1998). Although our study was limited to Maxwellian kick distributions we can conclude quite generally that the observed short-period transient LMXB population favors kick distributions with a dominant component at moderate mean velocities of about 100 km/s.   We conclude by noting that with future progress the observational sample of LMXBs will undoubtedly increase and improve in quality, and more detailed calculations of the secular evolution of LMXBs will become available. The model calculations we have presented here, of the fraction of short-period LMXBs and the transient fraction among them, can then be used to derive even tighter constraints on the supernova and evolutionary parameters."
    },
    "9803/astro-ph9803182_arXiv.txt": {
        "abstract": "The pulsar PSR B1259$-$63 is in a highly eccentric 3.4-yr orbit with the Be star SS 2883. Timing observations of this pulsar, made over a 7-yr period using the Parkes 64-m radio telescope, cover two periastron passages, in 1990 August and 1994 January. The timing data cannot be fitted by the normal pulsar and Keplerian binary parameters. A timing solution including a (non-precessing) Keplerian orbit and timing noise (represented as a polynomial of fifth order in time) provide a satisfactory fit to the data. However, because the Be star probably has a significant quadrupole moment, we prefer to interpret the data by a combination of timing noise, dominated by a cubic phase term, and $\\dot\\omega$ and $\\dot x$ terms. We show that the $\\dot\\omega$ and $\\dot x$ are likely to be a result of a precessing orbit caused by the quadrupole moment of the tilted companion star. We further rule out a number of possible physical effects which could contribute to the timing data of PSR B1259$-$63 on a measurable level. ",
        "introduction": "The pulsar PSR B1259$-$63 is a member of a unique binary system. Discovered using the Parkes telescope in a survey of the Galactic plane at 1.5 GHz \\cite{jlm+92a}, it was shown by Johnston et al.~(1992b) to be in a highly eccentric 3.4-yr orbit with a 10th-magnitude Be star, SS 2883. The pulsar period, $P$, is relatively short, 47.8 ms, and the measured period derivative, $\\dot{P}$, gives a pulsar characteristic age, $\\tau_c = P/(2\\dot{P})$, of $3.3 \\times 10^5$ yr and a surface magnetic field of 3.3$\\times 10^{11}$ G.  This therefore is a young system, which may evolve through an accretion phase to form a single or binary millisecond pulsar.  Rapidly spinning neutron stars can only accrete matter if the co-rotation velocity at the Alfv\\'en radius is less than the Keplerian velocity at the same radius \\cite{bv91}. Equality of these velocities defines the `spin-up line'. At present, PSR B1259$-$63 lies well to the left of the spin-up line, so that accretion onto the neutron star is not possible until either the pulsar slows down or the pulsar magnetic field decays.  Timing observations of PSR B1259$-$63, made over a 3.4-yr interval and covering the 1990 August periastron, were reported by Johnston et al.\\ \\shortcite{jml+94}. A phase-connected fit to these data gave parameters for the pulsar and its orbit, and showed that the next periastron would occur on 1994 January 9. This paper also reported optical observations which indicate that the companion star is of spectral type B2e, with a mass of $m_* \\sim 10$ M$_{\\sun}$ and radius $R_* \\sim 6$ R$_{\\sun}$. The break-up velocity at the equator, $v_{{\\rm max}}$, for B2e stars is not very well known; it is estimated to be $\\sim 380$ km s$^{-1}$ by Slettebak et al.\\ (1980) and $\\sim 480$ km s$^{-1}$ by Schmidt-Kaler (1982).  Recent work by Porter (1996) suggests that most Be stars rotate at $\\sim 70$ per cent of the break-up velocity.  From the mass function, a companion mass of 10 M$_{\\sun}$ and a pulsar mass, $m_p$, of 1.4 M$_{\\sun}$ imply an orbital inclination $i=36\\degr$. The orbital eccentricity is very high, 0.87, and the pulsar approaches within 24 $R_*$ of the companion star at periastron, passing through the circumstellar disk. Extensive observations of the pulsar were made at several radio frequencies before and after the 1994 January periastron, in order to probe the circumstellar environment of SS 2883 (Johnston et al.\\ 1996; Melatos, Johnston \\& Melrose 1995).  Observations made between 1990 January and 1994 October were well explained by step changes in the pulsar period at the two periastrons \\cite{mjl+95}, attributed to a propeller-torque spin-down caused by the interaction of the pulsar with the circumstellar matter at the Alfv{\\'e}n radius (Illarionov \\& Sunyaev 1975, King \\& Cominsky 1994, Ghosh 1995).  In this paper we report on additional timing observations made using the Parkes radio telescope over the past two years which, together with the earlier data, give a total timing data span of seven years. We find that the timing solution of Manchester et al. (1995) does not fit the recent data and discuss alternative models and their interpretation. ",
        "conclusions": "At present, the timing observations for the binary pulsar PSR B1259$-$63 span seven years. Because of the gaps in timing observations around the two periastrons and the large timing noise present in this young pulsar, we still are not able to derive a unique timing model to explain the TOAs.  Model 1 is a timing solution including a non-precessing Keplerian orbit and timing noise represented as a polynomial of fifth order in time. This model provides a satisfactory fit to the data. The remaining timing residuals are understood as short-term timing noise similiar to that seen in observations of other young pulsars (cf.\\ Foster et al.\\ 1994).  Equally good results were obtained by Model 2 and 3. Both timing models contain just a $\\ddot P$ term to account for the long-term behaviour of the timing noise, and $\\dot\\omega$ and $\\dot x$, which both are understood to result from a precession of the orbit. This orbital precession can be explained by the classical spin-orbit coupling caused by the quadrupolar nature of the main-sequence star companion. The corresponding advance of periastron is negative and thus the companion should be tilted by more than $30\\degr$ with respect to the orbital plane (See Fig.\\ 6). This can be explained by a birth kick for the pulsar (cf.\\ Kaspi et al.\\ 1996).  Tidal dissipation and frictional drag in the circumstellar matter is shown to be negligible. The influence of the mass loss of the companion is too small to be detectable unless the mass loss is $\\ga 10^{-5} M_\\odot$/yr. At the same time we can exclude a significant orbital period change in the TOAs of PSR B1259--63.  PSR B1259$-$63 should show the largest Einstein delay and largest Shapiro delay of all known binary pulsars. The small change in the longitude of periastron makes it impossible to isolate the Einstein delay. The Shapiro delay peaks sharply around periastron and is so far unobservable.  Again, we stress that the physical parameters given here for the companion star, $\\theta$, $\\Phi$ and $k$, should be understood as one possible explanation for the significant values of $\\dot\\omega$ and $\\dot x$ in Models 2A and 2B. If the long-term behaviour of the timing noise of PSR B1259$-$63 is not fully modelled by a cubic term ($\\ddot P$), it is possible that rather large fractions of these parameters are not explained by a precession of the orbital plane but have their origin in unmodelled timing noise. The parameters here show that, in principle, all of the $\\dot\\omega$ and $\\dot x$ can arise from classical spin-orbit coupling, for Model 2A in particular.  Although Model 1 gives a good fit to the TOAs without making use of the classical spin-orbit coupling, it seems unlikely that the classical spin-orbit coupling is of no importance for this system."
    },
    "9803/astro-ph9803019_arXiv.txt": {
        "abstract": "We present deep near--infrared images of high redshift radio galaxies (HzRGs) obtained with the Near Infrared Camera (NIRC) on the Keck I telescope.  In most cases, the near--IR data sample rest wavelengths free of contamination from strong emission lines and at $\\lambda_{\\rm rest} > 4000$\\AA, where older stellar populations, if present, might dominate the observed flux.  At $z > 3$, the rest--frame optical morphologies generally have faint, large--scale ($\\sim$50 kpc) emission surrounding multiple, $\\sim 10$ \\kpc~size components. The brightest of these components are often aligned with the radio structures.  These morphologies change dramatically at $2 < z < 3$, where the $K$ images show single, compact structures without bright, radio--aligned features.  The linear sizes ($\\sim 10$ \\kpc) and luminosities ($M(B_{\\rm rest}) \\sim -20$ to $-22$) of the {\\it individual} components in the $z > 3$ HzRGs are similar to the {\\it total} sizes and luminosities of normal, radio--quiet, star forming galaxies at $z = 3 - 4$ (Steidel et al.\\ 1996; Lowenthal et al.\\ 1997).  For objects where such data are available, our observations show that the line--free, near--IR colors of the $z > 3$ galaxies are very blue, consistent with models in which recent star formation dominates the observed light. Direct, spectroscopic evidence for massive star formation in one of the $z >3$ HzRGs exists (4C41.17, Dey \\etal 1997$a$).  Our results suggest that the $z > 3$ HzRGs evolve into much more massive systems than the radio--quiet galaxies and that they are qualitatively consistent with models in which massive galaxies form in hierarchical fashion through the merging of smaller star--forming systems.  The presence of relatively luminous sub--components along the radio axes of the $z > 3$ galaxies suggests a causal connection with the AGN.  We compare the radio and near--IR sizes as a function of redshift and suggest that this parameter may be a measure of the degree to which the radio sources have induced star formation in the parent objects.  We also discuss the Hubble diagram of radio galaxies, the possibility of a radio power dependence in the $K - z$ relation, and its implications for radio galaxy formation.  Finally, we present for the first time in published format basic radio and optical information on 3C~257 ($z=2.474$), the highest redshift galaxy in the 3C sample and among the most powerful radio sources known. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803255_arXiv.txt": {
        "abstract": "A linearity test shows $H_0$ to decrease by 7\\% out to $18\\,000\\kms$. The value at $10\\,000\\kms$ is a good approximation to the mean value of $H_0$ over very large scales. The construction of the extragalactic distance scale is discussed. Field galaxies, cluster distances relative to Virgo, and blue supernovae of type Ia yield $H_{0}$\\,(cosmic) with increasing weight; they give consistently $H_{0}=57\\pm7$ (external error). This value is supported by purely physical distance determinations (SZ~effect, gravitational lenses, MWB~fluctuations). Arguments for $H_{0}>70$ are discussed and shown to be flawed. ",
        "introduction": "\\label{sec:1} The calibration of the cosmic expansion rate $H_0$ consists of two steps. The first step is an investigation of the cosmic expansion field. How {\\em linear\\/} is the expansion? How large are {\\em systematic\\/} deviations from linearity in function of distance? What is the {\\em scatter\\/} of individual objects due to peculiar motions about the mean expansion? Only after these questions are solved can the second step be tackled, i.\\,e. the calibration of the expansion rate in absolute terms.  The procedural difference between the two steps is that only redshifts and {\\em relative\\/} distances are needed for an investigation of the characteristics of the expansion field, while the calibration of the present large-scale expansion rate $H_0$ requires in addition the {\\em true\\/} distance of at least one object which demonstratably partakes of the mean expansion.  Much confusion about the expansion rate has arisen from equating the velocity-distance ratio of a subjectively chosen object with $H_0$. The determination of $H_0$ from the Virgo, Fornax, or Coma clusters, for instance, is meaningful only if it is demonstrated that they reflect at their moderate distances the mean cosmic expansion. The long-standing problem of correcting the observed mean velocity of the Virgo cluster into the frame of the cosmic expansion field has become a classic (cf. Section~3.2). The Fornax cluster with a velocity of $v\\approx1200\\kms$ cannot be used for the determination of $H_0$, even if a useful distance was known for it, because its unknown peculiar velocity may well be as high as 20\\% of its observed velocity. And the Coma cluster at $v\\approx7000\\kms$, which is sometimes used for the determination of $H_0$, may still have a peculiar velocity component of 10\\%, as the peculiar velocity of $630\\kms$ with respect to the MWB of one other supercluster, i.\\,e. the Local Supercluster, would suggest.  The present paper outlines this two-step procedure. In Section~2 the available data are used to map the expansion rate in function of distance well beyond $30\\,000\\kms$, i.\\,e. out to distances where the truly cosmic character of $H_0$ cannot be questioned. Section~3 gives a summary of the various methods of determining distances of field galaxies, of the Virgo cluster, and --- most decisively for $H_0$ --- of distant blue SNe\\,Ia. Methods leading seemingly to $H_{0}>70$ are critically discussed in Section~4. A brief outlook is given in Section~5. ",
        "conclusions": "\\label{sec:5} A Test for the variation of $H_0$ with distance suggests a decrease by $\\sim\\!7\\%$ from $1000 < v \\le 18\\,000\\kms$. At $v=10\\,000\\kms$ $H_0$ goes through a value close to the mean over very large scales.  A system of three interconnected distance scales (field galaxies, cluster distances relative to the Virgo cluster, and most significantly blue SNe\\,Ia) give $H_{0}$ (cosmic) $=57\\pm7$ (external error). Physical distance determinations from the SZ effect, gravitationally lensed quasars, and MWB fluctuations scatter about the same value.  A discussion of proposed high values of $H_0$ shows that disagreement focuses on two topics: 1) the true distance of the Virgo cluster, and 2) the appreciation of the Malmquist bias. One may add as item 3) the distance of the E/S0 galaxies in the Fornax cluster; the latter has lower priority because the peculiar motion of this cluster is unknown, and it is poorly tied into the relative distance scale of other clusters.  \\bigskip \\noindent {\\small {\\bf Acknowledgement:} Financial support of the Swiss National Science Foundation is gratefully acknowledged. The author thanks his colleagues in the $HST$ team for the luminosity calibration of SNe\\,Ia, i.\\,e. Dres. A.~Sandage, A.~Saha, L.~Labhardt F.\\,D.~Macchetto, and N. Panagia, as well as the many collaborators behind the scenes at the STScI; much of the present understanding of $H_0$ depends on their work. He also thanks Mr.~Bernd Reindl for his excellent help in all computational and technical matters.}"
    },
    "9803/astro-ph9803313_arXiv.txt": {
        "abstract": "We have obtained medium-resolution spectra of seven UV-bright stars discovered on images of four southern globular clusters obtained with the  Ultraviolet Imaging Telescope (UIT).      Effective temperatures, surface gravities and helium abundances are derived from LTE and non-LTE model atmosphere fits. Three of the stars have sdO spectra,  including M4-Y453 (\\teff\\ = 58800~K, \\logg\\ = 5.15), NGC 6723-III60 (\\teff = 40600~K, \\logg\\ = 4.46) and NGC~6752-B2004 (\\teff\\ = 37000~K, \\logg\\ = 5.25). All seven stars lie along either post-extended horizontal branch (EHB) or post-early AGB evolutionary tracks.     The post-early AGB stars show solar helium abundances, while the post-EHB stars are helium deficient, similar to their EHB progenitors. ",
        "introduction": "Ultraviolet images of globular clusters are often dominated by  one or two hot, luminous, ``UV-bright'' stars.  The most luminous of these stars are believed to be post-asymptotic giant branch (post-AGB) stars, which go through a luminous UV-bright phase as they leave the AGB and move rapidly across the HR diagram toward their final white dwarf state.    Despite their short lifetimes ($\\sim 10^{5}$ yrs), hot post-AGB stars can dominate the total ultraviolet flux of an  old stellar population.    In particular, hot post-AGB stars probably make a  significant (although not the dominant) contribution  to the UV-upturn observed in elliptical galaxies (Brown et al.\\ \\cite{brown97}). However, a large uncertainty exists in modeling the contribution of hot post-AGB stars to the integrated spectrum of an old stellar population, due to the strong dependence of the post-AGB luminosity and lifetime on the core mass, which in turn depends  on when the stars leave the AGB (Charlot et al.\\ \\cite{charris96}).  Also the previous  mass loss on the red giant branch (RGB) plays an important r\\^ole here,  since it determines the fate of a star during and after the horizontal branch  stage: Stars with very low envelope masses settle along the extended  horizontal branch (EHB) and evolve from there directly to the white dwarf  stage,  whereas stars with envelope masses of more than 0.02~\\Msolar  will at least partly ascend the AGB.  A further uncertainty arises because theoretical post-AGB tracks have been only minimally tested for old, low-mass stars.  The last census of hot post-AGB stars in globular clusters was published by de Boer (\\cite{debo87}), but this list  is certainly incomplete.   The detection of hot post-AGB stars in optical  color-magnitude diagrams (CMD's) is limited by selection effects due to crowding in the cluster cores and to the large bolometric corrections for these hot stars.   More complete searches are possible for hot post-AGB stars in planetary nebulae, for example, by using \\ion{O}{III} imaging.   However, only four planetary nebulae (PNe) were discovered in a recent survey of 133  globular clusters (Jacoby et al.\\ \\cite{jamo97}), of which two were previously known (K648 in M~15, and IRAS 18333-2357 in M~22).        Jacoby et al. expected to find 16 planetary nebulae in their  sample, on  the basis of the planetary nebula luminosity function for metal-poor populations.    The origin of this discrepancy is not yet understood, but we mention two possible contributing factors.   First, the \\ion{O}{III} search of Jacoby et al. may have missed some old, faint  planetary nebulae.     Second, Jacoby et al. derive  the number of expected PNe from the total cluster luminosity, assuming  that all stars in a globular cluster will eventually go through the AGB phase.    But in a cluster such as NGC 6752, about 30\\% of the HB population consists of EHB stars  (with \\teff\\ $>$ 20,000~K),  which are predicted to evolve into white dwarfs without ever passing through the thermally pulsing AGB phase. The exact fraction of stars which follow such evolutionary will depend on the poorly known mass loss rates during the HB and early-AGB phases.  While globular clusters with a populous EHB are expected to be deficient in post-AGB stars, they should show a substantial population of less luminous (1.8 $<$ \\logl\\ $<$ 3) UV-bright stars, which can be either  post-EHB  stars or post-early AGB stars. The population of  post-EHB stars  is expected to be about 15--20 \\% of the population of EHB stars (Dorman et al. \\cite{dorm93}). The post-early AGB population arises from hot HB stars with sufficient envelope mass to return to the AGB, but which peel off  the AGB prior to the thermally pulsing phase (Dorman et al. \\cite{dorm93}).  During the two flights of the {\\em ASTRO} observatory in 1990 and 1995, the Ultraviolet Imaging Telescope (UIT, Stecher et al.\\ \\cite{stech97}) was used to obtain ultraviolet ($\\sim 1600$~\\AA) images of 14 globular clusters.   The solar-blind detectors on UIT suppress the cool star population, which allows UV-bright stars to be detected into the cluster cores, and the $40'$ field of view of UIT is large enough to image the entire population of most of the observed clusters.     Thus, the UIT images provide a complete census of the hot UV-bright stars in the observed clusters.   We have begun a program to obtain spectra of all the UV-bright stars found on the UIT images, in order to derive effective temperatures and gravities for the complete sample, for comparison with evolutionary tracks.       Several of the UV-bright stars found on the UIT images, such as ROB 162 in NGC~6397, Barnard 29 in M~13, and  vZ 1128 in M~3, were previously known and are well-studied.   Other UIT stars are too close to the cluster cores for ground-based spectroscopy, and will require HST observations for further study.    In this paper, we report on spectroscopy of those UIT UV-bright stars accessible for ground-based observations from the southern hemisphere. ",
        "conclusions": "The derived effective temperatures and gravities of the target stars are plotted in Figure 3, along with  ZAHB and post-HB evolutionary tracks for [Fe/H] = $-$1.48 from Dorman et al. (\\cite{dorm93}), and post-AGB (0.565~\\Msolar) and post-early AGB (0.546~\\Msolar) tracks from Sch\\\"onberner (\\cite{scho83}).  The stars NGC~6121-Y453, NGC~6723-III60, and NGC~6723-IV9 appear to fit the  post-early AGB track, while the remaining four targets are consistent with post-EHB evolutionary tracks.  In agreement with this scenario, the three post-early AGB stars have approximately solar helium abundances, while the post-EHB stars have subsolar helium abundances.    The latter stars are expected to have subsolar helium abundances because they are direct descendants of EHB stars, which are known to show helium deficiencies (Moehler et al. \\cite{mohe97}), most likely due to diffusion processes.     The post-early AGB stars, on the other hand, have evolved off the AGB,  where the convective atmosphere is expected to eliminate any previous abundance depletions  caused by diffusion. Curiously, no helium-rich (He/H $\\ge$ 1) sdO stars have yet been found in a  globular cluster, although such stars dominate the field sdO population (Lemke et al. \\cite{lehe98}).  As expected, the two clusters with a populous EHB (NGC~2808 and NGC~6752) have post-EHB stars but no post-AGB stars. The clusters NGC~6723 and M~4, on  the other hand,  do not have an EHB population, although they do have stars blueward of the RR Lyrae gap (which are  potential progenitors of post-early AGB stars). The lack of true post-AGB stars may be understood from the different lifetimes: The lifetime of  Sch\\\"onberner's post-early AGB track is about 10 times longer than his lowest mass post-AGB track.  Thus, even if only a small fraction of stars follow  post-early AGB tracks, those stars may be more numerous than true  post-AGB stars. Due to their relatively long lifetime, post-early AGB stars are unlikely to be observed as central stars of planetary nebulae (CSPNe) since any  nebulosity is probably dispersed before the central star is hot enough to  ionize it.   Additional detail on the individual stars is given below:  \\subsection{NGC 2808}  All three stars in NGC~2808 analysed in this paper are likely post-EHB stars. (Unfortunately, the best post-AGB candidate stars on the UIT image of NGC~2808 are too close to the cluster center to allow spectroscopy from the ground.) Although C2946 and C2947 could be separated in the long-slit optical spectra, they are too close together to estimate individual UV fluxes from the UIT image, and thus there is no UV luminosity determination in Table 2. Due to the well-populated EHB of NGC~2808 (Sosin et al., \\cite{sos97}), a large number of post-EHB stars are expected.  From their three-colour WFPC2 photometry of NGC 2808, Sosin et al.\\ (\\cite{sos97}) find a larger distance modulus  [(m-M)$_0$ = 15.25 -- 15.40] and lower reddening [\\Ebv\\ = 0.09 -- 0.16] than the values adopted here from Harris (\\cite{harris96}).    The use of the distance and reddening of Sosin et al.\\ would yield masses about 20\\% larger, and luminosities about 0.05 dex larger than the values given in Table 2.  \\subsection{NGC~6121 (M~4)}  Y453 is possibly the hottest  globular cluster star known so far. Other candidates are three central stars of planetary nebulae (IRAS 18333-2357 in M~22, Harrington \\& Paltoglou  \\cite{hapa93}; JaFu1 in Pal~6 and JaFu2 in NGC~6441, Jacoby et al., \\cite{jamo97}), which however lack model atmosphere analyses of their stellar spectra. Such a high \\teff\\ is not unexpected,  since according to the Sch\\\"onberner tracks, a post-AGB star will spend most of its  lifetime at temperatures greater than 30,000~K. However, as pointed out by  Renzini (\\cite{renz85}), the large bolometric corrections of such hot stars  in the visible have biased the discovery of post-AGB stars in favour of  cooler stars.  In the \\logt --\\logg\\ plot, Y453 fits well on the 0.546~\\Msolar\\ post-early AGB track of Sch\\\"onberner (\\cite{scho83}).  However, our derived luminosity (\\logl\\ = 2.6) is considerably lower than the Sch\\\"onberner track at that \\teff\\ and \\logg, and the derived mass of 0.16~\\Msolar\\ is astrophysically implausible. In order to obtain a mass of 0.55~\\Msolar, the value of \\logg\\ would need to be 5.68 instead of 5.15. This difference is too large to be accommodated by the spectral fitting, and still would not explain the discrepancy with the  theoretically expected luminosity.  Therefore, below we consider some other possible sources of error: \\begin{description} \\item [Line Blanketing:] The use of fully line-blanketed NLTE models for  the analysis of Y453 might result in a somewhat lower temperature  (0.05 dex) without any changes in surface gravity (Lanz et al. \\cite{lanz97}; Haas \\cite{haas97}). Such a lower temperature  would increase the derived mass by about 15\\%.  \\item [Differential Reddening:]  Cudworth \\& Rees (\\cite{cud90}) find a gradient in the reddening that would increase the adopted reddening for Y453 (\\magpt{0}{35})  by about \\magpt{0}{015}. Lyons et al. (\\cite{lyon95}) report a patchiness in  the reddening toward M~4 that is at least as significant as the gradient, and  find a total range of 0.16 mag in \\Ebv . An increase by such a large amount  would still lead to a mass  of only 0.3~\\Msolar.  Due to the non-standard reddening law toward M~4, the reddening correction for Y453 has a rather high uncertainty in any case.  \\item [Distance:] The adopted distance (1.72 kpc) to M~4 is  on the low side of the range of distance determinations but is supported by both a recent astrometric measurement (Rees \\cite{rees96}), and HST observations of the main-sequence (Richer et al.\\ \\cite{rich97}).     Use of the distance given by Harris (\\cite{harris96}; 2200 pc) would give a mass of  0.26~\\Msolar.  \\item [Photometry:]   The only ground-based photometry of Y453 of which we are aware is the photographic photometry  of Cudworth \\& Rees (\\cite{cud90}), who also derive a 99\\% probability  of cluster membership from its proper motion. Y453 is among the faintest  stars studied by Cudworth \\& Rees, so the photometric precision might be poorer than their quoted 0.025 mag (which corresponds to an error of 2\\% in M).  \\end{description}  \\subsection{NGC~6723}  The V and B$-$V magnitudes in Table 1 for III-60 are from Menzies (\\cite{menz74}); we are not aware of any other photometry of this star. The tabulated photometry for IV-9 is from L.K. Fullton (1997, priv. comm.), who also gives U--B $= -0.84$. Photometry for IV-9 was also obtained by  Menzies (\\cite{menz74}; V = 14.86, B$-$V $= -0.25$), and  Martins \\& Fraquelli (\\cite{mafr87}; V=14.69, B$-$V $= -0.142$).  III-60 and IV-9  fit well on the 0.546~\\Msolar\\ post-early AGB track (see above) of  Sch\\\"onberner, and also have luminosities (\\logl $\\sim 3.0$) consistent with  being post-early AGB stars.  The spectrum of IV-9, however, was difficult to fit with any single model. As can be seen from  Fig. 1 there is an absorption feature blueward of the \\ion{He}{I} line at  4713~\\AA , which matches the \\ion{He}{II} absorption line at  4686~\\AA\\ in wavelength.  The  \\ion{He}{II} line strength and the Balmer line profiles can be reproduced by a model with \\teff\\ =   30,000~K, \\logg\\ = 4.08, and \\loghe\\ = $-$0.89.  However, this model is inconsistent  with both the size of the Balmer jump and with the photometric indices  (optical and UV) of this star.  The photometric data indicate a temperature  of 20000 -- 21000~K instead. Excluding the absorption feature at 4686~\\AA\\ from the fit results in a temperature  of 20700~K and a \\logg\\ value of 3.34, in good agreement with the  photometric temperature.  The detection of metal lines (such as the \\ion{O}{II}  absorption lines in BD+33$^\\circ$2642 discussed by Napiwotzki et al., \\cite{nahe93}) could  help to decide between the two temperatures. Simulations with theoretical  spectra, however, show that due to the low resolution of our data we cannot  expect to see any metal lines. Any decision will therefore have to await  better data. We keep the cooler temperature for all  further analysis because of the good agreement with the photometric indices.  \\subsection{NGC~6752}  B2004 was one of only four post-EHB candidate stars present in the UIT color-magnitude diagram of  NGC~6752 reported by Landsman et al. (\\cite{land96}),  and the position of B2004 in the \\logt --\\logg\\ plot (Fig. 3) is consistent with post-EHB tracks.    Landsman et al.\\ estimated \\teff\\ = 45000~K and \\logl\\ = 2.12 for B2004 on the basis of IUE spectrophotometry. However, the IUE photometry of B2004 had large uncertainties due to the presence of the nearby ($2.5''$ distant) blue HB star B1995, and the \\teff\\ (37000~K) and luminosity (\\logl\\ = 1.94) of B2004 derived  here should be more accurate. Spectroscopic analyses of the other three post-EHB candidate stars (B852, B1754, and B4380) in NGC~6752  were presented by Moehler et al.\\ (\\cite{mohe97}). The four post-EHB stars in NGC~6752 occupy a fairly narrow range in temperature  (4.5 $< \\log$ \\teff\\ $<$ 4.6) and luminosity ($1.94 < $ \\logl\\ $ < 2.12$),  and are separated by a large luminosity gap (0.5~dex) from stars on the populous EHB.    As discussed by Landsman et al.\\ (\\cite{land96}), these two characteristics are consistent with the non-canonical HB models of Sweigart (\\cite{sweig97}), which include helium mixing on the RGB.    However, a more definitive test of EHB evolutionary tracks will require a larger sample of post-EHB stars."
    },
    "9803/astro-ph9803125_arXiv.txt": {
        "abstract": "We describe an automated method for detecting clusters of galaxies in imaging and redshift galaxy surveys.  The Adaptive Matched Filter (AMF) method utilizes galaxy positions, magnitudes, and---when available---photometric or spectroscopic redshifts to find clusters and determine their redshift and richness.  The AMF can be applied to most types of galaxy surveys: from two-dimensional (2D) imaging surveys, to multi-band imaging surveys with photometric redshifts of any accuracy (2$\\half$D), to three-dimensional (3D) redshift surveys.  The AMF can also be utilized in the selection of clusters in cosmological N-body simulations.  The AMF identifies clusters by finding the peaks in a cluster likelihood map generated by convolving a galaxy survey with a filter based on a model of the cluster and field galaxy distributions. In tests on simulated 2D and 2$\\half$D data with a magnitude limit of $r' \\approx 23.5$, clusters are detected with an accuracy of $\\Delta z \\approx 0.02$ in redshift and $\\sim$10\\% in richness to $z \\lesssim 0.5$.  Detecting clusters at higher redshifts is possible with deeper surveys.  In this paper we present the theory behind the AMF and describe test results on synthetic galaxy catalogs. ",
        "introduction": "Clusters of galaxies---the most massive virialized systems known---provide powerful tools in the study of cosmology: from tracing the large-scale structure of the universe (\\cite{Bahcall88}, \\cite{Huchra90}, \\cite{Postman92}, \\cite{Dalton94}, \\cite{Peacock94} and references therein) to determining the amount of dark matter on Mpc scales (\\cite{Zwicky57}, \\cite{Tyson90}, \\cite{Kaiser93}, \\cite{Peebles93}, \\cite{Bahcall95}, \\cite{Carlberg96}) to studying the evolution of cluster abundance and its cosmological implications (\\cite{Evrard89}, \\cite{Peebles89}, \\cite{Henry92}, \\cite{Eke96}, \\cite{Bahcall97}, \\cite{Carlberg97}, \\cite{Oukbir97}).  The above studies place some of the strongest constraints yet on cosmological parameters, including the mass-density parameter of the universe, the amplitude of mass fluctuations at a scale of 8~$\\hMpc$ and the baryon fraction.  The availability of complete and accurate cluster catalogs needed for such studies is limited.  One of the most used catalogs, the Abell catalog of rich clusters (\\cite{Abell58}, and its southern counterpart \\cite{Abell89}), has been extremely useful over the past four decades. This catalog, which contains $\\sim$4000 rich clusters to $z \\lesssim 0.2$ over the entire high latitude sky, with estimated redshifts and richnesses for all clusters, was constructed by visual selection from the Palomar Sky Survey plates, using well-defined selection criteria. The Zwicky cluster catalog (\\cite{Zwicky61}) was similarly constructed by visual inspection.  The need for new, objective, and accurate large-area cluster catalogs to various depths is growing, following the important use of clusters in cosmology.  Large area sky surveys using CCD imaging in one or several colors, as well as redshift surveys, are currently planned or underway, including, among others, the Sloan Digital Sky Survey (SDSS). Such surveys will provide the data needed for constructing accurate cluster catalogs that will be selected in an objective and automated manner. In order to identify clusters in the new galaxy surveys, a robust and automated cluster selection algorithm is needed.  We propose such a method here.  Cluster identification algorithms have typically been targeted at specific surveys, and new algorithms have been created as each survey is completed.  \\cite{Abell58} was the first to develop a well-defined method for cluster selection, even though the identification was carried out by visual inspection (see, e.g., \\cite{McGill90} for a analysis of this method).  Other algorithms have been created for the APM survey (\\cite{Dalton94}, \\cite{Dalton97}; see \\cite{Schuecker98} for a variant of this method), the Edinburgh-Durham survey (ED; \\cite{Lumsden92}), and the Palomar Distant Cluster Survey (PDCS; \\cite{Postman96}; see also \\cite{Kawasaki97} for a variant of this method; and \\cite{Kleyna97} for an application this method to finding dwarf spheroidals).  All the above methods were designed for and applied to two-dimensional imaging surveys.  In this paper we present a well defined, quantitative method, based on a matched filter technique that expands on some of the previous methods and provides a general algorithm that can be used to identify clusters in any type of survey.  It can be applied to 2D imaging surveys, 2$\\half$D surveys (multi-band imaging with photometric redshift estimates of any accuracy), 3D redshift surveys, as well as combinations of the above (i.e. some galaxies with photometric redshifts and some with spectral redshifts).  In addition, this Adaptive Matched Filter (AMF) method can be applied to identify clusters in cosmological simulations.  The AMF identifies clusters by finding the peaks in a cluster likelihood map generated by convolving a galaxy survey (2D, 2$\\half$D or 3D) with a filter which models the cluster and field galaxy distribution. The peaks in the likelihood map correspond to locations where the match between the survey and the filter is maximized.  In addition, the location and value of each peak also gives the best fit redshift and richness for each cluster.  The filter is composed of several sub-filters that select different components of the survey: a surface density profile acting on the position data, a luminosity profile acting on the apparent magnitudes, and, in the 2$\\half$D and 3D cases, a redshift cut acting on the estimated redshifts.  The AMF is adaptive in three ways. First, the AMF adapts to the errors in the observed redshifts (from no redshift information (2D), to approximate (2$\\half$D) or measured redshifts (3D)).  Second, the AMF uses the location of the galaxies as a ``naturally'' adaptive grid to ensure sufficient spatial resolution at even the highest redshifts. Third, the AMF uses a two step approach that first applies a coarse filter to find the clusters and then a fine filter to provide more precise estimates of the redshift and richness of each cluster.  We describe the theory of the AMF in \\S2 and its implementation in \\S3.  We generate a synthetic galaxy catalog to test the AMF in \\S4 and present the results in \\S5.  We summarize our conclusions in \\S6. ",
        "conclusions": "We have presented the Adaptive Matched Filter method for the automatic selection of clusters of galaxies in a wide variety of galaxy catalogs. The AMF can find clusters in most types of galaxy surveys: from two-dimensional (2D) imaging surveys, to multi-band imaging surveys with photometric redshifts of any accuracy (2$\\half$D), to three-dimensional (3D) redshift surveys.  The method can also be utilized in the selection of clusters in cosmological N-body simulations. The AMF is based on matching the galaxy catalog with a cluster filter that models the overall galaxy distribution. The model describes the surface density, apparent magnitude, and redshift of cluster and field galaxies. Convolving the data with the filter produces a cluster probability map whose peaks correspond to the location of the clusters. The probability peaks also yield the best fit redshift and richness of each cluster.  The heart of the AMF is the apparent overdensity $\\delta_i$ which is evaluated at each galaxy position and has a higher value for galaxies in clusters than galaxies in the field.   The apparent overdensity distills the entire description of the galaxy catalog into a single function. Two likelihood functions are derived, $\\sL_\\coarse$ and $\\sL_\\fine$, using different underlying model assumptions. The theoretical framework of the AMF allows  estimated redshifts to be included via a simple redshift filter, which effectively limits the sums in $\\sL_\\coarse$ and $\\sL_\\fine$ to those galaxies within a window around $z_c$.  The maxima in the likelihood functions are used to identify cluster positions as well as their redshifts and richnesses.  The AMF is adaptive in three ways.  First, it adapts to the errors in the estimated redshifts.  Second, it uses the locations of the galaxies as ``naturally'' adaptive grid to ensure sufficient resolution at even the highest redshifts.  Third, it uses a two step approach that applies a coarse filter to initially find the clusters and a fine filter to more precisely estimate the redshift and richness of each cluster.  We tested the AMF on a set of simulated clusters with different richnesses and redshifts---ranging from groups to rich clusters at redshifts 0.1 to 0.5; the clusters were placed in a simulated field of randomly distributed galaxies as well as in a non-random distribution produced by N-body cosmological simulations.  We find that the AMF detects clusters with an accuracy of $\\Delta z \\sim$0.02 in redshift and $\\sim$10\\% in richness to $z \\lesssim 0.5$ (for a simulated galaxy survey to $r' \\approx 23.5$).  In addition, robustness tests provide a strong indication that the AMF will perform well on observational data sets.  Detecting clusters at even higher redshifts will be possible in deeper surveys."
    },
    "9803/astro-ph9803038_arXiv.txt": {
        "abstract": "We present a new method for reconstructing two-dimensional mass maps of galaxy clusters from the image distortion of background galaxies. In contrast to most previous approaches, which directly convert locally averaged image ellipticities to mass maps (direct methods), our entropy-regularized maximum-likelihood method is an inverse approach. Albeit somewhat more expensive computationally, our method allows high spatial resolution in those parts of the cluster where the lensing signal is strong enough. Furthermore, it allows to straightforwardly incorporate additional constraints, such as magnification information or strong-lensing features. Using synthetic data, we compare our new approach to direct methods and find indeed a substantial improvement especially in the reconstruction of mass peaks. The main differences to previously published inverse methods are discussed. ",
        "introduction": "The reconstruction of projected cluster mass maps from the observable image distortion of faint background galaxies due to the tidal gravitational field is a new and powerful technique. Pioneered by Kaiser \\& Squires (1993), this method has since been modified and generalized to account for (a) strong tidal fields in cluster centers (Schneider \\& Seitz 1995; Seitz \\& Schneider 1995; Kaiser 1995); (b) finite and -- in some cases, e.g.~WFPC2 images -- very small data fields (Schneider 1995; Kaiser et al.\\ 1995; Bartelmann 1995; Seitz \\& Schneider 1996, 1998; Lombardi \\& Bertin 1998); and (c) the broad redshift distribution of background galaxies (Seitz \\& Schneider 1997). All of these are direct methods in the sense that a local estimate of the tidal field is derived from observed galaxy ellipticities, which is then inserted into an inversion equation to obtain an estimate of the surface mass density of the cluster.  Whereas these direct methods are computationally fast, can be treated as black-box routines, need only the observed ellipticities and a smoothing length $\\theta_0$ as input data, and yield fair estimates of the surface mass density, their application has several drawbacks:  \\begin{itemize}  \\item The data must be smoothed, and the smoothing scale is typically a free input parameter specified prior to the mass reconstruction. There are no objective criteria on how to set the smoothing scale, although some ad-hoc prescriptions for adapting it to the strength of the lensing signal have been given (Seitz et al.\\ 1996). In general, smoothing leads to an underestimate of the surface mass density in cluster centers or sub-condensations.  \\item The quality of the reconstruction is hard to quantify.  \\item Constraints on the mass distribution from additional observables (such as multiple images or giant arcs) cannot simultaneously be included. In particular, magnification information contained in the number density of background sources (Broadhurst et al.\\ 1995; Fort et al.\\ 1997) or in the image sizes at fixed surface brightness (Bartelmann \\& Narayan 1995), cannot be incorporated locally but only globally to break the mass-sheet degeneracy (Gorenstein et al.\\ 1988; Schneider \\& Seitz 1995).  \\end{itemize}  To overcome these drawbacks, a different class of methods should be used. Bartelmann et al.\\ (1996, hereafter BNSS) developed a maximum-likelihood (ML) technique in which the values of the deflection potential at grid points are considered as free parameters. After averaging image ellipticities and sizes over grid cells, local estimates of shear and magnification are obtained. The deflection potential at the grid points is then determined such as to optimally reproduce the observed shear and magnification estimates. Magnification information can be included this way. The smoothing scale in this method is given by the size of the grid cells, and can be chosen such that the overall $\\chi^2$ of the fit is of order unity per degree of freedom.  Squires \\& Kaiser (1996; hereafter SK) suggested several inverse methods. Their {\\em maximum probability method\\/} parameterizes the mass distribution of the cluster by a set of Fourier modes. If the number of degrees of freedom (here the number of Fourier modes) is large, the mass model tends to over-fit the data. This has to be avoided by regularizing the model, for which purpose SK impose a condition on the power spectrum of the Fourier modes. SK's {\\em maximum-likelihood method\\/} specifies the surface mass density on a grid and uses the Tikhonov-Miller regularization (Press et al.\\ 1992, Sect.\\ 18.5). The smoothness of the mass reconstructions can be changed by varying the regularization parameter, which is chosen such as to give an overall $\\chi^2\\approx1$ per degree of freedom.  Bridle et al.\\ (1998) have recently proposed an entropy-regularized ML method in which the cluster mass map is parameterized by the surface mass density at grid points. This method allows to restrict the possible mass maps to such with non-negative surface mass density.  This paper describes another variant of the ML method (Seitz 1997, Ph.D.\\ thesis). The major differences to the previously mentioned inverse methods are the following:  \\begin{itemize}  \\item The observational data (e.g.~the image ellipticities) are not smoothed, but each individual ellipticity of a background galaxy is used in the likelihood function. Whereas this modification complicates the implementation of the method, it allows larger spatial resolution for a given number of grid points, which is useful since the latter determines the computing time.  \\item The number of grid points can be much larger than in BNSS, and the likelihood function is regularized. This produces mass reconstructions of variable smoothness: Mass maps are smooth where the data do not demand structure, but show sharp peaks where required by the data. The resulting spatially varying smoothing scale is a very desirable feature. Fourier methods, such as SK's maximum probability method, have a spatially constant smoothing scale which is determined by the highest-order Fourier components. They always need to compromise between providing sufficient resolution near mass peaks and avoiding over-fitting of the data in the outer parts of a cluster.  \\item Following BNSS, we use the deflection potential to describe a cluster. This is an essential difference to Bridle et al.\\ (1998) who used the surface mass density at grid points. As we shall discuss below, working with the deflection potential has substantial fundamental and practical advantages.  \\end{itemize}  We describe our method in Sect.\\ 2, with details given in the Appendix. We then apply the method to synthetic data sets in Sects.\\ 3 \\& 4 to demonstrate its accuracy. In particular, we compare the performance of our ML method to that of direct methods. The results are then discussed in Sect.\\ 4, and conclusions are given in Sect.\\ 5, where we also discuss further generalizations of the method for, e.g., including constraints from strong lensing features. ",
        "conclusions": "We presented a new method for reconstructing projected mass distributions of galaxy clusters. The method uses image distortions of background galaxies and their size as a function of surface brightness. Our entropy-regularized ML method (Seitz 1997) is a further development of previously published inverse methods for the mass reconstruction. In particular, we describe the lens by its deflection potential $\\psi$ as suggested by BNSS. This is of major importance, for two reasons. First, if the surface mass density $\\kappa$ on a finite field $\\U$ is used to describe the lens, the shear on $\\U$ is incompletely specified by the model because the mass distribution outside $\\U$ can contribute to the shear. Second, the shear at the position of any galaxy depends only locally on $\\psi$, which allows a much faster minimization algorithm for a given number of grid points.  We regularize the method by an entropy term as suggested by Bridle et al.\\ (1998), but additionally adapt the prior to the current model of the mass distribution. This `moving prior' (Lucy 1994) allows a considerably higher resolution of mass peaks. The spatial resolution of the entropy-regularized ML method adapts itself to the strength of the lensing signal, producing mass distributions which are as smooth as possible, and as structured as the data require. In that respect, our method differs from that of BNSS and SK. We showed that the ML method is superior to the noise-filter method (Seitz \\& Schneider 1996) which was the most accurate of the presently known direct inversion methods (Seitz \\& Schneider 1996, 1998; SK; Lombardi \\& Bertin 1998).  Obviously, the method described here is not restricted to rectangular data fields, but can easily be adapted to any geometry of $\\U$ by covering $\\U$ with quadratic grid cells, and adding a boundary of grid points for $\\psi$ -- the rest is only a matter of labeling. Furthermore, we note that observational errors can be incorporated into the likelihood function. For example, if the measurement error of the ellipticity $\\chi$ is $\\sigma_{\\rm obs}$, one can replace $\\sigma_\\chi^2$ in (\\ref{eq:2.13} and \\ref{eq:2.14}) by $\\sigma_\\chi^2+\\sigma_{\\rm obs}^2$.  In contrast to the direct inversion methods, all of which are variants and generalizations of the original Kaiser \\& Squires (1993) method, the inverse methods allow to include additional information on top of the shear measured through image ellipticities. We demonstrated this here by adding magnification information derived from image sizes at given surface brightness, as discussed by Bartelmann \\& Narayan (1995). However, we could equally well use the change of number counts due to magnification bias (Broadhurst et al.\\ 1995) as an additional constraint. In that case, if the number counts of a certain (e.g.~color-selected) galaxy population have a cumulative slope of $-\\beta$, the expected number density of background galaxies at a position $\\vc\\theta$ is $n(\\vc\\theta)=n_0\\abs{\\mu(\\vc\\theta)}^{\\beta-1}$, where $n_0$ are the counts at the same flux limit in the absence of lensing. Assuming that galaxies are intrinsically randomly distributed, the probability of having $N$ galaxies within $\\U$ is a Poisson distribution $P_N(\\ave{N})$ with \\begin{equation} \\ave{N} = n_0\\int_\\U\\,\\d^2\\theta\\, \\abs{\\mu(\\vc\\theta)}^{\\beta-1}\\;. \\label{eq:5.1} \\end{equation} Consequently, the likelihood function could be augmented by a factor \\begin{equation} \\L_\\mu = P_N(\\ave{N})\\prod_{k=1}^N\\, \\abs{\\mu(\\vc\\theta_k)}^{\\beta-1}\\;. \\label{eq:5.2} \\end{equation} If galaxy clustering is important, the likelihood $\\L_\\mu$ cannot be written as a simple product over individual galaxies, but the joint probability distributions must be taken into account. The contribution of clustering effects to the likelihood function is somewhat uncertain, because an approximate expression has to be used due to lack of knowledge on the $N$-point correlation functions (see Broadhurst et al.\\ 1995 for further discussion).  Perhaps the most promising generalization of our method is the inclusion of strong lensing constraints. Since giant arcs and multiple images of background galaxies provide (nearly) exact constraints on the lens mass distribution, it is highly desirable to include them into a mass reconstruction. The obvious way to do this would be to augment the function $E$ by a term which measures the degree to which multiple images of the same source are mapped back to the same position in the source plane. In addition, the surface brightness profile of multiple images of extended sources can be incorporated, e.g.~in a similar manner as the spatially resolved multiple arc in the cluster Cl0024+16 (Colley at al.\\ 1996).  In some of the observed clusters, the lensing effects of individual galaxies are visible, in particular through deformations of giant arcs. Some of the most prominent examples are the triple arc in 0024+16 (Kassiola et al.\\ 1992), the multiple arc systems in A~2218 (Kneib et al.\\ 1996), and the distortion of the curvature in the arc of the galaxy cB58 in MS1512+36 (Seitz et al.\\ 1998). But even weaker lensing effects of individual (cluster) galaxies can be detected using a combination of cluster mass reconstruction and galaxy-galaxy lensing techniques. By adding two free parameters to the lens model, such as the mass-to-light ratio of cluster galaxies and their characteristic spatial extent, the size of halos of cluster galaxies can be investigated (Natarayan et al.\\ 1997; Geiger \\& Schneider 1998)."
    },
    "9803/astro-ph9803206.txt": {
        "abstract": "Differences between observed and theoretical eigenfrequencies of the Sun have characteristics which identify them as arising predominantly from properties of the oscillations in the vicinity of the solar surface: in the super-adiabatic, convective boundary layer and above. These frequency differences may therefore provide useful information about the structure of these regions, precisely where the theory of solar structure is most uncertain.  In the present work we use numerical simulations of the outer part of the Sun to quantify the influence of turbulent convection on solar oscillation frequencies.  Separating the influence into effects on the mean model and effects on the physics of the modes, we find that the main model effects are due to the turbulent pressure that provides additional support against gravity, and thermal differences between average 3-D models and 1-D models.  Surfaces of constant pressure in the visible photosphere are elevated by about 150 km, relative to a standard envelope model.  As a result, the turning points of high-frequency modes are raised, while those of the low-frequency modes remain essentially unaffected. The corresponding gradual lowering of the mode frequencies accounts for most of the frequency difference between observations and standard solar models.  Additional effects are expected to come primarily from changes in the physics of the modes, in particular from the modulation of the turbulent pressure by the oscillations. ",
        "introduction": "In standard solar models, the stratification of the convection zone is determined by mixing-length theory (MLT), thereby reducing the entire complexity of the turbulent hydrodynamics to a one-parameter family of models. MLT solar models suffer from several basic inconsistencies. For example, they predict that convective velocities, of several km$\\,$s$^{-1}$, disappear abruptly in a few tens of km, immediately below the solar surface. %at the edge of the convective zone. %Velocities are predicted to disappear In contrast, observations show convective cells with those same characteristic velocities, and with horizontal sizes of 1000 -- 5000 km, whose velocity fields obviously cannot vanish so abruptly. Indirect evidence from spectral line broadening indeed shows that the photosphere is pervaded by a velocity field with rms Mach numbers of the order of 0.3, and yet, in standard solar models these layers are assumed to be entirely quiescent.  Clearly, such a discrepancy between the theoretical description and the observations should be regarded as a warning not to take the quantitative predictions of the theory too seriously.  Helioseismology provides quantitative diagnostics that pertain precisely to these critical surface layers, since this is where the upper turning points of the majority of modes are located.  Thus, analysis of the observed properties of these modes may help clarify the consequences of the inconsistencies inherent in MLT and its more recent siblings (\\cite{CM91,CM92,Canuto+Goldman+Mazzitelli96}). Indeed, as we discuss in more detail below, adiabatic oscillations of MLT models show significant systematic discrepancies when compared with measured solar frequencies.  It is natural, therefore, to seek to use helioseismology applied to these differences to improve the theoretical description.  However, this procedure is undermined by our present uncertainty about the physics of the oscillations near the top of the convection zone where they are likely to be strongly coupled to both the convective and radiative fields. In the language of \\citetext{Balmforth92b}, the {\\emph extrinsic} (or {\\emph model}\\/) error in the mode frequencies (due to inaccurate modeling of the mean solar structure) cannot be accurately estimated while the {\\emph intrinsic} (or {\\emph modal}\\/) errors (due to uncertain mode physics) are largely unknown. %\\note [I have replaced extrinsic and intrinsic by model and %modal throughout; however, we may need another sentence here %to hammer the notation down. !jcd]  One approach to resolving this problem has been the time-dependent non-local non-adiabatic mixing-length theory of \\citetext{Gough77} and Balmforth (1992abc) in which the \\nocite{Balmforth92a,Balmforth92b,Balmforth92c} coupling of the oscillations to both the convection and the radiation is included within the framework of mixing-length theory. Another approach to the problem has been proposed by \\citetext{Zhugzhda+Stix94} who have developed a model of the modal effect on mode frequencies due to advection of the oscillations by spatially varying radial flows. This approach has not yet been developed to the stage where it can be usefully applied to realistic solar models with stratification and turbulent pressure. Finally, \\citetext{Rudiger+97} have used turbulence closure assumptions to parameterize the propagation of acoustic disturbances through a convecting medium.  In the present work, we use an alternative technique for estimating the model effects, based on the results obtained from numerical simulations of turbulent convection in a radiating fluid. We show that p modes can be calculated from a mean model with hydrostatic and thermodynamic stratification obtained by appropriately weighted averages of the simulation results. We proceed by making simplifying, and certainly very naive, assumptions about the modal effects, postponing their detailed study to subsequent papers.  We begin (Section 2) with a brief discussion of the helioseismic data and the discrepancy between measured frequencies and those calculated from MLT models.  We then discuss (Section 3) the averaging techniques needed to analyze the radial oscillations of a convecting layer, describe the model computations (Section 4), and investigate the resulting oscillation frequencies (Section 5).  Finally (Section 6), we discuss the relevance and limitations of the results, and indicate future plans. ",
        "conclusions": "In the classical theory of solar structure, a one-parameter family of models (MLT) is calibrated against the known radius of the Sun. It is well known that MLT is based on a number of fundamentally inapplicable and inconsistent assumptions, and so a large number of alternative models of stellar convection have been proposed.  Non-local mixing-length theory (\\cite{Gough77}; Balmforth 1992abc\\nocite{Balmforth92a,Balmforth92b,Balmforth92c}) attempts to improve on standard MLT by incorporating into it the effect of the finite size of turbulent eddies. \\citetext{Forestini+91}  have produced an MLT-type model incorporating a measure of asymmetry between upflows and downflows as found in the simulations. The models of \\citeauthor{Lydon+92} (\\citeyear{Lydon+93a}, \\citeyear{Lydon+93b}) are essentially attempts to parameterize a wide range of convection simulations using a formalism similar to MLT but incorporating also the contribution of the kinetic energy flux to the flux-balance equation. \\citetext{CM91} have taken an approach based on modern theories of turbulence, and have attempted to produce a parameterized expression based on such theories. Finally, Canuto (\\citeyear{Canuto92,Canuto93}) has produced an ambitious model of convection based on a Reynolds' stress formalism.  These more elaborate models of convection are potentially very useful, if it can be shown that they capture essential aspects of the full 3-D convection in terms of much simpler equations.  In particular, non-local and time-dependent models of convection such as the ones by \\citetext{Gough77}, Buchler (1993), and Houdek (1997) would be very \\nocite{Houdek97,Buchler93} useful, for example in the modeling of pulsating stars, since full 3-D simulations are too expensive to be used in that context. The most obvious way to proceed is by use of numerical simulations such as the one used here or those of \\citetext{Kim+95}, treating the simulations as data against which the models are to be tested and calibrated.  A second approach to testing the simplified models, and one which has been more widely adopted so far, is the helioseismic approach in which the frequencies of oscillations of a solar model constructed with a given theory are compared with the observed frequencies. Such a procedure, however, must deal with the difficulty that the complete structure of the surface layers is certainly underdetermined by the (low- and intermediate-degree) oscillation frequencies since, as \\citetext{JCD+Thompson97} have shown, the near-surface contribution to the frequencies depends only on $\\upsilon$ and not separately on, e.g., $p_0$ and $\\Gamma_1$. Thus improved agreement with measured mode frequencies cannot, by itself, be taken as evidence that a given model of convection is a better description of reality.  In the present work we have attempted to improve on this approach by analyzing the problem in such a way as to obtain a physically justifiable description of the oscillations. Moreover, by using a numerical simulation of convection we give ourselves no free parameters with which to calibrate our model. Given this approach, the fact that we are able to construct a complete solar envelope model with essentially the correct convection-zone depth is, in itself, a considerable achievement for the simulation.  We have here investigated model effects on the mode frequencies, primarily those due to changes in pressure support of the atmosphere and 3-D radiation transfer.  The effects we have found are robust; there is no question that the averaging of 3-D fluctuations results in differences of this sense and order of magnitude; the turbulent pressure elevation is constrained by the observed photospheric velocity field, and the mean thermal difference is an inevitable consequence of the temperature dependence of the opacity.  One might attempt to estimate the elevation effect from turbulent pressure using a local convection model such as MLT or the model of Canuto \\& Mazzitelli (1991, 1992). However, as discussed by \\citetext{Antia+Basu97score}, such models cannot accurately account for the effect of turbulent pressure because they result in an artificially abrupt upper boundary of the convection zone and therefore a serious overestimate of the turbulent-pressure gradient there.  In addition, as discussed in section \\ref{pturb.sec}, a 1-D model with the correct pressure stratification unavoidably has a surface radiation flux that corresponds to an incorrect effective temperature.  Thus, there are inherent limitations in simplified 1-D models of convection.  The principal remaining uncertainty in determining the oscillation frequencies from 3-D models lies in the uncertain mode physics. In particular, non-adiabatic effects, and the response of the turbulent pressure to the compression and rarefaction in the oscillations needs to be understood. % We have attempted to quantify this uncertainty by calculating two models: the RGM in which the turbulent pressure is assumed to be unaffected by the oscillations and the GGM in which it is assumed to respond in exactly the same way as does the gas pressure. The frequency discrepancies for the RGM are almost exactly twice those for the GGM. The GGM produces frequencies that are closer to those observed, but this should of course not be taken as evidence that the turbulent pressure responds in exactly the same way as the gas pressure.  The actual depth- and time-dependent mode response of the turbulent pressure produces, together with the response of the gas pressure, a complex-valued, and frequency-dependent $\\Gamma_1(z,\\omega)$. % At any particular frequency, the real part of $\\Gamma_1$ may be expected to have a different depth dependence than that assumed in both the RGM and GGM models; therein lies an essential part of the modal effects, and thus a potential explanation for part of the remaining differences between the observed and calculated oscillation frequencies.  Preliminary investigations of the relation between $\\delta \\ln \\bra \\rho \\ket$ and $\\delta \\ln \\bra P \\ket$ in numerical 3-D simulations overlaid with initial radial eigenmodes show that non-adiabatic effects are indeed significant. The effective gamma appears to be closer to unity than to $5/3$ in the optically thin parts of the photosphere, while in the very surface layers the effective $\\Gamma_1$ can become quite large ($\\sim 8$), because of a localized reduction of $\\delta\\rho$. A more quantitative analysis of the non-adiabatic effects will require much more work, though, and will appear in a subsequent paper.  Additional differences (in particular the ones reflected in the discrepancy of the f-mode frequencies) are expected to come from true 3-D effects; differences between mode behavior in a homogeneous and inhomogeneous medium.  Again, a quantitative investigation of such effects requires elaborate, differential comparisons between non-radial mode behavior in 1-D and 3-D models, and is beyond the scope of the present paper.  However, a pre-requisite for studying mode physics effects is a sufficient accuracy of the mean model; only if one includes the model effects with sufficient accuracy does it make sense to use remaining discrepancies to diagnose modal effects.  How accurate, then, is the pressure stratification obtained from the present numerical simulations?  The tests at various numerical resolutions show that the location, shape, and width of the peak of the turbulent pressure (relative to the total pressure) are quite insensitive to the numerical resolution, while the magnitude of the turbulent pressure increases slightly with increasing numerical resolution (\\Fig{F8}).  The magnitude scaling of the turbulent pressure is, however, tightly constrained by spectral line-broadening observations (\\Fig{F9}). The existence of turbulent pressure support of about the magnitude found here thus cannot be doubted.  Moreover, part of the elevation effect is due to the thermal difference between 1-D and average 3-D models.  Any calculation of mode frequencies that does not include a turbulent and thermal pressure elevation of the upper turning points of the modes is thus neglecting a significant effect.  If parameter fitting for such a calculation leads to near agreement with the observations it merely illustrates that it is quite possible to obtain ``the right result for the wrong reason''.  Finally we must emphasize that while our understanding of convective effects on radial oscillations may seem rudimentary, it is in fact highly sophisticated by comparison with our understanding of their influence on nonradial modes. Indeed if we consider the most nonradial mode of all, the f mode, in which radial and nonradial motions are of equal magnitude, we note that no explanation which seeks to replace convection modeling with a hydrostatic solar model can ever explain the measured frequency residuals because f-mode frequencies are largely insensitive to hydrostatic structure. Thus, both the f-mode frequency discrepancies and the remaining discrepancies in \\Fig{mobsdiff} are likely to be caused by modal effects, rather than by the stratification effects that we have uncovered in the present paper. Evidently the behaviour of both nonradial and radial modes needs more study.  In order to address the modal effects on nonradial modes it will be necessary to invoke a more elaborate technique than the simple planar averages we have used, a result already evident from the simplified model calculations of \\citetext{Zhugzhda+Stix94}. On the one hand this emphasizes the gulf which still exists between theory and measurement in mode physics but, on a more positive note, it suggests that the remaining frequency discrepancies may have great power as diagnostic probes of the structure of turbulent convection."
    },
    "9803/nucl-th9803026_arXiv.txt": {
        "abstract": "The properties of hot matter are studied in the frame of the relativistic Brueckner-Hartree-Fock theory. The equations are solved self-consistently in the full Dirac space. For the interaction we used the potentials given by Brockmann and Machleidt. The obtained critical temperatures are smaller than in most of the nonrelativistic investigations. We also calculated the thermodynamic properties of hot matter in the relativistic Hartree--Fock approximation, where the force parameters were adjusted to the outcome of the relativistic Brueckner--Hartree--Fock calculations at zero temperature. Here, one obtains higher critical temperatures, which are comparable with other relativistic calculations in the Hartree scheme. ",
        "introduction": "The properties of hot and dense nuclear matter play an essential role in the understanding of high-energy heavy-ion collisions, supernova explosions and proto-neutron stars. For that reason the problem of hot nuclear matter has been studied over the last decades in several investigations, which however were predominantly performed within the nonrelativistic scheme \\cite{1a}, using either effective density dependent interactions \\cite{1a,1,2,23} or the Brueckner approach \\cite{1a,3,4,5,5a}. In the relativistic approach investigations of the equations of state for $T\\neq 0$ are relatively scarce. The majority of such calculations were performed in the relativistic Hartree approximation (RH), where the extension to finite temperatures is straightforward. Details of this scheme are given, for instance, in Refs.\\,\\cite{23}, \\cite{6}-\\cite{11}. More complicated are the relativistic Hartree--Fock-- \\cite{13,11} and the Brueckner--Hartree--Fock approximation \\cite{12}, and the application to finite nuclei \\cite{14}.  In this contribution we will concentrate on the relativistic Brueckner--Hartree--Fock treatment (RBHF) of symmetric and asymmetric nuclear matter generalizing the formalism as described in Refs.\\,\\cite{15,15a} to $T\\neq 0$. The RBHF--approach seems to be of special interest, since it is known for $T = 0$ that the resulting EOSs are much stiffer than their nonrelativistic counterparts \\cite{15b}. To our knowledge such an investigation has been only performed so far by the Groningen group for symmetric nuclear matter \\cite{12}. As described in more details in Refs.\\,\\cite{12,15}, their method uses a nonunique ansatz for the T--matrix in terms of five Fermi invariants, which can lead to ambiguous results for the self-energies (see, e.g., Refs.\\,\\cite{17,18}). In order to avoid this problem we solve, according to the original scheme of the Brooklyn group \\cite{19}, the RBHF--approximation in the full Dirac space, which is more tedious (for a more detailed discussion, see Ref.\\,\\cite{15}). Since the formal structure of the problem is the same as for $T=0$, where one has to solve three coupled equations, namely the Dyson equation for the one-body Green's function $G$, the (reduced) Bethe--Salpeter equation for the effective scattering matrix $T$ in matter and the equation for the self--energy $\\Sigma$, we will not repeat here the equations. As in Refs.\\,\\cite{12,15,15a} we will restrict ourselves to the incorporation of intermediate positive--energy nucleon states only, where now the Fermi step functions are replaced by the Fermi distribution functions $n_{\\vec p}\\,(T)$ at finite  temperature $T$ (for details, see Ref.\\,\\cite{21}).  The Green's function obeys for $T\\neq 0$ the spectral representation \\cite{11,20} \\begin{equation} \\label{I.1} G(p) = \\int d\\omega\\,A(\\vec{p},\\omega) \\left\\{\\frac{f(\\omega)}{p_\\rho -\\omega - i\\eta} + \\frac{f(-\\omega)}{p_\\rho -\\omega+i\\eta}\\right\\} ~, \\end{equation} with \\begin{equation} \\label{I.2} f(\\omega) = (e^{\\beta\\omega}+1)^{-1} ~ {\\stackrel{T=0}{_{\\mbox{$\\longrightarrow$}}}} ~ \\Theta(-\\omega) ~. \\end{equation} The formal structure of the spectral function $A(\\vec{p},\\omega)$ \\cite{15} remains unaltered to the case for  $T=0$.   A further difference in comparison with Ref.\\,\\cite{12} is that we take the momentum dependence of the self--energies into account. Since the pole of the quasi--particle propagators occurs for $T\\neq 0$ in the integration domain of the intermediate states, one obtains, in principle, complex effective scattering matrix elements and self--energies. It was checked in Ref.\\,\\cite{12} that the imaginary part of $\\Sigma$ turned out to be small. Therefore we neglect also Im~$\\Sigma$ in the calculations. For the one--boson--exchange interaction we used the modern potentials constructed by Brockmann and Machleidt \\cite{22}. We select for the presentation  the so-called potential $B$, which gives the best results for the EOS ($E/A = -15.73$~MeV; $\\rho_0 = 0.172$~fm$^{-3}; K = 249$~MeV; $J = 32.8$~MeV) at zero temperature in RBHF--calculations (see Refs.\\,\\cite{15,15a}, for the potential $A$ the outcome is similar, see Ref.\\,\\cite{21}). For the sake of comparison we  also treated the RHF--approximation, where we adjusted the force parameters to the outcome for the EOS for symmetric and asymmetric nuclear matter at $T=0$ \\cite{15a,21}. The RHF--approximation has in comparison with the RH--approximation the advantage to resemble in its formal structure more to the RBHF--approximation with the benefit of a much simpler numerical treatment than in the RBHF case.  For finite temperatures one needs for the determination of the pressure the free energy per baryon, defined as \\begin{equation} \\label{I.3} f = u - T s ~, \\end{equation} with the entropy per baryon: \\begin{equation} \\label{I.4} s = - \\frac{2}{\\rho h^3} \\sum_\\tau \\int d^3p\\,[n(\\vec{p})\\; {\\rm ln} \\,n(\\vec{p}) + \\left(1 - n(\\vec{p}\\right)\\; {\\rm ln}\\,\\left(1 - n(\\vec{p}\\right)]~. \\end{equation} The pressure is given by: \\begin{equation} \\label{I.5} P = \\rho \\sum_\\tau \\rho_\\tau \\left(\\frac{\\partial f} {\\partial\\rho_\\tau}\\right)_{T,\\rho_{-\\tau}} ~. \\end{equation} ",
        "conclusions": "In conclusion, we have performed a calculation of hot symmetric and asymmetric nuclear matter within the relativistic Brueckner--Hartree--Fock scheme using modern OBE--interactions constructed by Brockmann and Machleidt. It turned out that the critical temperatures are smaller than it is the case for the majority of nonrelativistic treatments. We have additionally treated the relativistic Hartree--Fock approximation at $T\\neq 0$, where the Lagrangian parameters were adjusted to the outcome of the RBHF--treatment for $T=0$. Here the critical temperature is in the range of other relativistic treatments performed in the Hartree scheme."
    },
    "9803/astro-ph9803042_arXiv.txt": {
        "abstract": "We report on the discovery of a new kind of thermal pulse in intermediate mass AGB stars. Deep dredge-up during normal thermal pulses on the AGB leads to the formation of a long, unburnt tail to the helium profile. Eventually the tail ignites under partially degenerate conditions producing a strong shell flash with very deep subsequent dredge-up.  The carbon content of the intershell convective region (X$_{C}$ $\\sim$ 0.6) is substantially higher than in a normal thermal pulse (X$_{C}$ $\\sim$ 0.25) and about 4 times more carbon is dredged-up than in a normal pulse. ",
        "introduction": "It is now just over 30 years since thermal pulses were discovered in AGB stars by \\cite{Sch65} and over 20 years since third dredge-up was discovered by \\cite{Ibe75}, and we are still learning about the consequences of these events for stellar evolution and nucleosynthesis. Although much is known  there are still many uncertainties, especially concerning third dredge-up (\\cite{Fro96}) and mass-loss. We are presently studying the effects of hot bottom burning (HBB) on intermediate mass AGB stars during their thermally pulsing evolution. During these calculations we found a new kind of thermal pulse which we report on in this paper. ",
        "conclusions": "We believe that the development of degenerate thermal pulses is inevitable provided two criteria are met:~1)~dredge-up is very deep, providing cooling of the helium shell soon after ignition;~and~2)~the star lives long enough on the AGB for the helium shell to reach (partially) degenerate conditions. Whether these conditions are realised in real stars remains an open question. The two aspects of AGB evolution which are the most uncertain are the extent of dredge-up and the rate of mass-loss (hence AGB lifetime), precisely the two phenomena which govern the occurrence or not of degenerate thermal pulses.   It seems prudent to look for some observational consequence of degenerate thermal pulses. The nucleosynthetic consequences are currently being investigated, and will be reported elsewhere.  \\clearpage"
    },
    "9803/astro-ph9803274_arXiv.txt": {
        "abstract": "Deep $J, H,$ and $K$ images are used to probe the evolved stellar contents in the central regions of the Sculptor group galaxies NGC55, NGC300, and NGC7793. The brightest stars are massive red supergiants (RSGs) with $K \\sim 15 - 15.5$. The peak RSG brightness is constant to within $\\sim 0.5$ mag in $K$, suggesting that NGC55, NGC300, and NGC7793 are at comparable distances. Comparisons with bright RSGs in the Magellanic Clouds indicate that the difference in distance modulus with respect to the LMC is $\\Delta \\mu = 7.5$. A rich population of asymptotic giant branch (AGB) stars, which isochrones indicate have ages between 0.1 and 10 Gyr, dominates the $(K, J-K)$ color-magnitude diagram (CMD) of each galaxy. The detection of significant numbers of AGB stars with ages near 10 Gyr indicates that the disks of these galaxies contain an underlying old population. The CMDs and luminosity functions reveal significant galaxy-to-galaxy variations in stellar content. Star-forming activity in the central arcmin of NGC300 has been suppressed for the past Gyr with respect to disk fields at larger radii. Nevertheless, comparisons between fields within each galaxy indicate that star-forming activity during intermediate epochs was coherent on spatial scales of a kpc or more. A large cluster of stars, which isochrones suggest has an age near 100 Myr, is seen in one of the NGC55 fields. The luminosity function of the brightest stars in this cluster is flat, as expected if a linear luminosity-core mass relation is present. ",
        "introduction": "Recent studies of the structural characteristics of spiral galaxies suggest that the disk and spheroid do not evolve in isolation, but interact throughout the lifetime of a galaxy (e.g. Andredakis, Peletier, \\& Balcells 1995). The observational signatures of these interactions are readily apparent in late-type spiral galaxies. For example, Courteau, de Jong, \\& Broeils (1996) find a scale-free Hubble sequence among late-type spirals, a result which could be explained if the central light concentrations \\footnote{ The presence of a traditional bulge in late-type spirals galaxies has been challenged by Bothun (1992) and Regan \\& Vogt (1994), who concluded that the `bulge' in M33 is actually the central extension of the halo. However, Minniti, Olszewski, \\& Rieke (1993) and Mighell \\& Rich (1995) resolved the innermost regions of M33 into stars, and detected a significant intermediate age population. Minniti {\\it et al.} (1993) argue that, on the basis of stellar content alone, M33 contains a central component that is distinct from the halo. Given this debate, in the current paper the term `central light concentration' is used to refer to what has traditionally been called the `bulge' in late-type spirals.} formed after disks. WFPC1 images discussed by Phillips {\\it et al.} (1996) reveal that the central light concentrations of late-type spirals are structurally distinct from the bulges of early-type spirals, suggesting differences in evolutionary histories.  Surveys of the bright stellar content in nearby galaxies provide a direct means of studying the evolution of their central regions. Photometric studies of first ascent and asymptotic giant branch (AGB) populations are of particular interest, as these stars probe evolution during early and intermediate epochs, when the basic properties of the disk and spheroid were being imprinted. While efforts to resolve the central light concentrations of these systems into stars require near-diffraction limited image quality to overcome crowding, the inner disks of many nearby systems can easily be resolved into stars from the ground. Although the Local Group contains the closest, most obvious objects for stellar content surveys, the number of targets is limited to three morphologically diverse spiral galaxies: the Milky-Way, M31, and M33. This limited sample makes it necessary to study more distant systems and, as the nearest collection of galaxies outside the Local Group, the Sculptor Group offers a number of lucrative targets that can be resolved from the ground.  In the current paper, deep $J, H$, and $K$ images are used to investigate the photometric properties of cool stars in the inner disks of the Sculptor galaxies NGC55, NGC300, and NGC7793. The morphological types and integrated brightnesses of these galaxies, as assigned by Sandage \\& Tammann (1987), are summarized in Table 1. Throughout this study it is assumed that these galaxies are equidistant with $\\mu_0 = 26.0$, as derived for NGC300 by van den Bergh (1992) from a number of different standard candles. The Galactic reddening towards these galaxies is negligible (Burstein \\& Heiles 1984).  There are a number of advantages to conducting photometric surveys of luminous, cool evolved stars at wavelengths longward of $1\\mu$m. Not only is the contrast between bright cool stars and fainter unresolved objects in the disk enhanced at infrared wavelengths, but it is also possible to overcome the effects of line blanketing, which can affect the spectral energy distributions of moderately metal-rich giants at optical wavelengths (e.g. Bica, Barbuy, \\& Ortolani 1991), and complicate efforts to derive bolometric corrections. In addition, near-infrared two-color diagrams can also be used to identify foreground stars, contamination from which may be significant at faint visible magnitudes.  NGC55 and NGC300 have been the targets of earlier photometric investigations. Deep broad- and narrow-band surveys of NGC55 (Pritchet {\\it et al.} 1987) and NGC300 (Richer, Pritchet, \\& Crabtree 1985; Zijlstra, Minniti, \\& Brewer 1996) have revealed that the outer disks of these galaxies contain rich AGB populations. A comparison of the AGB luminosity functions suggests that the star-forming histories of NGC55 and NGC300 during intermediate epochs were similar, but not identical (Pritchet {\\it et al.} 1987). Freedman (1984) and Pierre \\& Azzopardi (1988) used $B$ and $V$ photometry to survey the bright young stellar content of NGC300, while Kiszhurno-Kozrey (1988) used CCD observations to construct $(V, B-V)$ CMDs of two fields in NGC55. There is no published photometric study of the stellar content of NGC7793, although Catanzarite {\\it et al.} (1995) report the discovery of Cepheids in this galaxy. The only published infrared survey of these galaxies was carried out by Humphreys \\& Graham (1986), who obtained $JHK$ aperture measurements of red supergiant (RSG) candidates in NGC300. Spectroscopy revealed that almost half of the candidate objects were cool Galactic main sequence stars.  The paper is structured as follows. The observations, reduction techniques, and methods used to measure stellar brightnesses are discussed in \\S 2. The luminosity functions (LFs), two-color diagrams (TCDs), and color-magnitude diagrams (CMDs) derived from these data are presented and compared in \\S 3. In \\S 4 the data are used to search for radial population gradients in NGC300 and NGC7793. A summary of the results follows in \\S 5. ",
        "conclusions": ""
    },
    "9803/astro-ph9803104_arXiv.txt": {
        "abstract": "We report the results of a 5-GHz southern-hemisphere snapshot VLBI observation of a sample of blazars. The observations were performed with the Southern Hemisphere VLBI Network plus the Shanghai station in 1993 May. Twenty-three flat-spectrum, radio-loud sources were imaged. These are the first VLBI images for 15 of the sources. Eight of the sources are EGRET ($>$~100~MeV) $\\gamma$-ray sources. The milliarcsecond morphology shows a core-jet structure for 12 sources, and a single compact core for the remaining 11. No compact doubles were seen. Compared with other radio images at different epochs and/or different frequencies, 3 core-jet blazars show evidence of bent jets, and there is some evidence for superluminal motion in the cases of 2 blazars. The detailed descriptions for individual blazars are given. This is the second part of a survey: the first part was reported by Shen \\etal (1997). ",
        "introduction": "Blazar is the collective name for BL Lac objects, optically violent variables and highly polarized quasars, all of which share extreme observational properties that distinguish them from other active galactic nuclei. These properties include strong and rapid variability, high optical polarization, weak emission lines, and compact radio structure (cf. Impey 1992). About 200 blazars have been identified (cf. Burbidge \\& Hewitt 1992). A possible explanation for the blazar phenomenon within a unified scheme for active galactic nuclei is that their emission is beamed by the relativistic motion of the jets traveling in a direction close to the observer's line of sight. This beaming argument is strengthened by the recent CGRO (Compton Gamma Ray Observatory) discovery that most of the detected high-latitude $\\gamma$-ray sources are blazars (e.g. Dondi \\& Ghisellini 1995). A comprehensive theoretical review of these sources has been made by Urry \\& Padovani (1995).  Blazars are an important class of active galactic nuclei because they are thought to be sources with relativistic jets seen nearly end-on. Such sources generally have very compact, flat-spectrum radio cores, which are appropriate for VLBI study.  Pearson \\& Readhead (1988) have undertaken a survey of a complete sample consisting of 65 strong northern-hemisphere radio sources. They provided the first well-defined morphological classification scheme, based primarily on the large-scale radio structure and radio spectra of the sources.  Most surveys to date, however, including the recent Caltech--Jodrell Bank VLBI Surveys (Polatidis \\etal 1995; Thakkar \\etal 1995; Xu \\etal 1995; Taylor \\etal 1994; Henstock \\etal 1995), have been restricted to northern-hemisphere sources. For example, all the confirmed superluminal radio sources, except the well-known equatorial source 3C~279 (1253$-$055), are in the northern sky (Vermeulen \\& Cohen 1994). This reflects the paucity of southern VLBI observations, the notable exceptions being the systematic one-baseline surveys by Preston \\etal (1985) and Morabito \\etal (1986), and the more extensive SHEVE survey (Preston \\etal 1989 and references therein).  Since 1992 we have been carrying out a program to address this deficiency, using VLBI at 5 GHz to study southern radio sources. In an earlier paper (Shen \\etal 1997, hereafter Paper I), we reported the results from the first observing session in 1992 November, and presented images of 20 strong sources selected on the basis of their correlated fluxes on intercontinental baselines. In 1993 May we observed a second sample of southern sources, which is the subject of this paper.  Section~2 introduces this blazar sample; Section~3 briefly describes the observations and data reduction procedures; Section~4 presents the results; the summary and conclusions are presented in Section~5.  Throughout the paper, we define the spectral index, $\\alpha$, by the convention {S$_{\\small \\nu}\\,\\propto\\,{\\nu}^{\\alpha}$}, and assume {H$_0=100$~km~s$^{-1}$~Mpc$^{-1}$} and q$_0=0.5$. ",
        "conclusions": "In this paper we have defined a sample of southern-hemisphere core-dominated blazars.  Of the 24 blazars in the sample, 3 were observed earlier with the same array. The other 21 in the sample and 2 other sources were observed in 1993 May with the Southern VLBI Network plus the Shanghai radio telescope. This is part of the Southern Hemisphere 5-GHz VLBI Survey project, the aim of which is to improve the study of southern extragalactic radio sources (see Paper I). Our study also adds significantly to the number of sources whose structures can be compared on arcsecond (kpc) and milliarcsecond (pc) scales (Table~5). The misalignment of jet-like structures on these scales is an important unsolved problem for the understanding of compact sources.  The main conclusions presented in this paper can be summarized as follows:  \\begin{enumerate} \\item We have detected and imaged all 23 radio sources, of which 15 are first-epoch VLBI images. These are PKS~0118$-$272, PKS~0332$-$403, PKS~0426$-$380, PKS~0454$-$234, PKS~0823$-$223, PKS~1034$-$293, PKS~1244$-$255, PKS~1514$-$241, PKS~1936$-$155, PKS~1954$-$388, PKS~2005$-$489, PKS~2155$-$152, PKS~2240$-$260, PKS~2243$-$123 and PKS~2355$-$534. \\item Most of the blazars are resolved and display simple morphology, with 12 having core-jet structures and 11 having single-core structures. Observations with increased sensitivity will probably reveal many more core-jet structures (e.g. 2243--123). We have compared our VLBI images with other radio images. Only 3 (PKS~0438$-$436, PKS~0537$-$441 and PKS~1226+023) of the 12 core-jet blazars were found to have curved jets. Superluminal motion was inferred from two-epoch observations for 2 sources (PKS~0208$-$512, PKS~2243$-$123). \\item Eight of these blazars (PKS~0208$-$512, PKS~0454$-$234, PKS~0521$-$365, PKS~0537$-$441, PKS~1127$-$145, PKS~1226+023, PKS~1424$-$418 and PKS~2005$-$489) have been detected at $>100$~MeV $\\gamma$-ray energies. Together with the other 5 EGRET sources observed in 1992 November (Paper I), a total of 13 southern $\\gamma$-ray-loud blazars have now been imaged by our survey project. A systematic study of the VLBI properties of these $\\gamma$-ray blazars and comparison with other non-$\\gamma$-ray sources will improve our understanding of the beaming characteristics in blazars and the properties of EGRET sources. \\end{enumerate}"
    },
    "9803/astro-ph9803332_arXiv.txt": {
        "abstract": "A quantitative method is presented to compare observed and synthetic colour-magnitude diagrams (CMDs). The method is based on a $\\chi^2$ merit function for a point $(c_i,m_i)$ in the observed CMD, which has a corresponding point in the simulated CMD within $n\\sigma(c_i,m_i)$ of the error ellipse. The $\\chi^2$ merit function is then combined with the Poisson merit function of the points for which no corresponding point was found within the $n\\sigma(c_i,m_i)$ error ellipse boundary.\\hfill\\break Monte-Carlo simulations are presented to demonstrate the diagnostics obtained from the combined ($\\chi^2$, Poisson) merit function through variation of different parameters in the stellar population synthesis tool. The simulations indicate that the merit function can potentially be used to reveal information about the initial mass function. Information about the star formation history of single stellar aggregates, such as open or globular clusters and possibly dwarf galaxies with a dominating stellar population, might not be reliable if one is dealing with a relatively small age range. ",
        "introduction": "In the last decade the simulation of synthetic Hertzsprung-Russell and colour-magnitude diagrams (hereafter respectively referred to as HRDs and CMDs) has advanced at a rapid pace. It has been applied successfully in various studies of young clusters in the Magellanic Clouds (Chiosi \\etal\\ 1989; Bertelli \\etal\\ 1992; Han \\etal\\ 1996; Mould \\etal\\ 1997; Vallenari \\etal\\ 1990, 1992, and 1994ab), dwarf galaxies in the Local Group (Aparicio{\\muspc\\&\\muspc}Gallart 1995, 1996; Aparicio \\etal\\ 1996, 1997{\\mmuspc}a,b; Ferraro \\etal\\ 1989; Gallart \\etal\\ 1996a{\\to}c; Han \\etal~1997; Tolstoy~1995, 1996; Tosi \\etal~1991), open clusters in our Galaxy (Aparicio \\etal~1990; Carraro \\etal~1993, 1994; Gozzoli \\etal~1996), and the structure of our Galaxy (Bertelli \\etal~1994, 1995, 1996; Ng 1994, 1997{\\mmuspc}a,b; Ng{\\muspc\\&\\muspc}Bertelli 1996{\\mmuspc}a,b; Ng \\etal~1995, 1996{\\mmuspc}a,b, 1997). \\par Generally all studies focus in the first place on matching the morphological structures at different regions in the CMDs (Gallart 1998). In the recent years a good similarity is obtained between the observed and the simulated CMDs. Unfortunately, the best fit is in some cases distinguished by eye. The morphological differences are large enough to do this and the eye is actually guided by a detailed knowledge of stellar evolutionary tracks. However, the stellar population technique has improved considerably and an objective evaluation tool is needed, to distinguish quantitatively one model from another. \\par Bertelli \\etal~(1992,\\muspc1995) and Gallart \\etal~(1996c) defined ratios to distinguish the contribution from different groups of stars. The ratios are defined so that they are sensitive to the age, the strength of the star formation burst and/or the slope of the initial mass function. Vallenari \\etal~(1996{\\mmuspc}a,b) demonstrated that this is a good method to map the spatial progression of the star formation in the Large Magellanic Cloud. Robin \\etal~(1996), Han \\etal~(1997) and Mould \\etal~(1997) use a maximum likelihood method to find the best model parameter(s), while Chen (1996{\\mmuspc}a,b) adopted a multivariate analysis technique. Different models can be quantitatively sampled through Bayesian inference (Tolstoy 1995; Tolstoy{\\muspc\\&\\muspc}Saha 1996) or a chi-squared test (Dolphin 1997). \\par In principle one aims with stellar population synthesis to generate a CMD which is identical to the observed one. The input parameters reveal the evolutionary status of the stellar aggregate under study. To obtain a good similarity between the observed and the simulated CMD one needs to implement in the model in the first place adequately extinction, photometric errors and crowding. It goes without saying that the synthetic population ought to be comparable with the age and metallicity (spread) of the stellar aggregate. Only with a proper choice of these parameters, one can start to study in more detail the stellar initial mass function and the star formation history for an aggregate. \\par This paper describes a quantitative evaluation method based on the combination of a chi-squared and Poisson merit function% \\footnote{The method introduced in Sect.~2 actually mimics closely the procedure used in the `fit by eye' method.}. It allows one to select the best model from a series of models. In the next section this method is explained and Monte-Carlo simulations are made, to display how the diagnostic diagrams of the residual points can be employed. It is demonstrated that the non-fitting, residual points provide a hint about the parameter that needs adjustment in order to improve the model. This paper ends with a discussion about the method and the diagnostic diagrams together with their limitations. \\hfill\\break It is emphasized that this paper deals with a description of a quantitative evaluation method for CMDs. Aspects related with the implementation of an automated CMD fitting program and comparison with simulated or real data sets will not be considered, because they are not relevant for the general validity of the method described in the following section. ",
        "conclusions": "A method based on the $\\chi^2$ merit function is presented to compare `observed' with synthetic single stellar populations. Monte-Carlo simulations have been performed to display the diagnostic power from a CMD containing the points for which no corresponding synthetic point was found within a reasonable error ellipse. The simulations indicate that the CMDs of residual points might provide hints about model parameters to be improved. The simulations further indicate that one ought to be cautious with the analysis of stellar luminosity functions and that strong hints can be obtained about the shape of the initial mass function."
    },
    "9803/astro-ph9803326_arXiv.txt": {
        "abstract": "We present WFPC2--HST photometry of the resolved stellar population in the post-starburst galaxy NGC~1569. The color-magnitude diagram (CMD) derived in the F439W and F555W photometric bands contains $\\sim$2800 stars with photometric error $\\leq$ 0.2 mag down to \\mb, \\mv~$\\simeq$~26, and is complete for \\mv~$\\lsim$~23. Adopting the literature distance modulus and reddening, our CMD samples stars more massive than $\\sim$4~\\MSUN, allowing us to study the star formation (SF) history over the last $\\sim$0.15~Gyr. The data are interpreted using theoretical simulations based on stellar evolutionary models. The synthetic diagrams include photometric errors and incompleteness factors. Testing various sets of tracks, we find that the ability of the models to reproduce the observed features in the CMD is strictly related to the shape of the blue loops of the sequences with masses around 5~\\MSUN. The field of NGC~1569 has experienced a global burst of star formation of duration $\\gsim$0.1~Gyr, ending $\\sim$ 5$-$10~Myr ago. During the burst, the SF rate was approximately constant, and, if quiescent periods occurred, they lasted less than $\\sim$10~Myr. The level of the SF rate was very high: for a single-slope initial mass function (IMF) ranging from 0.1 to 120~\\MSUN\\ we find values of 3, 1, and 0.5~\\Myr\\ for $\\alpha = 3$, 2.6, and 2.35 (Salpeter), respectively. When scaled for the surveyed area, these rates are approximately 100 times larger than found in the most active dwarf irregulars in the Local Group. The data are consistent with a Salpeter IMF, though our best models indicate slightly steeper exponents. We discuss the implications of our results in the general context of the evolution of dwarf galaxies. ",
        "introduction": "Understanding the status and evolution of dwarf galaxies is of crucial importance in current astrophysics and cosmology. They are important ingredients in common scenarios of galaxy formation, either as building blocks or as left-overs of the formation process (Silk 1987). In addition, dwarf galaxies forming at redshift around and below 1 have been suggested to be responsible for the excess of faint blue galaxies seen in deep photometric surveys (Babul $\\&$ Rees 1992; Babul $\\&$ Ferguson 1996).  Dwarf galaxies come in three main classes: (i) dwarf spheroidals, which contain population~II stars and have low surface brightness (e.g., Wirth \\& Gallagher 1984); (ii) dwarf irregulars, with various levels of star-formation activity and a high gas mass (Hunter \\& Gallagher 1986); (iii) starbursting dwarfs, including blue compact dwarfs (Thuan 1991), H~II galaxies (Terlevich \\etal\\ 1991), and blue amorphous galaxies (Sandage \\& Brucato 1979). The boundary between dwarf and giant galaxies is usually taken to be around $M_{\\rm{B}} = -16$ mag (Tammann 1994).  An intriguing property of dwarf galaxies is the pronounced dichotomy between the two main classes of dwarf spheroidals and dwarf irregulars: most dwarfs are either gas-poor spheroids or gas-rich disk-like systems, with very few transition objects (van den Bergh 1977). The relationship between the two classes is still poorly understood: attempts to unify the complex zoo of dwarf galaxies have not been fully successful (Binggeli 1994). Some effort has been made to understand the interrelationship between individual classes in terms of an evolutionary sequence (e.g., Gallagher, Hunter, \\& Tutukov 1984). Starbursting dwarfs can play a key role: a powerful starburst might strip a gas-rich dwarf irregular of its gas and transform it into a gas-poor dwarf-elliptical (Kormendy 1985; Marlowe 1997).  Understanding the star-formation histories in starbursting dwarf galaxies is a prerequisite for understanding evolutionary connections between different object classes (e.g., Gallagher 1996) since during cosmologically brief periods of time their star formation increases dramatically. This has profound implications for the interstellar medium. Stellar winds and supernovae inject energy and may cause a `blow-out', initiating a galactic superwind (Heckman 1995). Numerous examples of galactic superwinds are known from observations (e.g., Heckman, Armus, \\& Miley 1990) and their cosmological implications have been discussed (e.g., De Young \\& Heckman 1994). However, attempts to relate the superwind properties and the star-formation histories in individual galaxies are quite rare and do not go beyond a qualitative level (Heckman \\etal\\ 1990; Leitherer, Robert, \\& Drissen 1992). Yet establishing such a relation is crucial for theoretical models of the superwind hydrodynamics (Suchkov \\etal\\ 1994; Tenorio-Tagle \\& Mu\\~noz-Tu\\~non 1997).  NGC~1569 (= UGC~3056 = Arp~210 = VII~Zw~16 = IRAS~4260+6444) has been indicated by Gallagher, Hunter, $\\&$ Tutukov (1984) as an outstanding object in their sample of active star forming galaxies. Indeed, it is a prime candidate for a detailed study of the star-formation history in a starburst galaxy since at $D = 2.2 \\pm 0.6$~Mpc it is the closest starburst galaxy known (Israel 1988a). Adopting the literature reddening $E(B-V)$ = 0.56 (Israel 1988a), photometry of individual stars is feasible with HST down to $M_{V,0}\\simeq -1.5$. This corresponds to stars with mass of about 3~\\Ms\\ in their core-helium burning phase, and consequently to a look-back time of $\\sim$0.4~Gyr (e.g., Schaller \\etal\\ 1992).  The properties of NGC~1569 are typical of a dwarf galaxy. With a distance modulus of $(m-M)_0$ = 26.71, its total absolute B magnitude is $M_{B.0} \\sim -17$ (Israel 1988a), which is intermediate between the two Magellanic Clouds. Its total mass and hydrogen content are estimated $M \\simeq 3.3 \\times 10^8$~\\Ms\\ and $M_H \\simeq 1.3 \\times 10^8$~\\Ms, respectively (Israel 1988a). Numerous studies of the chemical composition exist. The published range of oxygen abundances is very narrow: $12 + \\log (O/H) = 8.25$ (Hunter, Gallagher, \\& Rautenkranz 1982), 8.37 (Calzetti, Kinney, \\& Storchi-Bergmann 1994), 8.26 (Devost, Roy, \\& Drissen 1997), 8.29 (Gonz\\'alez-Delgado et al. 1997), and 8.19 (Kobulnicky \\& Skillman 1997; Martin 1997). Searches for chemical composition gradients were done by Devost et al. and Kobulnicky \\& Skillman. No significant evidence for chemical inhomogeneities was found. If [O/Fe]~=~0.0 is assumed, the average metallicity of NGC~1569 is $Z \\simeq 0.25$~\\Zsun, with about 0.2~dex uncertainty. Thus, NGC~1569 is a gas-rich system with SMC-like composition in a relatively early stage of its chemical evolution.  The most detailed study of the warm ($\\sim$10$^4$~K) gas of NGC~1569 was done by Waller (1991), who found evidence for a link between the star-formation history and the morphology and kinematics of extended gas structures. The same connection is suggested by hot ($\\sim$10$^7$~K) gas observed with the ROSAT and ASCA satellites (Heckman \\etal\\ 1995; Della Ceca \\etal\\ 1996). The extended X-ray emission is consistent with a starburst driven galactic superwind which could in principle lead to a large-scale disruption of the interstellar medium. Evidence for large numbers of supernovae was also found from the non-thermal radio spectrum (Israel 1988b).  One of the most spectacular properties of NGC~1569 are its `super star clusters'. These high-density star clusters were first detected and discussed by Arp \\& Sandage (1985) and Melnick \\& Moles, \\& Terlevich (1985) and later studied by O'Connell, Gallagher, \\& Hunter (1994), Ho \\& Filippenko (1996), De~Marchi \\etal\\ (1997), and Gonz\\'alez-Delgado \\etal\\ (1997). The super star clusters reflect a recent starburst in which at least $10^5$~\\Ms\\ of gas was transformed into stars within about 1~pc. Super star clusters are similar to young Galactic globular clusters observed shortly after their formation (Meurer 1995).  Pre-Costar HST studies of the field population were done by O'Connell \\etal\\ (1994) and Vallenari \\& Bomans (1996) (hereinafter VB). According to the latter authors, NGC~1569 experienced a strong starburst until a few Myr ago, when its activity subsided. The repaired HST offers the possibility of significant improvement over these previous studies. We have therefore embarked on an extensive spectroscopic and photometric HST study in an attempt to confront the observational data with the predictions of theoretical models, and derive detailed informations on the star formation (SF) history in NGC~1569. In a first paper (De Marchi \\etal\\ 1997) we discussed the photometry of the SSCs, while in this paper we present the results for the resolved stars in the same field, which are interpreted in terms of the recent SF history via theoretical simulations. The data and their analysis are described in Section~2. The derived color-magnitude diagram and luminosity function are in Section~3. Our theoretical models are described in Section~4. They are compared with the observations in Section~5. In Section~6 we discuss our results and their implications for the evolution of dwarf galaxies, and we present in Section~7 our general conclusions. ",
        "conclusions": "In this paper we have presented HST WFPC2 F555W and F439W photometry of resolved stars in the dwarf irregular galaxy NGC~1569. About 2800 stars were measured with photometric error smaller than \\sigmada = 0.2\\,mag out of a total of $\\sim$7000 objects detected. The corresponding CMD extends down to \\mb, \\mv~$\\simeq$~26, and is complete for \\mv~$\\lsim$~23. It samples stars more massive than $\\sim$4~\\Ms\\, allowing us to derive the SF history over the last 0.15~Gyr.  The interpretation of the data has been performed via theoretical simulations, based on stellar evolutionary tracks. We have considered several sets of stellar tracks from the Geneva and Padova groups, and were able to reach a satisfactory representation of the data adopting either the Geneva tracks with $Z$~=~0.001, or the Padova tracks with $Z$~=~0.004. Since the applicability of these tracks is related to the shape of the blue loops, which is sensitive to details in the input parameters used in the computations, we consider it premature to dismiss alternative sets of evolutionary tracks.  The two sets provide a similar picture of the recent SF history in the field of NGC~1569. The galaxy has experienced a general burst of SF, of duration $\\Delta t \\gsim 0.1$~Gyr, at an approximately constant rate. The burst seems to have stopped $\\sim$5~--~10~Myr ago but SF continues in the HII regions and in the SSCs. If this general burst consists of successive episodes, these must have occurred at a similar rate, and be separated by short quiescent periods. Qualitatively, this behavior looks similar to that inferred for dwarf irregulars in the Local Group. Quantitatively, though, the level of the SFR in this recent burst in NGC~1569 is approximately 2 orders of magnitude higher.  We find that  the Salpeter IMF is consistent with the observed CMD and LF, but our best simulations are characterized by slightly steeper slopes (2.6 -- 3). This seems to disagree with observations of other starbursts which generally indicate a Salpeter slope (Leitherer 1997). The present study, however, refers to the field population of stars less massive than $\\sim$30~\\Ms, whereas most starburst IMFs are derived for the nuclear burst and for dense OB clusters in the mass range above $\\sim$20~\\Ms. Since most of the SF activity in the NGC~1569 field has subsided about 7~--~10~Myr ago, we have no direct information on the IMF of those stars with lifetimes less than this value. The corresponding minimum masses are around 30~\\Ms. The field-star IMF in our Galaxy and in the Magellanic Clouds is steeper (biased against massive stars) than the cluster IMF (Massey, Johnson, $\\&$ DeGioia-Eastwood 1995a; Massey \\etal\\ 1995b), with the cluster value being close to Salpeter. This could be due to a richness effect, where the most massive stars are not observed simply because of their small expected numbers. NGC~1569 may be another case where the most massive stars with masses above $\\sim$30~\\Ms\\ are preferentially formed in clusters, but we emphasize the different mass range sampled by our and that of LMC/SMC studies. Our preferred IMF is quite similar to that derived for field stars of comparable mass in our Galaxy. This is in agreement with the notion of a virtually constant IMF, at least as far as the shape is concerned, and does not support the expectations of flatter slopes in starbursting regions (e.g., Padoan, Nordlund, $\\&$ Jones 1997), or in low metallicity environments, as sometimes invoked to better reproduce the properties of elliptical galaxies (e.g., Vazdekis \\etal\\ 1996).  In the past, the SF in NGC~1569 is likely to have proceeded at a substantially lower average rate than in the recent burst. This follows from the estimate of the gas exhaustion timescale, and from the relatively low metallicity of the ISM in this galaxy. This leaves the possibilities of either an approximately constant SFR, or short strong bursts, with long interburst periods. Detailed chemical evolution models are needed to explore quantitatively the possible evolutionary paths for NGC~1569, and this will be the subject of a forthcoming paper.  The recent burst of SF is certainly not the first episode in NGC~1569 (see also VB) but may be the last. The velocity field of the ionized gas (Tomita, Ohta, $\\&$ Saito 1994) and energetics arguments support this possibility (Heckman \\etal\\ 1995). If this was the case, in the future NGC~1569 may turn into a dwarf spheroidal galaxy. In addition, its luminosity evolution would be characterized by passive fading, which seems required by the Babul $\\&$ Ferguson (1996) bursting dwarfs model to explain the excess faint galaxy counts. However, it is also possible that the outflowing gas falls back onto the galaxy, triggering a successive burst of SF, an option which has to be explored with the computation of detailed hydrodynamic modelling for bursting dwarf galaxies (D'Ercole $\\&$ Brighenti 1998). In spite of these uncertainties, our results shows that dwarf galaxies are capable of sustaining bursts of SF at large enough rates to be relevant for the interpretation of the faint galaxy counts."
    },
    "9803/astro-ph9803110_arXiv.txt": {
        "abstract": "The velocity distribution $f(\\bvel)$ of nearby stars is estimated, via a maximum-likelihood algorithm, from the positions and tangential velocities of a kinematically unbiased sample of 14\\,369 stars observed by the \\hip\\ satellite. $f$ shows rich structure in the radial and azimuthal motions, $v_R$ and $v_\\varphi$, but not in the vertical velocity, $v_z$: there are four prominent and many smaller maxima, many of which correspond to well known moving groups. While samples of early-type stars are dominated by these maxima, also up to about a quarter of red main-sequence stars are associated with them. These moving groups are responsible for the vertex deviation measured even for samples of late-type stars; they appear more frequently for ever redder samples; and as a whole they follow an asymmetric-drift relation, in the sense that those only present in red samples predominantly have large $|v_R|$ and lag in $v_\\varphi$ w.r.t.\\ the local standard of rest (LSR). The question arise, how these old moving groups got on their eccentric orbits? A plausible mechanism known from the solar system dynamics which is able to manage a shift in orbit space is sketched. This mechanism involves locking into an orbital resonance; in this respect is intriguing that Oort's constants, as derived from \\hip\\ data, imply a frequency ratio between azimuthal and radial motion of exactly $\\Omega:\\kappa=3:4$.  Apart from these moving groups, there is a smooth background distribution, akin to Schwarzschild's ellipsoidal model, with axis ratios $\\sigma_R: \\sigma_\\varphi:\\sigma_z\\approx1:0.6:0.35$. The contours are aligned with the $v_r$ direction, but not w.r.t.\\ the $v_\\varphi$ and $v_z$ axes: the mean $v_z$ increases for stars rotating faster than the LSR. This effect can be explained by the stellar warp of the Galactic disk. If this explanation is correct, the warp's inner edge must not be within the solar circle, while its pattern rotates with frequency $\\gtrsim13\\kmskpc$ retrograde w.r.t.\\ the stellar orbits. ",
        "introduction": "\\else ",
        "conclusions": "\\else"
    },
    "9803/astro-ph9803260_arXiv.txt": {
        "abstract": "There are now about fifty known radio pulsars in binary systems, including at least five in double neutron star binaries. In some cases, the stellar masses can be directly determined from measurements of relativistic orbital effects. In others, only an indirect or statistical estimate of the masses is possible.  We review the general problem of mass measurement in radio pulsar binaries, and critically discuss all current estimates of the masses of radio pulsars and their companions. We find that significant constraints exist on the masses of twenty-one radio pulsars, and on five neutron star companions of radio pulsars.  All the measurements are consistent with a remarkably narrow underlying gaussian mass distribution, $m=1.35\\pm0.04M_\\odot$. There is no evidence that extensive mass accretion ($\\Delta m\\gsim0.1M_\\odot$) has occurred in these systems. We also show that the observed inclinations of millisecond pulsar binaries are consistent with a random distribution, and thus find no evidence for either alignment or counteralignment of millisecond pulsar magnetic fields. ",
        "introduction": "Neutron stars have been the subject of considerable theoretical investigation since long before they were discovered as astronomical sources of radio and X-ray emission (\\cite{bz34b,ov39,whe66}).  Their properties are determined by the interplay of all four known fundamental forces---electromagnetism, gravitation, and the strong and weak nuclear forces---but neutron stars remain sufficiently simple in their internal structure that realistic stellar modeling can be done. Measurements of their masses and radii (as well as detailed study of their cooling histories and rotational instabilities) provide a unique window on the behavior of matter at densities well above that found in atomic nuclei ($\\rho_{\\rm nuc}\\approx 2.8\\times10^{14}\\mbox{g~cm}^{-3}$). Observations of neutron stars also provide our only current probe of general relativity (GR) in the ``strong-field'' regime, where gravitational self-energy contributes significantly to the stellar mass.  The most precisely measured physical parameter of any pulsar is its spin frequency. The frequencies of the fastest observed pulsars (PSR~B1937+21 at 641.9~Hz and B1957+20 at 622.1~Hz) have already been used to set constraints on the nuclear equation of state at high densities (e.g., \\cite{fidp88}) under the assumption that these pulsars are near their maximum (breakup) spin frequency.  However, the fastest observed spin frequencies may be limited by complex accretion physics rather than fundamental nuclear and gravitational physics. A quantity more directly useful for comparison with physical theories is the neutron star mass.  The basis of most neutron star mass estimates is the analysis of binary motion.  Soon after the discovery of the first binary radio pulsar (\\cite{ht75a}), it became clear that the measurement of relativistic orbital effects allowed extremely precise mass estimates. Indeed, the measurement uncertainties in several cases now exceed in precision our knowledge of Newton's constant $G$, requiring masses to be quoted in solar units $GM_\\odot$ rather than kilograms if full accuracy is to be retained.  After several recent pulsar surveys, there are now about fifty known binary radio pulsar systems, of which five or six are thought to contain two neutron stars.  It is thus possible for the first time to consider compiling a statistically significant sample of neutron star masses.  It is our purpose here to provide a general, critical review of all current estimates of stellar masses in radio pulsar binaries. The resulting catalog, with a careful, uniform approach to measurement and systematic uncertainties, should be of value both to those who wish to apply mass measurements to studies of nuclear physics, GR, and stellar evolution, and as a guide to the critical observations for observational pulsar astronomers. We begin with a discussion of known methods for pulsar mass determination (\\S\\ref{sec:methods}), including a new statistical technique for estimating the masses of millisecond pulsars in non-relativistic systems.  In \\S\\ref{sec:estimates} we review all known mass estimates, including new data and analysis where possible.  Statistical analysis of the available pulsar mass measurements is presented in \\S\\ref{sec:stat}. We summarize in \\S\\ref{sec:summ}.  A second paper will consider mass estimates for neutron stars in X-ray binary systems (\\cite[Paper~II]{ct98}). A detailed discussion of the implications of the combined results of this work and Paper~II for studies of supernovae and neutron star formation, mass transfer in binary evolution, the nuclear equation of state, and GR will occur elsewhere (Paper~III). ",
        "conclusions": "\\label{sec:stat}  For a dozen neutron stars, useful mass constraints are available with no assumptions beyond the applicability of the general relativistic equations of orbital motion to binary pulsar systems.  Ten of these stars are members of double neutron star binaries.  With the possible exception of PSR~B2127+11C, in the globular cluster M15, the pulsar in each system is believed to have undergone a short period of mass accretion during a high-mass X-ray binary phase ($\\Delta m\\sim10^{-3}M_\\odot$, Taam and van den Heuvel 1986).  The companion stars have not undergone accretion; their masses most directly preserve information about the initial mass function of neutron stars.  Only two ``millisecond'' pulsars, the end products of extended mass transfer in low-mass X-ray binaries, have interesting mass estimates based on GR alone: PSRs~B1802$-$07 and B1855+09.  Because such pulsars must accrete $\\sim0.1M_\\odot$ to reach millisecond periods (\\cite{tv86}), and much more ($\\sim0.7M_\\odot$) in some field decay models (e.g., \\cite{vb95a}), obtaining additional mass measurements of millisecond pulsars is of particular interest in testing evolutionary models and in locating the maximum neutron star mass.  As noted in \\S\\ref{sec:cmrr}, the $P_b$--$m_2$ relation can be used to estimate the companion mass in recycled binary systems with circular orbits and orbital periods $P_b\\gtrsim 3$~d.  There are now thirteen such millisecond ($P_{\\rm spin}<10$~ms) pulsars known, excluding those in globular clusters (where gravitational interactions may have significantly perturbed the orbital parameters since spin-up).  In each case, the measured mass function and the inferred companion mass, together with the requirement that $\\sin i<1$, then yields an upper limit on the mass of the pulsar itself.  A number of systems in which this upper limit is particularly constraining have been mentioned in \\S\\ref{sec:nswd}.  Additional constraints on the neutron star mass in these systems can be derived using statistical arguments, given a prior assumption about the distribution of binary inclinations.  The simplest such assumption is that the binaries are randomly oriented on the sky, though biases toward high or low inclinations are possible in some models (\\S\\ref{sec:uniform}). However, as discussed below, we believe there is currently no evidence for such a bias, so for the remainder of this discussion we assume random orbital orientations.  For an individual system, we are interested in the probability distribution\\footnote{We adopt the notation that $p(x ; A)$ is the (marginal) probability density for the random variable $x$, where $x$ depends upon the parameter $A$.  Also, $p(x|y; A)$ is the conditional probability density for the random variable $x$ for a given value of the random variable $y$ and parameter $A$.}  $p(m_1; f, P_b)$ for the neutron star mass $m_1$ given the measured mass function $f$ and binary period $P_b$.  We can neglect the measurement uncertainty in $f$ and $P_b$.  Then, the probability distribution for $m_1$ can be written schematically as \\begin{equation} p(m_1; f, P_b) = \\int_0^1 d(\\cos i) \\int_{m_{2,{\\rm min}}(P_b)}^{m_{2,{\\rm max}}(P_b)} dm_2\\, p(m_2; P_b)\\, p(\\cos i)\\, p(m_1|m_2,\\cos i; f)  , \\end{equation} where $m_1=f^{-1/2}(m_2 \\sin i)^{3/2}-m_2$ is restricted to positive values.  We have evaluated this numerically for each system, assuming that $p(\\cos i)$ is uniform between zero and unity and that $p(m_2; P_b)$ is uniformly distributed within the appropriate factor (see \\S\\ref{sec:cmrr}) of the $m_2$ implied by equations (9)--(11).  Not surprisingly, the width of the distribution $p(m_1; f, P_b)$ is dominated by the range of allowed $\\cos i$ rather than the uncertainty in $m_2$ for a given $P_b$.  For each of the 13 binaries, we have plotted the cumulative distribution CDF$(m_1)=\\int_0^{m_1} p(m_1')\\, dm_1'$ in Figure~\\ref{fig:cumprob}.\\placefigure{fig:cumprob}  The median and 68\\% and 95\\% confidence regions for each pulsar mass is given in Table~\\ref{tab:pbm2}.  \\placetable{tab:pbm2} Although several of the pulsars have, under the assumptions made, most likely masses well above $2M_\\odot$, some such results are expected even if all the masses are quite low.  In fact, in only one case of the 13 pulsars does $1.35M_\\odot$ lie outside the 95\\% central confidence region (J1045$-$4509), and in 6 cases of 13 is $1.35M_\\odot$ excluded at 68\\% confidence, consistent with chance.  It is interesting to ask whether a single, simple distribution of neutron star masses is consistent with all of our observational constraints.  We considered two models for this question: a Gaussian distribution of masses with mean $\\hat{m}$ and standard deviation $\\sigma$, and a uniform distribution of masses between $m_l$ and $m_u$ ({\\it cf.} Finn 1994).  A maximum likelihood analysis was used to estimate the parameters $\\hat{m}$, $\\sigma$, $m_l$, and $m_u$ (assuming a uniform prior distribution for all four parameters).  The resulting 68\\% and 95\\% joint confidence limits on $\\hat{m}$ and $\\sigma$ are shown in Figure~\\ref{fig:all26}, and on $m_l$ and $m_h$ in Figure~\\ref{fig:all26uni}.  \\placefigure{fig:all26} \\placefigure{fig:all26uni} In each model, the distribution of neutron star masses is remarkably narrow: the maximum likelihood solutions are $\\hat{m}=1.35M_\\odot$ and $\\sigma=0.04M_\\odot$, and $m_l=1.26M_\\odot$ and $m_u=1.45M_\\odot$.  Of course, {\\em any} model (even a poor one) will yield maximum likelihood parameters for a given data set.  However, it is obvious by inspection that both the Gaussian and uniform distributions for the neutron star mass are good fits to the extremely narrow observed range of neutron star masses in the double neutron star binaries.  While it is difficult to quantify the goodness-of-fit for the entire data set, because of the diverse assumptions made in the various mass estimates and the sometimes highly non-gaussian error estimates, we can easily test some neutron star subsamples against the maximum likelihood gaussian model $m_1=1.35\\pm0.04M_\\odot$.  For the thirteen neutron-star--white-dwarf binaries, we used a Monte Carlo technique to evaluate the fit quality.  For each binary (with its measured $P_b$ and $f$), we simulated a large number of Monte Carlo trials where the neutron star mass $m_1$ was drawn from the maximum likelihood model, $m_2$ was drawn from the appropriate uniform distribution implied by $P_b$, and $\\cos i$ was drawn from a uniform distribution.  The Monte Carlo trials were then used to construct the probability distribution for the mass function, and this distribution was used to compute the cumulative probability for the measured mass function, $p(f^\\prime<f)$.  If the model and the associated assumptions are correct, then the cumulative probabilities for the 13 measured mass functions should be consistent with a uniform distribution betwen zero and unity ({\\it cf.} a classical $V/V_{\\rm max}$ test, Schmidt 1968). A KS test of the distribution shows consistency with a uniform distribution at the improbably good 99\\% level (Figure~\\ref{fig:ks}). \\placefigure{fig:ks}  For the ten stars for which gaussian error estimates $\\sigma_e$ are available (both stars in the relativistic binaries B1534+12, B1913+16, and B2127+11C, as well as J1012+5307, J1713+0747, B1855+09, and J0045$-$7319), we can calculate a $\\chi^2$ statistic, $\\sum(m-\\hat{m})^2/(\\sigma^2+\\sigma_e^2)=7.5$, consistent with expectations for a chi-square distribution with $10-2$ degrees of freedom.  We conclude, therefore, that at least in the radio pulsar systems, there is no evidence for neutron star masses above about $1.45M_\\odot$.  Indeed, the data appear very well modeled by very narrow distributions centered around $1.35M_\\odot$."
    },
    "9803/astro-ph9803056_arXiv.txt": {
        "abstract": " ",
        "introduction": "Soft gamma repeaters are astrophysical sources which exhibit long periods of quiescence, often spanning years, punctuated by periods of intense bursting activity during which many brief (durations $<1$ s) and intense (luminosities L$\\sim1-10^{3}$ L$_{Edd}$) bursts are emitted by the source (Norris et al. 1991). Besides the burst emission, SGRs are also characterized by steady - also referred to as ``quiescent'' - emission. Believed to be neutron stars, the mechanism(s) for both the steady and bursting X--ray emission is still not well understood (Thompson \\& Duncan 1995).  SGR 1806-20 is the most prolific SGR, and it has been studied in the X--ray (Sonobe et al. 1994), optical (van Kerkwijk et al. 1995), infrared (Kulkarni et al. 1995), and radio (Kulkarni et al. 1994) bands. The source became active again during the Fall of 1996, emitting many powerful bursts that were first detected with BATSE (Kouveliotou et al. 1996). A target of opportunity observation by the {\\it Rossi X--ray Timing Explorer} (RXTE) was initiated on November 5, 1996. The data analyzed here were taken during that $50$ ks observation, which spanned the time interval starting at 10:53:20 UT (5/11/96) and ending at 10:52:00 UT (6/11/96). In addition, followup {\\it RXTE} data were taken during the time interval 06:45:20 UT (15/7/97) to 09:06:24 (15/7/97), in which the source was not bursting. ",
        "conclusions": "Through analysis of {\\it RXTE} TOO data, we have determined that the persistent emission from SGR 1806--20 is consistent with a constant spectral shape and intensity both during and away from the active bursting periods. The spectrum is best-fit by a nonthermal power law shape, with thermal bremsstrahlung and Raymond-Smith functional forms producing much worse fits to the {\\it RXTE} data. The mean power law spectral index obtained by {\\it RXTE} is $2.30\\pm0.02$, which is consistent with the {\\it ASCA} value of $\\alpha=2.2\\pm0.2$ (Sonobe et al. 1994).  The nonthermal nature of the SGR 1806--20 quiescent spectrum supports the idea that the X--ray emission is due to a compact synchrotron nebula, or plerion, that derives its power from either a rapidly spinning-down pulsar (Kulkarni et al. 1994) or energetic particles ejected in the SGR bursts (Tavani 1994). The lack of an iron line in the SGR spectrum (Sonobe 1994) and short term time variability argue against the source being a low luminosity X--ray binary system (White, Nagase, \\& Parmar 1995), although a more extensive monitoring campaign is necessary to rule out this hypothesis.  \\paragraph*{Acknowledgements.} We thank NASA for support under grants NAS5-30720 (D.M. and R.E.R.), NAG5-2560 (S.D. and C.K.), and NAG5-4878 (JvP)"
    },
    "9803/nucl-th9803032_arXiv.txt": {
        "abstract": "The nucleus $^{54}$Mn has been observed in cosmic rays.  In astrophysical environments it is fully stripped of its atomic electrons and its decay is dominated by the $\\beta^-$ branch to the $^{54}$Fe ground state.  Application of $^{54}$Mn based chronometer to study the confinement of the iron group cosmic rays requires knowledge of the corresponding halflife, but its measurement is impossible at the present time. However, the branching ratio for the related $\\beta^+$ decay of $^{54}$Mn was determined recently. We use the shell model with only a minimal truncation and calculate both $\\beta^+$ and $\\beta^-$ decay rates of $^{54}$Mn. Good agreement for the $\\beta^+$ branch suggests that the calculated partial halflife of the $\\beta^-$ decay, $(4.94\\pm 0.06)\\times 10^5$ years, should be reliable.  However, this halflife is noticeably shorter than the range 1-$2\\times 10^6$~y indicated by the fit based on the $^{54}$Mn abundance in cosmic rays. We also evaluate other known unique second forbidden $\\beta$ decays from the nuclear $p$ and $sd$ shells ($^{10}$Be, $^{22}$Na, and two decay branches of $^{26}$Al) and show that the shell model can describe them with reasonable accuracy as well. ",
        "introduction": "The nucleus $^{54}$Mn decays in the laboratory dominantly by electron capture to the $2^+$ state in $^{54}$Cr with the halflife of 312 days. However, as a component of cosmic rays, $^{54}$Mn will be fully stripped of its atomic electrons, and this mode of decay is therefore impossible. The $^{54}$Mn nuclei were in fact detected in cosmic rays using the Ulysses spacecraft~\\cite{Ulysses1,Ulysses2}. They offer an attractive possibility to use their measured abundance as a chronometer for the iron group nuclei (Sc - Ni) in cosmic rays in analogy to the chronometers based on the abundances of other long lived isotopes ($^{10}$Be, $^{26}$Al, and $^{36}$Cl). With them one can, in turn, determine the mean density of interstellar matter, a quantity of considerable interest. The use of the long lived nuclei as cosmic ray chronometers is reviewed in Ref.~\\cite{Simpson}. The importance of $^{54}$Mn for the understanding of propagation of the iron group nuclei that are products of explosive nuclear burning have been stressed in Refs.~\\cite{Grove,Leske}.  For this program to succeed, however, one must know the halflife of the stripped $^{54}$Mn ($I^{\\pi} = 3^+$). The decay scheme of $^{54}$Mn is shown in Fig.~\\ref{fig:decay}; the dashed lines indicate the decay paths of the stripped $^{54}$Mn.  In two recent difficult and elegant experiments the very small branching ratio for the $\\beta^+$ decay to the ground state of $^{54}$Cr has been measured: $(2.2 \\pm 0.9) \\times 10^{-9}$~\\cite{b+1} and $(1.20 \\pm 0.26) \\times 10^{-9}$~\\cite{b+2}. By taking the weighed mean of these values we extract the averaged branching ratio of $(1.28 \\pm 0.25) \\times 10^{-9}$.  Combining it with the known halflife for $^{54}$Mn of 312.3(4)~d~\\cite{nds54}, it corresponds to an experimental partial $\\beta^+$ halflife of $(6.7 \\pm 1.3) \\times 10^8$ years.  As explained in~\\cite{Ulysses1,b+1,b+2} one expects, however, that the decay of the fully stripped $^{54}$Mn will be dominated by the at present unobservable $\\beta^-$ decay to the $^{54}$Fe ground state.  Previously, the partial $\\beta^-$ halflife was estimated assuming that the $\\beta^-$ and $\\beta^+$ form factors are identical. Very recently, in Ref. \\cite{b+2}, the ratio of the $\\beta^-$ and $\\beta^+$ form factors was calculated using a very truncated shell-model and extending it by comparison with similar calculations in the $sd$-shell. The estimated $\\beta^-$ halflife is $(6.3 \\pm 1.3) \\times 10^5$~y \\cite{b+2}.  In this work we will use the state of the art shell model and evaluate not only the $EC$ decay rate of the normal $^{54}$Mn, but also both decay branches of the unique second forbidden transitions $^{54}$Mn$(3^+) \\rightarrow {}^{54}$Cr$(0^+)$ and $^{54}$Mn$(3^+) \\rightarrow {}^{54}$Fe$(0^+)$.  By comparing the calculated $\\beta^+$ decay halflife (or branching ratio) to the measured one we hope to judge the reliability of the calculation. We then proceed to calculate the halflife of the unknown $\\beta^-$ decay.  The decays of stripped $^{54}$Mn are unique second forbidden transitions which depend on a single nuclear form factor (matrix element).  Halflives of several such decays in the $p$ shell ($^{10}$Be) and $sd$ shell ($^{22}$Na, and two decay branches of $^{26}$Al) are known and have been compared to the nuclear shell model predictions in Ref.~\\cite{War}. For the $sd$ shell nuclei, however, only calculations in a severely truncated space were performed in~\\cite{War}.  Since that time computation techniques and programming skills have improved considerably.  Thus, in order to further test our ability to describe this kind of weak decays, we repeat the analysis~\\cite{War}, using the exact shell model calculations without truncation. At the same time the availability of new (and different) experimental data for the $^{10}$Be~\\cite{be10exp} and $^{26}$Al~\\cite{al26exp} decays make necessary a new comparison between experiment and calculations.  In order to evaluate the decay rate we use the formulation of~\\cite{tables}.  The number of particles with momentum $p$ emitted per unit time is: \\begin{equation} N(p_e) dp_e = \\frac{g^2}{2 \\pi^3} p_e^2 p_{\\nu}^2 F(Z,W_e) C(W_e) dp_e, \\end{equation} where $g$ is the weak coupling constant, $p_e$ and $W_e$ are electron (or positron) momentum and energy, $p_{\\nu}$ is the neutrino (or antineutrino) momentum, and $Z$ is the atomic number of the daughter nucleus. All momenta and energies are in units where the electron mass is unity.  For the Fermi function $F(Z,W_e)$ we use the tabulated values, and the shape factor $C(W_e)$ for the case of the unique second forbidden transitions is of the form \\begin{equation} C(W_e) = \\frac{R^4}{15^2} \\left| ^AF_{321}^{(0)} \\right|^2 \\left[ p_\\nu^4 + \\frac{10}{3} \\lambda_2 p_\\nu^2 p_e^2 + \\lambda_3 p_e^4 \\right]. \\end{equation} The nuclear form factor, in turn, is defined as \\begin{equation} ^AF_{321}^{(0)} = g_A \\sqrt{\\frac{4\\pi}{2J_i + 1}} \\frac{\\langle f || r^2 [\\bbox{Y}_2 \\times \\bbox{\\sigma}]^{[3]}\\bbox{t}_\\pm || i \\rangle}{R^2}, \\label{eq:formfactor} \\end{equation} where $i$ denotes the initial state and $f$ the final one; the matrix element is reduced with respect to the spin space only (Racah convention~\\cite{edmons}); $\\pm$ refers to $\\beta^\\pm$ decay; $\\bbox{t}_\\pm = (\\bbox{\\tau}_x \\pm i \\bbox{\\tau}_y)/2$, with $\\bbox{t}_+ p = n$; $g_A = -1.2599 \\pm 0.0025$~\\cite{towhar}; and $R$ is the nuclear radius (the final expression for $C(W_e)$ is obviously independent of $R$).  The functions $\\lambda_2$ and $\\lambda_3$ are tabulated in~\\cite{tables}.  Integrating the rate formula up to the spectrum endpoint we obtain the expression for $1/\\tau$ and, respectively, for the halflife ($T_{1/2} = \\ln(2) \\tau$) in terms of the nuclear form factor squared. (For the stripped atoms we correct the endpoint energy accordingly.). For the quantity $2 \\pi^3(\\ln 2)/g^2$ we use the value $6146\\pm 6$~s~\\cite{towhar}.  Note the usual $f t$ value, commonly used to characterize a decay, uses the integrated phase space factor $f$ of Eq. (1), however without the constant $g^2/(2 \\pi^3 15^2)$) and the radius factor $R^4$. ",
        "conclusions": "Our shell model calculations reproduce the experimental halflives of the unique second forbidden beta decays within a factor of less than two. No clear evidence for the quenching of the corresponding form factors emerges. For the stripped $^{54}$Mn decays, the shell model describes the $\\beta^+$ branch within errors. It predicts that the form factor for the $\\beta^-$ decay is larger than the one for the $\\beta^+$ decay. The calculated $\\beta^-$ halflife (and therefore also the total halflife) is noticeably shorter than the range based on the observation of $^{54}$Mn in cosmic rays. This conflict, albeit relatively mild, makes attempts to determine the branching ratio for the $\\beta^-$ decay experimentally even more compelling."
    },
    "9803/astro-ph9803285_arXiv.txt": {
        "abstract": "Analysis of several recent ROSAT HRI observations of the gravitationally lensed system Q0957+561 has led to the detection at the 3$\\sigma$ level of the cluster lens containing the primary galaxy G1. The total mass was estimated by applying the equation of hydrostatic equilibrium to the detected hot intracluster gas for a range of cluster core radii, cluster sizes and for different values of the Hubble constant.  X-ray estimates of the lensing cluster mass provide a means to determine the cluster contribution to the deflection of rays originating from the quasar Q0957+561. The present mass estimates were used to evaluate the convergence parameter $\\kappa$, the ratio of the local surface mass density of the cluster to the critical surface mass density for lensing. The convergence parameter, $\\kappa$, calculated in the vicinity of the lensed images, was found to range between 0.07 and 0.21, depending on the assumed cluster core radius and cluster extent. This range of uncertainty in $\\kappa$ does not include possible systematic errors arising from the estimation of the cluster temperature through the use of the cluster luminosity-temperature relation and the assumption of spherical symmetry of the cluster gas. Applying this range of values of $\\kappa$ to the lensing model of Grogin \\& Narayan (1996) for Q0957+561 but not accounting for uncertainties in that model yields a range of values for the Hubble constant: $ 67 <$ H$_{0}$ $<$ 82 $\\,$ km $\\,$ s$^{-1}$ $\\,$ Mpc$^{-1}$, for a time delay of 1.1 years. ",
        "introduction": "Refsdal (1964a 1964b 1966) outlined a method for determining a global value for the Hubble constant by measuring the light travel delay between different images of a gravitationally lensed system. Together with a model that describes the gravitational potential of the lens, the time delay defines the geometrical dimensions of the lens system and therefore may lead to the determination of the Hubble constant.  The criteria that make a gravitational lens system (GLS) suitable for measuring the Hubble constant can be summarized as follows: (a) The sky positions and redshifts of the lens and of the images of the lensed object should be measurable accurately; (b) The mass distribution of the lens should be subject to accurate and reliable estimation. When the lens has multiple components, as for Q0957+561, this estimation becomes especially difficult and benefits from many different types of observations and also from many and extended images, such as rings. For example, if the types of observations include radio and optical emission lines, which are less sensitive to microlensing, flux ratio measurements are more useful in constraining the lens model; and (c) The differences in propagation time from the lensed object to us through (at least two) different images (``time delay'') should be measurable with high accuracy. The source, i.e., the lensed object, must therefore be variable on time scales far shorter than the time delays, to allow accurate estimates to be made of the latter.  Several potential candidates for determining $H_{0}$ are the quasar Q0957+561 with two images and a measured time delay of 1.1 years (Pelt et al. 1996; Haarsma et al. 1997; Kundi\\'{c} et al. 1997), the quadruple image system PG 1115+080 with a measured time delay between components B and C of about 25 days (Schechter et al. 1997; Bar-Kana 1997), and B0218+357 with a time delay of 12 $\\pm$ 3 days (Geiger \\& Schneider 1996; Corbett et al. 1996). Combining measured time delays with a detailed lensing model, Refsdal's suggestion has been applied to Q0957+561 (Falco, Gorenstein, \\& Shapiro 1991; Grogin \\& Narayan 1996, hearafter GN) and PG1115+080 (Keeton \\& Kochanek 1997).   In this paper we focus on a study of the properties of Q0957+561 (Walsh, Carswell \\& Weyman 1979), one of the most extensively studied lensing systems which has been monitored since its discovery in both the optical and radio (see, for example, Schild \\& Thomson 1995; Vanderriest et al. 1989; Haarsma et al. 1997). The primary lensing components are an elliptical cD galaxy at a redshift of 0.355, usually referred to as the G1 galaxy, a cluster of galaxies containing the G1 galaxy, and a group of galaxies at a redshift of 0.5.  The most thorough optical spectroscopic and photometric study to date of the field around Q0957+561 has been presented by Angonin-Willaime, Soucail, \\& Vanderriest (1994).   One of the major concerns in modeling Q0957+561 is a mass sheet degeneracy which arises from the presence of the cluster of  galaxies at z=0.36. In particular if one were to modify the assumed radial mass density profile $\\kappa(\\theta)$ of the lens to ${\\lambda}{\\kappa(\\theta)}$ + (1 - $\\lambda$), where 1 - $\\lambda$ represents  a constant surface mass density sheet term, then the new mass distribution would also satisfy the observational constraints. However the resulting value for $H_{0}$ would be scaled by the constant 1 - $\\lambda$. Since $\\lambda$ must always be positive one may always determine an upper limit on the Hubble constant for $\\lambda$ = 0 (Falco et al. 1991).  As also noted by Falco et al. (1991) and later GN, this degeneracy may be broken if one were to measure the velocity dispersion or total mass distribution of either the cluster of galaxies or the principal lensing galaxy G1. Recently Falco et al. (1997) have obtained a velocity dispersion for the central lens galaxy G1 of $\\sigma_{v} = 266 \\pm 12$ km s$^{-1}$ or $\\sigma_{v} = 279 \\pm 12$ km s$^{-1}$ depending on the interpretation of the spectroscopic slit data.\\\\  In this paper we focus on improving the cluster model and we present an estimate of the total mass of this lensing cluster of galaxies through X-ray measurements of thermal emission from the intracluster gas. In section 2 we describe the spatial analysis of the X-ray data of Q0957+561 obtained from several ROSAT HRI observations and we present an estimate for the convergence parameter $\\kappa(\\theta)$ of the cluster. Section 3 describes the spectroscopic analysis of ROSAT position sensitive proportional counter (PSPC) and ASCA Gas Imaging Spectrometer (GIS) X-ray observations at different epochs in order to investigate the variability of different spectral components, and presents various scenarios to explain the observed variability in the X-ray flux. Finally section 4 presents our conclusions from the analysis of the X-ray observations of the  Q0957+561 GL system. In particular, we discuss how our results allow bounds to be placed on the Hubble constant, albeit with these bounds still subject to systematic errors that we cannot yet evaluate accurately. ",
        "conclusions": "The need for a detailed measurement of the mass distribution of the cluster lens has been emphasized by Falco et al.(1991) and by GN (1996) in analyses of lens models for Q0957+561 that incorporate the cluster contribution through a convergence parameter $\\kappa(r)$. The most important result of the present analysis of the ROSAT HRI data of Q0957+561, that we hope will eventually lead to an accurate determination of the mass distribution of the cluster lens, is the 3$\\sigma$ detection of X-rays from the cluster of galaxies that contains the principal lens galaxy G1. \\\\  As indicated in Figure 3, the uncertainty in our estimated values for the mass distribution of the cluster spans a large range as do the consequent estimates of the convergence parameter of the lens (Figure 6). The main contributors to this uncertainty are the lack of knowledge of the shape of the cluster, the core radius of the cluster, the temperature profile of the hot gas, the cluster extent, and the location of the cluster center as well as the sensitivity to the value of $H{_0}$ used for deriving the cluster hot gas temperature from the X-ray luminosity. The present analysis yields a total cluster mass in the range 1.5 - 3.2 $\\times$ 10$^{14}$M${\\odot}$, and does not include allowances for errors in several of the above characteristics, but does consider the effects of cluster core radii between 5$''$ and 45$''$, and cluster extents between 120$''$ and 280$''$. The calculations of the X-ray luminosity were performed for Hubble constants of 50 and 75 km s$^{-1}$ Mpc$^{-1}$. The average convergence parameter $\\overline{\\kappa} = ({\\kappa}_{A} + {\\kappa_B})/2$ of the cluster at the image locations A and B was found to lie between 0.07 and 0.21. \\\\  One of the other assumptions in the present calculation is that the temperature of the cluster follows the luminosity temperature relation of David et al. (1993) for a large sample of clusters of galaxies. The derived cluster temperature depends weakly on the cluster X-ray luminosity. An error of a factor of 2 in the X-ray luminosity derived from the HRI data will propagate as an error of about 30$\\%$ in the cluster temperature and about 7$\\%$ (for $\\kappa$ = 0.2) in the derived value of the Hubble constant. The residual rms scatter in the $L_{X}$ - T relation along the temperature axis is $\\sigma_{logT}$ = 0.104/2.1 according to Markevitch (1998) and propagates as an error of about 3$\\%$ in the derived value of the Hubble constant.   Our calculations for the deflection angles produced from considering the influence of the cluster alone (not including G1) found the deflection for images A and B to lie in the range of 2$''$ to 16$''$ as shown in Figure 6; the differences in these deflections for images A and B are much smaller. The resulting percent contribution, ${\\Delta}S_{AB}$, to the observed 6.$''$1 separation of images A and B due to the cluster alone, lies in the range of 4\\% to 22\\%. This range includes only the effects of varying the cluster core radius between 5$''$ and 45$''$ and the cluster limit between 120$''$ and 280$''$. This result is in contrast with the commonly accepted notion that the large separation of the two images of Q0957+561 is an indicator that the lens contains a cluster of galaxies. To test the sensitivity of this result to the adopted cluster center position, we computed the cluster contribution to the image separation for cluster center locations within a distance of 5$''$ of the location derived from the optical observations by Angonin-Willaime et al. (1994). As expected, when the cluster center is moved closer to the locations of images A and B, the derived cluster contribution to the image separation becomes more significant reaching a value as high as ${\\Delta}S_{AB}$ = 31\\% for a cluster center at a distance of about 16$''$ from image B. \\\\     Recently the mass distribution of the lens of Q0957+561 was determined by Fischer et al. (1997) based on the distortion of background galaxies produced by the lens. Their estimate of 3.7 $\\pm$ 1.2 $\\times$ 10$^{14}$M${\\odot}$ within a radius of 1Mpc is consistent with the X-ray estimates. As pointed out by Schneider \\& Seitz (1995) however, the Kaiser \\& Squires (1993) mass reconstruction technique also is insensitive to constant density sheets of matter.   An interesting conclusion from the mass reconstruction analysis by Fischer et al. (1997) is the relatively low value of 5$''$ derived for the cluster core radius. The cluster core radius, r$_{cF}$, as defined in Fischer et al. (1997), however, differs from the cluster core radius, r$_{c}$, used in this paper and defined in equation 2. The relation between r$_{c}$ and r$_{cF}$ is given approximately by r$_{cF}$ = -0.987 + 0.685r$_{c}$ for 5$''$ $<$ r$_{c}$ $<$ 50$''$. We will have to await future X-ray observations of Q0957+561 to compare more accurately  X-ray and weak-lens results for the cluster core radius. \\\\  The present ROSAT detection of cluster emission suggests that future X-ray observations will greatly add to our understanding of Q0957+561. AXAF observations should provide direct determination of the temperature profile, the core radius and the shape of the cluster lens, reveal possible clumpiness in the total mass profile, and ultimately provide a tighter constraint on the Hubble constant.\\\\"
    },
    "9803/astro-ph9803146_arXiv.txt": {
        "abstract": "Many disk galaxies are lopsided: their brightest inner parts are displaced from the center of the outer isophotes, or the outer contours of the H{\\sc i} disk.  This asymmetry is particularly common in small, low-luminosity galaxies.  We argue here that long-lived lopsidedness is a consequence of the disk lying off-center in the potential of the galaxy's extended dark halo, and spinning in a sense retrograde to its orbit about the halo center.  The stellar velocity field predicted by our gravitational $N$-body simulations is clearly asymmetric. ",
        "introduction": "Many galaxies display lopsided, rather than bisymmetric, structure, on both large and small scales.  Hubble Space Telescope observations show that the nuclei of M31 and NGC 4486B do not lie at the center of the bulge isophotes (Lauer \\etal\\ 1993, 1996; Davidge \\etal\\ 1997).  In our own Milky Way, the distribution of molecular gas as measured in CO is offset about $1^{\\circ}$ in longitude and +15\\kms\\ in velocity from the center (Binney \\etal\\ 1991); other mass concentrations in that region also appear to be off-center (Blitz 1995).  On kiloparsec scales, many galaxies are asymmetric in both their optical appearance and their gas distribution (\\eg\\ Baldwin, Lynden-Bell \\& Sancisi 1980).  Working from a sample of 1700 spectra, Richter \\& Sancisi (1994) estimate that about half of all late-type spiral galaxies show clearly asymmetric H{\\sc i} profiles, indicating a lopsided gas disk; about half of this sample are nearby field galaxies, and none lie in rich clusters where interactions between galaxies are frequent.  Zaritsky \\& Rix (1997) estimate that about 50\\% of spiral galaxies in the field are significantly lopsided at optical and near-infrared wavelengths; the asymmetry involves both young and old stellar populations, and most of their lopsided galaxies lack any obvious companions.  Asymmetric bars are particularly common in late type or small disk galaxies, less luminous than $M_B \\sim 18$ (\\eg\\ Feitzinger 1980; Odewahn 1996; Matthews \\& Gallagher 1997). While some lopsidedness in the outer parts of galaxies is undoubtedly due to recent accretion, the fact that lopsided asymmetries are so common suggests that they can be long-lived.  The origin and persistence of these disk asymmetries remain a mystery.  Baldwin, \\etal\\ (1980) suggest that the lopsided distortions seen in the outer parts of galaxies are kinematic features; in that case they should wind themselves into a leading spiral, since the slow pattern speed $\\Omega - \\kappa$ is negative. In fact the one-armed spiral often has the same sense of winding as the two-armed components, which presumably trail (Colin \\& Athanassoula 1989; Phookun \\etal\\ 1992, 1993).  The only analytic disk models which are strongly unstable to lopsided distortions are those with many retrograde-streaming particles (Zang \\& Hohl 1978; Sawamura 1988; Sellwood \\& Merritt 1994); but observed counter-rotation is rare in disk galaxies (\\eg\\ Kuijken, Fisher \\& Merrifield 1996).  Adams, Ruden \\& Shu (1989) suggested that a protostellar disk could become unstable to a lopsided distortion; in the galactic context, this would correspond to the stellar disk being off-center in the dark halo. However, Heemskerk, Papaloizou \\& Savonije (1992) showed that the $m=1$ distortion is generally evanescent in the disk, so that resonant growth will not occur.  Matthias (1993), investigating orbits in a weakly `egg-shaped' tumbling potential, found that prograde orbits were distorted so as to oppose the `eggness'.  Miller \\& Smith (1992) describe an $N$-body simulation of an elliptical galaxy in which a central tilted disk developed of an off-center motion, tracing an oscillation in the underlying particle distribution, and Taga \\& Iye (1998) have done simulations in which a massive object wanders off-center in a dense stellar system; but no similar results have been reported in the literature. By contrast, kinematic models for lopsided systems have been quite successful; orbits in the potential of an off-center bar rotating rigidly in an axisymmetric disk can account for the rotation curves of Magellanic barred systems (de Vaucouleurs \\& Freeman 1973; Christiansen \\& Jeffreys 1976), while the gas response has a trailing one-armed form (Colin \\& Athanassoula 1989).  We propose that the key to the puzzling ubiquity of lopsided galactic disks is that they are not dynamically isolated inside the galaxy. Late-type spirals and dwarf galaxies, in which larger-scale lopsidedness is most frequently observed, are most likely to be dominated by dark halo mass: \\eg\\ Casertano \\& van Gorkom (1991); C\\^ot\\'e, Carignan \\& Sancisi (1991); Broeils (1992a, b).  (Athanassoula, Bosma \\& Papaoiannou 1987 suggested that in later-type galaxies the halo core is larger in relation to the radius of the stellar disk, although this was not confirmed by Broeils 1992a, b.)  If a disk found itself off-center in a dominant dark halo with a core of nearly constant density that was large compared to the disk dimensions, most of the disk mass would lie in a region where the angular speed of an orbit about the halo center did not vary strongly.  Self-gravity might then act to maintain the disk's coherence as it orbits the center of the halo potential well. This paper presents our investigation of just such a model using a self-consistent, self-gravitating disk embedded in a static dark halo. ",
        "conclusions": "Our experiments imply that a gravitating galactic disk can remain off-center in a halo potential as long as it orbits in a region where the halo has nearly constant density; the effect is more pronounced if the spin of the disk is in the opposite sense to its orbit around the halo center.  This behavior can be understood in the limit where the disk dimensions are small, and its self-gravity is weak: a particle in near-circular orbit in the halo then follows an epicycle in a sense retrograde to the orbit.  A collection of particles with the same angular momentum, and hence with a common guiding-center radius, would stay together as they orbited the halo center.  In the halo core, where the orbital frequency does not vary strongly with radius, the disk's self-gravity can fairly easily maintain coherence in approximately this configuration.  The relatively large fraction of lopsided systems found at redshifts $z \\sim 1$, which was apparently a period of rapid galaxy formation (\\eg\\ Driver, Windhorst \\& Griffiths 1995), suggests that many disks can slowly lose their off-center character as they settle down.  In dwarf galaxies, the halo is relatively more massive (Broeils 1992a, b), so lopsidedness is more easily maintained than in more luminous systems. The relation between a lopsided disk and the presence of a central bar has not been studied systematically, but both these features are characteristic of Magellanic irregular systems.  The lopsided Sc galaxy NGC~1637 (Block \\etal\\ 1994) appears prominently barred in an infrared image, while at optical wavelengths the bar is hidden by dust and confused by star formation.  The spatial scale of the predicted asymmetry is approximately that of the constant-density core; consistent with this idea, the off-center nuclei of M31 and the elliptical NGC~4486B lie within resolved central cores in those galaxies.  Our results are consistent with those of Bontekoe (1988), who investigated the sinking of a satellite galaxy into a larger system. In his unpublished fifth chapter, he reports that the satellite never sank all the way to the center, but the orbit shrank until it enclosed about 10\\% of the main galaxy's mass. The larger system was represented by a series of polytropes; the more concentrated the polytrope, the closer the satellite sank to its center.  Our study also has aspects in common with work on galaxy mergers by Miller \\& Smith (1995), who found that if two polytropic `cores' orbited each other in isolation, dynamical friction caused them to merge within a couple of orbits.  But in the presence of a `halo' of constant density, whether it was represented by a fixed potential, or consisted of `live' simulation particles, the cores survived for $15-30$ orbits without coming noticeably closer, although the interpretation of the `live halo' model was complicated by an overstable oscillation in the cores' orbital radii.  Persistent off-center disks have not been reported in published $N$-body simulations following a galactic disk in an imposed external bulge potential (see \\eg\\ Sellwood \\& Wilkinson 1993); but in these computations, the mass distribution of the bulge was in general more concentrated than the disk.  We must be concerned about the effect of dynamical friction on the disk, from the particles which would make up a `live' galactic halo. In the context of M31's nucleus, King, Stanford \\& Crane (1995) point out that this drag could well be small, if the bulge stars stream rapidly in the same direction as the orbiting cluster's motion. It is unclear even how far the usual estimates of dynamical friction are to be trusted. Bontekoe's (1988) satellite stopped sinking at some distance from the center of the system into which it was accreting. The physical situation of an off-center disk has much in common with that of a rotating bar, where estimates of the drag from the particles of a spherical halo suggest that the bar should rapidly spin down (Weinberg 1985), and gravitational $N$-body experiments also indicate strong braking. But observed galactic bars seem to be fast-rotating: Sellwood (1996) summarizes this confusing situation. The long-lived pairs of {\\it counter-rotating} bars which develop in the gravitational $N$-body simulations of Sellwood \\& Merritt (1994) and Friedli (1996) are clearly not discouraged by dynamical friction.  It may well be more useful, instead of calculating only the back-reaction of the halo on the orbiting disk, to look for long-lived modes of oscillation in the combined system.  Weinberg (1991, 1994) has developed an analytic method for doing this; examining spherical King models with an isotropic velocity dispersion, he found $m=1$ modes which, once excited, require many orbital times to decay.  The distortions lasted longest in the models in which the core was largest in relation to the total extent of the system.  He then used gravitational $N$-body simulations to confirm the existence of these modes.  Similar lopsided modes may well exist for a disk within a galactic halo.  We would then expect a relatively long-lived lopsided structure to result when material is accreted such as to push it off-center in the underlying potential. In the central regions of a galaxy, where the gravitational force is provided largely by a `hot' stellar system such as a bulge, an inner disk or nuclear star cluster could remain off-center over many orbital periods.  Gas might be captured by a galactic disk such as to push it off-center in the halo; alternatively, accretion of dark matter onto the halo may result in the disk lying off-center.  If the orbit of the disk is retrograde with respect to its spin, the lopsidedness is likely to be stronger and more persistent.  To maintain appreciable asymmetry, the near-constant-density halo core should be large enough to encompass a substantial fraction of the disk mass."
    },
    "9803/astro-ph9803236_arXiv.txt": {
        "abstract": "By redshift of $10$, star formation in the first objects should have produced considerable amounts of Carbon, Nitrogen and Oxygen. The submillimeter lines of C, N and O redshift into the millimeter and centimeter bands ($0.5\\,{\\rm mm}$--$1.2\\,{\\rm cm}$), where they may be detectable. High spectral resolution observations could potentially detect inhomogeneities in C, N and O emission, and see the first objects forming at high redshift. We calculate expected intensity fluctuations and discuss frequency and angular resolution required to detect them. For CII emission, we estimate the intensity using two independent methods: the line emission coefficient argument and the luminosity density argument. We find they are in good agreement. At $1+z \\sim 10$, the typical protogalaxy has a velocity dispersion of $30 \\,{\\rm km\\,s^{-1}}$ and angular size of $1 \\,{\\rm arcsecond}$. If CII is the dominant coolant, then we estimate a characteristic line strength of $\\sim 0.1 \\,{\\rm K\\, km\\, s}^{-1}$. We also discuss other atomic lines and estimate their signal. Observations with angular resolution of $10^{-3}$ can detect moderately nonlinear fluctuations of amplitude $2 \\cdot 10^{-5}$ times the microwave background. If the intensity fluctuations are detected, they will probe matter density inhomogeneity, chemical evolution and ionization history at high redshifts. ",
        "introduction": "When did the first stars and galaxies form? How can we detect them? Quasars were once the most distant objects that have ever been observed (\\cite{sch91}). Now galaxies are beginning to take their place. Discovery of a large number of galaxies at $z>3$ (\\cite{ste96}), and even up to $z\\sim 5$ (\\cite{fra97,dey98}), suggests that galaxies existed well before quasars did. In hierarchical models, star formation starts relatively early, by $z\\sim30$ (\\cite{cou86,fuk94,gne96,teg97}), so even these $z \\sim 5$ galaxies are not representative of the first generation of objects.  The purpose of this paper is to investigate the possibility of detecting primordial objects through the atomic lines of carbon, nitrogen and oxygen, which should have been produced as a result of star formation. At low redshift, these submillimeter lines are important coolants (e.~g., \\cite{ben94,isr95,stu97,you97,barv97,mal97}). If the source is at redshift $1+z \\sim 10$--$20$, the lines are redshifted to radio frequencies, where the atmosphere is more transparent and ground--based detectors with very high sensitivity and resolution exist.  If the source was uniformly distributed through space, then its emission would merely distort the spectrum. Since heavy elements produced by stars are distributed inhomogeneously, the observed emission should have intensity fluctuations as a function of both position and frequency. With high enough resolution one can in principle distinguish them from background. If we take $1 h^{-1}\\,{\\rm Mpc}$ as a characteristic comoving scale of the inhomogeneity, this corresponds to $\\sim 1$ arcmin in angular scale, and $\\sim 10^{-3}$ in frequency resolution. On smaller scales, velocity dispersion of matter limits the scale that can be resolved in frequency space. In this paper we formulate the intensity fluctuations and estimate their magnitude. This signal should give us information on number density fluctuations of specific source elements at $1+z>10$. Detailed determination of the peak of this signal will also tell us at which epoch elements were ionized.  The rest of this paper is organized as follows. In the next section, we present our formalism for estimating the fluctuating signal. Then we estimate its magnitude in the case that the matter density fluctuations are calculated from the linear power spectrum in section \\ref{sect:linear}. We also estimate optimal frequency and angular resolution for an object search. We focus on atomic line redshifted into present radio band, investigating carbon, nitrogen and oxygen emission lines at rest frame wavelength of hundreds of micrometers. In section \\ref{sect:nonlin}, we discuss the size and the number density of nonlinear clumps at the reionization epoch, and estimate intensity from nonlinear clumps and the interval between enhance of intensity fluctuations. Section \\ref{sect:paremis} gives discussion of merits and demerits as the signal of specific emission. After alternative estimation of CII luminosity in section \\ref{sect:altcii}, the final section gives our conclusion. ",
        "conclusions": "We have proposed a new and promising method to detect the first objects: search for carbon, nitrogen and oxygen line emission at $1+z=10$--$20$. We conclude that the intensity fluctuations originated from the emission are detectable in radio band ($0.5\\,{\\rm mm}$--$1.2\\,{\\rm cm}$). If observation is carried out with angular resolution $\\theta_{\\rm res} = 1 \\,{\\rm arcmin}$ and spectral resolution $\\Delta \\nu / \\nu = 10^{-3}$, the magnitude of the smoothed intensity fluctuations corresponding with linear matter density fluctuations is predicted as $\\Delta \\bar{I}_{\\rm line} / I_{\\rm tot}= (1$ and $3) \\times 10^{-6}$ (in CI $609$ and $370 \\,\\mu {\\rm m}$ cases) from sources at $1+z=10$, in the standard CDM model. The predictions for other lines are $\\Delta \\bar{I}_{\\rm line} / I_{\\rm tot} = (20, 70, 3, 20) \\times 10^{-6}$, respectively, in [$158$(CII), $146$(OI), $204$(NII), $122$(NII)] $\\,\\mu {\\rm m}$ cases.  Observation with $\\theta_{\\rm res} = 1 \\,{\\rm arcsec}$ and $\\Delta \\nu / \\nu = 10^{-4}$ resolve individual protogalaxies with velocity dispersion of $30\\,{\\rm km\\,s^{-1}}$. We predict line intensities relative to the microwave background of $\\Delta I_{\\rm line} / I_{\\rm tot} = (1, 4, 30, 80, 4, 30) \\times 10^{-4}$ or equivalently $\\Delta T_{\\rm line} v_{\\rm typ} = (0.8, 2, 7, 20, 1, 6) \\times 10^{-2} \\,{\\rm K\\, km\\, s}^{-1}$ for [$609$(CI), $370$(CI), $158$(CII), $146$(OI), $204$(NII), $122$(NII)] $\\,\\mu {\\rm m}$ lines, respectively. If detected, these features will provide useful information on matter inhomogeneities at $1+z=10$--$20$. They will also probe chemical evolution and reionization time $z_{\\rm re}$ itself."
    },
    "9803/astro-ph9803222_arXiv.txt": {
        "abstract": "We define a complete sample of thirty-three  GHz-Peaked-Spectrum (GPS) radio sources based on their spectral properties. We present measurements of the radio spectra and polarization of the complete sample and a list of additional GPS sources which fail  one or more criteria to be included in the complete sample.  The majority of the data have been obtained from quasi-simultaneous multi-frequency observations at the Very Large Array (VLA) during 3 observing sessions. Low frequency data from the Westerbork Synthesis Radio Telescope (WSRT) and from the literature have been combined with the VLA data in order to better define the spectral shape.  The objects presented here show a rather wide range of spectral indices at high and low frequencies, including a few cases where the spectral index below the turnover is close to the theoretical value of 2.5 typical of self-absorbed incoherent synchrotron emission. Faint and diffuse extended emission is found in about 10\\% of the sources.  In the majority of the GPS sources, the fractional polarization is found to be very low, consistent with the residual instrumental polarization of 0.3 $\\%$. \\footnote{Tables 4 and 5 are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/Abstract.html} ",
        "introduction": "The GHz-peaked-spectrum (GPS) radio sources are characterized by a simple convex spectrum which peaks in a range of about a decade around 1 GHz.  General discussions on the general properties of these objects are given by O'Dea \\etal\\  (1991) O'Dea and Baum (1997), and O'Dea (1998) where an exhaustive   bibliography can also be found.  Common characteristics of the bright sample of GPS radio sources are: small size ($\\lae 1$ kpc), high radio luminosity, low fractional polarization, and apparently low variability. They are a mixed class of quasars and galaxies. Galaxies tend to be $L_*$ or brighter and at redshifts $0.1 \\lae z \\lae 1$  (O'Dea \\etal\\ 1996) while quasars are often found at  very large redshift $1 \\lae z \\lae 4$, (O'Dea 1990).  Currently, there are two main competing hypotheses to explain the origin of GPS radio sources and their possible evolution.  In the ``young source\" scenario, first suggested by Phillips and Mutel (1982), GPS radio sources with compact double (CD) morphology  or  compact symmetric morphology (Compact Symmetric Objects: CSO) are classical double radio sources at the very first stage of their lives.  In the ``frustration'' scenario, GPS radio sources will never become as large as the classical doubles since they are confined to the sub-kpc scale by a dense and turbulent ambient medium (e.g., O'Dea \\etal\\ 1991, Carvalho 1994, 1998). However, it is unclear whether the gas density and its distribution in the nuclear region is sufficient to confine the radio source for times of the order of 10$^7$ years (see e.g., O'Dea 1998).  Recent results support the young source  hypothesis. Fanti \\etal\\  (1995)  presented a model for the evolution of  the galactic-scale Compact Steep Spectrum (CSS) sources into large scale classical doubles.  They argued that the objects showing symmetric morphology are probably not confined by a dense and clumpy medium (see also De Young 1993). They also suggest that the typical CSS source age is of the order of 10$^6$ years and that the radio source luminosity decreases by an order of magnitude as the size of the radio source grows from a few kpc to hundreds of kpc (see also Begelman 1996). The same conclusions are reached by Readhead et al. (1996a,b) in their study of a few CSO's.  O'Dea and Baum (1997), combining the complete sample of GPS presented here with the CSS sources by Fanti \\etal\\ (1990) reach similar conclusions. They further note note that the luminosity evolution required makes it likely that  GPS and CSS  radio sources will evolve into objects less powerful than the most powerful classical doubles.  The detection of arc-second scale faint extended emission around $\\sim 10\\%$ of GPS sources  (Baum \\etal\\ 1990, Stanghellini \\etal\\ 1990) motivated Baum \\etal\\ (1990) to suggest that nuclear activity is recurrent in these sources. In this hypothesis, we see the relic of a previous epoch of activity as faint diffuse emission surrounding the current young nuclear source.  In order to achieve a deeper understanding of the GPS radio sources we pursued the following project.  (1) We created  a complete sample of GPS radio sources by means of a preliminary bibliographic research, checked with subsequent new multi-frequency data from the Very Large Array (VLA) and the Westerbork Synthesis Radio Telescope (WSRT).  (2) We determined the properties of the radio spectrum and the polarization.  (3) We obtained optical imaging to determine the host galaxy properties.  (4) We obtained  VLBI observations to study the  milliarcsecond radio morphology.  In this paper we present the selected sample and the results from the VLA and WSRT radio observations. The optical properties of GPS radio sources are discussed by O'Dea, Baum and Morris (1990), Stanghellini \\etal\\ (1993), and O'Dea \\etal\\  (1996). The milliarcsecond morphology is presented in Stanghellini \\etal\\  (1997). Constraints on radio source evolution based on our observations of the complete GPS sample are discussed by O'Dea \\& Baum (1997).  H$_0$=100 km sec$^{-1}$ Mpc$^{-1}$, and q$_0$=0.5 have been used in this paper. ",
        "conclusions": "In this paper we present the observational results and the primary qualitative conclusions. In a future paper we will present a quantitative  discussion of the results.   \\subsection{The radio spectra}  The flux densities are presented in Tables 4 and 5 and plotted in Figures 1 to 5.     In order to extract information from the radio spectra, we fitted them with a hyperbola (which is a curve tending asymptotically to a straight line at the extrema). We first fit the spectral indices in the thick and thin part of the spectrum independently, and then fixed  these two values (which correspond to the angular coefficients of the asymptotes) and we performed a least square fit, solving for the other parameters of the hyperbola.  We calculated the frequency of the spectral peak (Tables 1 and 2) using  the fitted curves.  We show the distributions of the observed and rest frame turnover frequency for the 33 sources of the complete sample in Figures 6c and 6d. In Fig. 6a and 6b we show  the high frequency (above the peak) and low frequency (below the peak) spectral index distribution, respectively.  \\subsection{The spectral indices and the turnover frequency}  The GPS radio sources are a mixed group of galaxies and quasars, with some remarkable differences between the two classes. The histogram in Fig. 16a shows that the redshift distribution is very different for galaxies and quasars. The galaxies have a typical redshift of $\\sim$ 0.5, and none has a redshift higher than 1.  For the galaxies without redshift information, we note that only 0316+161 has an optical magnitude slightly fainter than the galaxy 2128+048 at redshift 0.99 (Table. 1).  Since these galaxies follow the Hubble diagram (O'Dea \\etal\\ 1996; Snellen \\etal\\ 1996) it is unlikely they will be found at a redshift much higher than the others in the sample. The quasars  are instead found at any redshift (we included the galaxy 1404+286 (OQ208), which has a Seyfert 1 nucleus, in the quasar class) with the majority at very high z (see also O'Dea 1990).  \\begin{figure} \\vbox{\\psfig{figure=figure7.ps,width=8cm,height=9cm}}\\par \\caption[]{$0248+430$ at 1.35 GHz. The restoring beam is 1.41 $\\times$ 1.32 arcsec in P.A. $-78^\\circ$. The r.m.s. noise on the image is 0.2 mJy. The peak flux is 810 mJy/beam. The contour levels for all the images are -3, 3, 6, 12, 25, 50, 100, 200, 500, 1000 $\\times$ the r.m.s noise} \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure8.ps,width=8cm,height=9cm}}\\par \\caption[]{$0528+134$ at 1.66 GHz. The restoring beam is 1.14 $\\times$ 1.01 arcsec in P.A. $+2^\\circ$. The r.m.s. noise on the image is 0.25 mJy. The peak flux is 2088 mJy/beam. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure9.ps,width=8cm,height=9cm}}\\par \\caption[]{$0738+313$ at 1.33 GHz. The restoring beam is 2 $\\times$ 2 arcsec. The r.m.s. noise on the image is 0.3 mJy. The peak flux is 1943 mJy. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure10.ps,width=8cm,height=6cm}}\\par \\vskip -1cm \\caption[]{$0941-080$ at 1.33 GHz. The restoring beam is 2.44 $\\times$ 1.54 arcsec in P.A. $+29^\\circ$. The r.m.s. noise on the image is 0.5 mJy. The peak flux is 2088 mJy/beam. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure11.ps,width=8cm,height=9cm}}\\par \\caption[]{$2134+004$ at 1.33 GHz. The restoring beam is 1.91 $\\times$ 1.55 arcsec in P.A. $-5^\\circ$. The r.m.s. noise on the image is 0.4 mJy. The peak flux is 3249 mJy/beam. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure12.ps,width=8cm,height=9cm}}\\par \\caption[]{$2223+210$ at 1.33 GHz. The restoring beam is 1.37 $\\times$ 1.32 arcsec in P.A. $-7^\\circ$. The r.m.s. noise on the image is 0.3 mJy. The peak flux is 1663 mJy/beam. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure13.ps,width=8cm,height=9cm}}\\par \\caption[]{$2223+210$ at 5 GHz. The restoring beam is 0.37 $\\times$ 0.35 arcsec in P.A. $+5^\\circ$. The r.m.s. noise on the image is 0.3 mJy. The peak flux is 1034 mJy/beam. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure14.ps,width=8cm,height=9cm}}\\par \\caption[]{$2230+114$ at 1.33 GHz. The restoring beam is 1.46 $\\times$ 1.36 arcsec in P.A. $+1^\\circ$. The r.m.s. noise on the image is 0.8 mJy. The peak flux is 6444 mJy/beam. } \\end{figure}  \\begin{figure} \\vbox{\\psfig{figure=figure15.ps,width=8cm,height=9cm}}\\par \\caption[]{$2230+114$ at 5 GHz. The restoring beam is 0.38 $\\times$ 0.36 arcsec in P.A. $-11^\\circ$. The r.m.s. noise on the image is 0.5 mJy. The peak flux is 4137 mJy/beam. } \\end{figure}  \\begin{figure*} \\vbox{\\psfig{figure=figure16.ps,width=16.5cm,height=22cm}}\\par \\vskip -6cm \\caption[]{Histograms for the complete sample: a) redshift distribution; b) fractional polarization at 1.3 GHz; c) fractional polarization at 4.9 GHz; d) fractional polarization at 8.5 GHz } \\end{figure*}    The high frequency spectral index ranges from 0.5 (the limit set in the selection criteria) to 1.3 with the galaxies having perhaps slightly steeper values (but the 2 objects with the steepest spectral indices are quasars). The low frequency spectral index ranges from -0.2 to -2.1 without any clear difference between galaxies and quasars, but the higher values are biased since in several objects the low frequency part of the spectrum is under-sampled and the spectral index is calculated close to the turnover frequency where it is likely to be flatter. Similar results were found by De Vries \\etal\\ (1997) though their poorer frequency coverage resulted in a somewhat smaller range in spectral index.  The quasars tend to peak at higher frequencies than the galaxies in both the rest frame and observed frame and  some quasars have a turnover frequency in their rest frame exceeding 10 GHz. This suggests that, on the assumption that the turnover is caused by synchrotron self-absorption, the GPS quasars are more compact than the galaxies. This effect has been also found by De Vries \\etal\\ (1997) in a bright heterogeneous sample and by Snellen (1997) in a fainter sample. In addition, VLBI images of several  sources belonging to the complete sample show that quasars are more compact than galaxies, and in general exhibit different morphologies (Stanghellini \\etal\\ 1997).  These results suggest that either GPS galaxies and GPS quasars are different types of objects, or that beaming of compact components plays a role in the quasars (see also O'Dea 1998).    \\begin{table*} \\caption{Polarization for the complete sample } \\begin{flushleft} \\begin{tabular}{rrrrrrrrrrrrrrr} \\hline \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} &   1380  &&      1640   &&     4835   &&     4885&& && &&&\\\\ name   &  $\\%$  &   PA  &  $\\%$   &  PA&    $\\%$  &   PA &   $\\%$  &   PA &   &   &   &&& \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} 0019$-$000&   0.5&  +0 &  0.3   &-2 &  $<$0.1&    &     $<$0.1& & & & &&&\\\\ 0237$-$233&   2.0&  +7 &  1.2   &+4 &  4.0&  -35  & 4.0    &-35 & & & &&&\\\\ 0500+019&  $<$0.2&     & $<$0.2   &   &  $<$0.1&    &     $<$0.1& & & & &&&\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} &   1380&&        1630&&        4815&&        4865 &&       8435&&         8485&&&\\\\ name   &  $\\%$  &   PA  &  $\\%$   &  PA&    $\\%$  &   PA &   $\\%$  &   PA &   $\\%$ &    PA &    $\\%$  &  PA&&\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0108+388& $<$0.2&      &  0.5  &     & $<$0.2&&   $<$0.1&&       $<$0.1&       &  $<$0.1&     & &\\\\ 0316+161&  0.5&  -63 &  0.7  & -60 & $<$0.1&&     $<$0.1&&       $<$0.1&       &  $<$0.1&    &&\\\\ 0428+205&  0.4&  -59 &  0.6  & -58 & $<$0.1&&     $<$0.1&&       $<$0.1&       &  $<$0.1&   &&\\\\ 0710+439& $<$0.1&      & $<$0.3  &     & $<$0.1&& $<$0.1&&       $<$0.1&       &  $<$0.1& &&\\\\ 1031+567& $<$0.2&      & $<$0.2  &     & $<$0.3&& $<$0.2&&       $<$0.2&       &  $<$0.2&  && \\\\ 1323+321&  0.3&  -35 &  0.4  & -01 & $<$0.1&&     $<$0.1&&        0.9&   +14 &   0.8&  +15&&\\\\ 1404+286&  0.4&  -50 & $<$0.3  & -23 & $<$0.2&&   $<$0.1&&       $<$0.2&       &  $<$0.2&&&\\\\ 1518+047& $<$0.2&      & $<$0.2  &     & $<$0.3&& $<$0.1&&       $<$0.2&       &  $<$0.2& & & \\\\ 1607+268& $<$0.3&      & $<$0.2  &     & $<$0.1&& $<$0.1&&       $<$0.1&       &  $<$0.1& &&\\\\ 2008$-$068&  1.7&  -19 &  2.5  & -21 & $<$0.3&&     $<$0.3&&       $<$0.3&       &  $<$0.3& &  \\\\ 2352+495& $<$0.2&      & $<$0.3  & -34 & $<$0.1&& $<$0.1&&       $<$0.1&       &  $<$0.1& &&\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} &  1335&&        1665&&        4535 &&       4985&&        8085 &&        8465&&RM(rad/m$^2$)&\\\\ name    & $\\%$  &   PA &   $\\%$   &  PA  &  $\\%$ &    PA &   $\\%$  &   PA &   $\\%$ &    PA&     $\\%$ &   PA&obs&rest\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0248+430&  1.7&  +52? &  1.5  & +23  & 2.0&   -18 &  2.0 &  -27 &  0.5&   -40 &   0.4&  -44&131&703\\\\ 0457+024& $<$0.1&      & $<$0.1  &      &$<$0.3&       &  0.3 &  +47 &  1.0&   +14 &   1.2&  +12&258&2954\\\\ 0738+313& $<$0.1&      & $<$0.3  &      & 1.7&   +41 &  1.6 &  +67 &  3.0&   -2  &   3.2&  +2&-813&-2160\\\\ 0742+103& $<$0.2&      & $<$0.2  &      &$<$0.1&       & $<$0.2 &      & $<$0.1&       &  $<$0.1& & & \\\\ 0743$-$006&  0.4&  -76 &  0.5  & -7   &$<$0.2&       & $<$0.2 &      &  0.8&   +1  &   1.0&  -18& &\\\\ 0941$-$080& $<$0.2&      & $<$0.1  &      &$<$0.1&       & $<$0.1 &      & $<$0.2&       &  $<$0.2&    & &\\\\ 1117+146& $<$0.2&      & $<$0.1  &      &$<$0.2&       & $<$0.3 &      & $<$0.2&       &  $<$0.2&    && \\\\ 1127$-$145&  4.3&  -84 &  3.9  & +61  & 2.6&   -26 &  2.8 &  -24 &  3.3&   -17 &   3.3&  -18 &-49&-234\\\\ 1143$-$245&  1.9&  +72 &  2.1  & +80  & 0.7&   -18 &  0.7 &  -27 &  1.3&   -43 &  1.4 &  -46&146&1270\\\\ 1245$-$197& $<$0.2&      & $<$0.1  &      &$<$0.1&       & $<$0.1 &      & $<$0.1&       & $<$0.1 &  && \\\\ 1345+125& $<$0.2&      & $<$0.1  &      &$<$0.1&       & $<$0.2 &      & $<$0.1&       & $<$0.1 & & & \\\\ 1358+624& $<$0.2&      & $<$0.1  &      &$<$0.1&       & $<$0.3 &      & $<$0.2&       & $<$0.2 & &&\\\\ 1442+101&  0.8&  -74 &  0.8  & -78  & 2.0&   +60 &  1.7 &  +64 &  1.4&   +82 &  1.6 &  +78&-116&-2395\\\\ 1600+335& $<$0.2&      & $<$0.1  &      &$<$0.1&       & $<$0.1 &      & $<$0.1&       & $<$0.1 & &  & \\\\ 2126$-$158& $<$0.2&      & $<$0.2  &      &$<$0.1&       & $<$0.2 &      & $<$0.1&    & $<$0.2 & & &\\\\ 2128+048& $<$0.1&      & $<$0.1  &      &$<$0.1&       & $<$0.1 &      & $<$0.2&       & $<$0.2 & & &\\\\ 2134+004& $<$0.1&      & $<$0.1  &      & 0.9&   +60 &  0.9 &  +43 &  0.5&   -4  &  0.5 &  +0&349&3008\\\\ 2210+016& $<$0.1&      & $<$0.1  &      &$<$0.1&       & $<$0.1 &      & $<$0.2&       & $<$0.2 &   && \\\\ 2342+821& $<$0.2&      & $<$0.1  &      &$<$0.1&       &  0.5 &  -63 & &&&&&\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table*}   \\begin{table*} \\caption{Polarization of additional objects } \\begin{flushleft} \\begin{tabular}{rrrrrrrrrrrrrrr} \\hline \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} &  1335&&        1665&&        4535&&        4985&&        8085&&         8465&&RM(rad/m$^2$)&\\\\ name   &  $\\%$  &   PA  &  $\\%$  &   PA   & $\\%$     &PA   & $\\%$     &PA    &$\\%$     &PA     &$\\%$    &PA&obs&rest\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0404+768& $<$0.2&       &$<$0.1 &       &$<$0.1&     &   $<$0.3&      &       &     &       &    &&\\\\ 0440$-$003&  0.9&  +49  & 0.9 & -27   & 2.3&  +61&    2.1&   +58&    2.6&  +55&     2.6 & +56&28&\\\\ 0528+134&  0.6&  -56  &$<$0.3 &       & 2.2&  -31&    2.6&   -33&    3.9& -39 &    4.0  & -41&52&\\\\ 2149+056& $<$0.2&       &$<$0.2 &       &$<$0.1&     &   $<$0.1&      &  $<$0.2 &     &   $<$0.2  &   && \\\\ 2223+210& 13.3&  -28  &10.5 & -82   &10.0&  -64&    9.1&   -58&       &     &        &-123&-1077\\\\ 2230+114&  2.5&  -87  & 2.6 & -33   & 1.2&  +61&    1.5&   +75&       &     &       &-59&-244\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} &   1380& &       1630&  &      4815&&        4865 &&       8435&&         8485&&RM(rad/m$^2$)&\\\\ name   &  $\\%$  &   PA  &  $\\%$  &   PA   & $\\%$     &PA   & $\\%$     &PA    &$\\%$     &PA     &$\\%$    &PA&obs&rest\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0026+346&  0.4&  -45  & 0.6 & -43    &$<$0.2&     &   $<$0.2&     &   $<$0.1&     & $<$0.1&&&\\\\ 0201+113&  0.7&  -48  & 0.9 & -38    & 1.4&  +33&    1.4&  +33&    0.5&  +4 &  0.5&  +5&131&2724\\\\ 0552+398&  0.4&  -52  & 1.0 & +78    & 0.8&  +43&    0.6&  +53&    0.8&  +66&  0.8&  +66&-1344&-15218\\\\ 0703+468& $<$0.1&       &$<$0.2 &        &$<$0.2&     &   $<$0.2&     &   $<$0.3&     &  $<$0.3&&&\\\\ 0711+356&  1.3&  +5   & 1.0 & -33    & 1.8&  +83&    1.8&  +80&    1.9&  +76&  2.0& +79&40&275\\\\ 0904+039& $<$0.2&       &$<$0.2 &        &$<$0.5&     &   $<$0.5&     &   $<$0.9&     &  $<$0.9&&&\\\\ 0914+114& $<$0.2&       &$<$0.2 &        &$<$0.8&     &   $<$0.8&     &   $<$2.0&     &  $<$2.0&&&\\\\ 1543+005& $<$0.1&       &$<$0.2 &        &$<$0.2&     &   $<$0.2&     &   $<$0.2&     &  $<$0.1&&&\\\\ 1732+094& $<$0.2&       &$<$0.3 &        & 1.5&  +16&    1.4&  -1 &   $<$0.3&     &  $<$0.3&&&\\\\ 2015+657&  4.5&  -53  & 3.6 & -71    & 3.2&  +54&    3.1&  +55&    4.3&  +51&  4.3&  +51&30&\\\\ 2021+614&  0.6&  -53  & 1.0 & -51    &$<$0.2&     &   $<$0.2&     &   $<$0.1&     &  $<$0.1&&&\\\\ 2050+364& $<$0.2&       & 0.4 & -77    &$<$0.2&     &   $<$0.1&     &   $<$0.1&     &  $<$0.1&&&\\\\ 2137+209& $<$0.2&       &$<$0.1 &        &$<$0.2&     &   $<$0.2&     &    0.4&  +68&  $<$0.3&&&\\\\ 2337+264& $<$0.2&       &$<$0.3 &        &$<$0.3&     &   $<$0.2&     &   $<$0.2&     &  $<$0.2&&&\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table*}    \\subsection{Extended emission}   We have detected extended emission (both diffuse and compact) close to the compact radio source in some cases at 21 cm. In the remaining sources, our upper limits on extended emission is typically 1 mJy/beam at 21 cm.  0248+430 (Figure 7) has a compact emitting region 15 arcsec east of the main component and a hint of weak emission 5 arcsec to the south. 0528+134 is resolved, showing an extension in the NW direction (Figure 8). Murphy \\etal\\ (1993) present an image of 0738+313 at 20 cm showing 2 emitting regions resembling 2 weak hot-spots and lobes on the opposite sides of the dominant component. In our image (Figure 9) these 2 weak components are almost completely resolved out and only a hint of emission has been detected 30 arcsec north and south of the compact region. 0941-080 shows a slightly resolved secondary component 20 arcsec east of the main one (Figure 10). 2134+004 has very weak and diffuse emission around the strong compact component (Figure 11). 2223+210 has a secondary component 4 arcsec away from the main one in the SW direction in our image at 1.35 GHz (Figure 12); the main component itself is resolved in a core-jet structure oriented  NE with a possible counter jet in the image at 5 GHz (Figure 13). 2230+114 at 1.35 GHz (Figure  14) shows an elongated structure in the NW-SE direction with a hint of emission bending to SW, while in the 4.9 GHz image (Figure 15) the elongated structure turns out to be a core-jet structure  with  the possible presence of a counter jet.  Stanghellini \\etal\\ (1990) report several cases of extended emission around GPS radio sources. Of the objects presented here showing extended  emission, 0528+134 and 2223+210 are not true GPS objects (see also section 4.4 for a discussion of the case of 0528+134). In a couple of sources (0248+430, 0941-080 both belonging to the complete sample) it is difficult to say whether the secondary emission is related to the GPS radio source and further observations are probably needed. The extended emission found around 0738+313, 2134+004, and 2230+114 (the first 2 objects belong to the complete sample) is likely to be related to the GPS object. In the complete sample, 0108+388 is known to have extended emission, so there are 3 to 5 objects out of 33  with known extended emission so far. This percentage of 9 to 15$\\%$ is slightly smaller but consistent with that previously claimed by Stanghellini \\etal\\ (1990). It is clear that the vast majority ($\\sim 90\\%$) of the GPS sources appear to be truly isolated and have no emission beyond the kpc scale at the current limits.  \\subsection{Variability}   Waltman \\etal\\ (1991) presented monitoring observations at 2.7 and 8.1 GHz for several GPS sources covering the time range from 1979 to 1988. Some sources as 0237-233, 1245-197, 1345+125 were found to be very stable in flux density. Others were found to be variable: 0552+398 shows a variation of $\\sim$ 30 $\\%$ at 8.1 GHz. 2134+004 has a variability of about 15-20$\\%$ at 8.1 GHz. 2352+495 has a variability below 10$\\%$ at 8.1 GHz. Also 0742+103 is slightly variable.  Wehrle \\etal\\ (1992) also report variability for some GPS objects in the time range 1985-1991 at 4.8, 8, and 14.5 GHz, from the University of Michigan Radio Astronomy Observatory monitoring program for several sources, some of which are GPS objects. 0552+398 shows an increase in flux density exceeding 50$\\%$ at 8.4 and 14.5 GHz, 1127-145 shows a quasi periodical variability of approximately 1 Jy at all the 3 frequencies. The source 2230+114 also shows a rather remarkable flux density variability at all the 3 frequencies with an amplitude of 0.5-1 Jy and is a well known low frequency variable source (Bondi \\etal\\ 1996).  The variability of 1404+286 has been discussed by Stanghellini \\etal\\ (1996).  In conclusion,  we find that some  GPS sources (mainly quasars) show mild to high flux density variability at cm and mm wavelengths. However, without uniform monitoring of the complete sample it will not be possible to determine how common this variability is. We also note a couple sources where the spectral shape is variable and at some times the spectrum was peaked, and at other times it was not, 0528+134, the well known gamma-ray source (Mukherjee \\etal\\ 1996), and 0201+110. Both of these sources show a rather flat spectrum in the VLA observations from the second or the third session and they would be very easily discarded as GPS radio sources. But 0528+134 has been included in the class of GPS radio sources because of its GPS-like spectrum from the literature (O'Dea \\etal\\ 1991), and 0201+113 really shows a convex spectrum in the  data in the first VLA session published in O'Dea \\etal\\  (1990). This behavior is not surprising in highly variable radio sources as we may well expect that the presence of new radio components will change the spectral shape.  Thus, there are sources which show a peaked spectrum only part of the time. This aspect of the GPS phenomenon  deserves more attention as it could be related to the remarkably different properties found between GPS galaxies and (some?) GPS quasars.   \\subsection{Polarization}  Due to the low level of observed polarization, the errors are dominated by the r.m.s noise on the polarization images (typically 0.5 mJy) and by contamination from residual unpolarized emission (estimated about 0.2-0.3$\\%$ of the total flux density). The error in the polarized flux may then be calculated as $\\sigma _P=\\sqrt {0.5^2 + (0.003\\times S_{mJy})^2}mJy$. We considered any measurement of the fractional polarization below 0.3$\\%$ to be an upper limit (Tables 6 and 7). The errors in the position angle have a systematic contribution due to the uncertainty in the determination of the position angle of the calibrator (3C286), assumed to be around 2-3 degrees. This is the dominant contribution in the position angle error for most of the objects with detected polarized flux density.  In Fig. 16 we show the histograms of the fractional of polarization (mostly upper limits) for the complete sample at 1.3, 4.9 and 8.5 GHz. The fractional polarization is in general low at all the frequencies. Only a few quasars have a fractional polarization above 1 $\\%$ at 4.9 or 8.5 GHz.  When polarized emission has been detected we attempted a linear fit to the polarization angles versus the squared observed wavelength, as is expected from the Faraday effect on polarized radiation propagating through a magnetized and ionized medium.  The low level or even the lack of detection of polarized flux limited us to only a few sources  (all quasars). The fits are generally rather good and are given in Table 6 and 7, and in Figures 17 and 18. Due to the better frequency coverage in the present observations, our estimated rotation measures supersede those reported by O'Dea \\etal\\ (1990), though we cannot rule out that some of the difference is due to variability. We find Faraday rotation measures in the rest frame above 1000 rad/m$^2$ for 5 quasars of the complete sample. We also found a very high value ($>10^4$ rad/m$^2$) for 0552+398 which does not belong to the complete sample but has a GPS shape.  Sometimes the frequencies which give a good fit include those close to the turnover, and in the case of 0552+398 are all below the turnover. This implies that the region emitting the polarized emission is different from that responsible for the optically thick emission or that the turnover is not caused by synchrotron self absorption.    \\subsection{Summary and conclusions}  We have presented a bright flux-density-limited complete sample of 33 GPS radio sources selected on the basis of  their peaked radio spectra. The sample selection was based on observations with the VLA, WSRT, and other instruments.  Additional GPS sources not belonging to the complete sample have also been observed. We present our results on the polarization and radio spectrum. We found remarkable differences in the properties of  quasars and galaxies, the latter having lower turnover frequencies,  mostly undetectable polarization and lower redshifts.  In the few objects where polarization has been detected at many frequencies, the Faraday rotation measure in the rest frame often exceeds 1000 rad/m$^2$.  In about 10\\% of the sources we detect weak diffuse extended emission. In the remaining $\\sim 90\\%$ any extended emission has a peak surface brightness is less than about 1 mJy/beam at 21 cm.  In a following paper we will discuss the implications of the properties of the complete sample in the framework of the scenarios proposed to explain the existence of the GPS radio sources."
    },
    "9803/astro-ph9803152_arXiv.txt": {
        "abstract": "s{I discuss the interplay between inflation and microwave background anisotropies, stressing in particular the accuracy with which inflation predictions need to be made, and the importance of inflation as an underlying paradigm for cosmological parameter estimation.} ",
        "introduction": "Because of calculational simplicity, and because it provides a good fit to the observational data, an initial spectrum of adiabatic density perturbations is normally assumed responsible for all the observed structures in the Universe, such as galaxy clusters and microwave background anisotropies. The inflationary cosmology provides a natural explanation for such an initial spectrum, and indeed the causal generation of large-scale adiabatic perturbations requires a period of inflation.\\cite{L95}  In studying the microwave background anisotropies, we believe we have a tool through which all of the cosmological parameters, such as the Hubble constant $h$ and the density parameter $\\Omega_0$, can be probed. This is because the evolution of perturbations in the matter and radiation fields depends on all the cosmological parameters. On the other hand, the microwave background anisotropies are entirely due to linear perturbation theory, and hence entirely dependent on the initial perturbations you have in the first place. The only reason that one can indeed hope to extract cosmological parameters is because one believes that the initial spectrum takes on a simple form, perhaps a power-law, which can be parametrized simply and then those initial parameters thrown into the melting pot along with the cosmological ones and fitted by the data. One of the delights of inflation is that the initial spectrum is indeed typically predicted to take on a simple form, and furthermore one which can readily be calculated to high precision for a given inflationary model. ",
        "conclusions": "Inflation as a paradigm is both eminently and imminently testable by upcoming microwave background observations. For example, the prediction of a peak structure is extremely generic and quite specific to the situation where perturbations begin their evolution on scales much larger than the Hubble radius, and details such as the peak spacing promises a very strong test~\\cite{HW}. Something as simple as an observed spectrum without multiple peaks appears sufficient to rule out inflation (see e.g.~Barrow and Liddle~\\cite{BL97}). If inflation passes these tests, then detailed fitting to the observations promises startlingly high quality information about the inflationary mechanism.  However, the main purpose of my presentation is to provide a reminder of the important role inflation has in underpinning the microwave background endeavour. I stressed at the start that we can only get highly quality constraints from the present radiation power spectrum if we have a simple form, preferably motivated by theory, for the initial perturbations. Since the observations aim to be accurate at the percent level, the input information needs this accuracy too, and inflationary theory is now in a position where predictions at this level of accuracy can be made for all known models.  This can be contrasted with the situation for topological defect models, where it has proven much harder to make accurate theoretical predictions. Less accurate theoretical predictions will naturally lead to much more poorly determined cosmological parameters. [In fact, Pen (these proceedings) has also argued that in a defect model the observed spectrum is less sensitive to the cosmological parameters, implying poorer parameter estimation even if the theoretical calculations can after all be made more accurate.]  If all goes well with the observations, and inflation proves to be right, we can indeed look forward to the tiny error bars one hears about so often. If inflation is not correct, the results will still be spectacular but yet after the hype the constraints may seem disappointing."
    },
    "9803/astro-ph9803291_arXiv.txt": {
        "abstract": "We report on the analysis of the ASCA and BeppoSAX X--ray observations of the polar system BL Hyi, performed in October 94 and September 96, respectively.  Emission from both poles is apparent from the folded light curves of both observations; the emission from the second pole varies from cycle to cycle, indicating non--stationary accretion there.  The temperature of the post--shock region is estimated to be about 10 keV. Inclusion of both complex absorption and Compton reflection significantly improves the quality of the fit. No soft X--ray component is observed; the BeppoSAX/LECS upper limit to the soft component is in agreement with theoretical expectations for this low magnetic field system. ",
        "introduction": "Polar systems, a subgroup of magnetic Cataclysmic Variables (mCVs), contain a highly magnetized white dwarf with polar field strengths ranging from  $\\sim$ 10 MG $\\sim$ to 230 MG (see Beuermann 1997 and references therein), and accreting material from a late type main sequence star. The magnetic field of the white dwarf is strong enough to phase--lock its rotation with the orbital period. These systems are strong X-ray emitters in both soft and hard X-ray bands (see review by Cropper 1990).  While hard X-rays are emitted from a standing shock above the white dwarf surface, soft X-rays originate from hot photospheric regions heated either by irradiation from the post-shock plasma (Lamb \\& Masters 1979) or by dense plasma blobs carrying their kinetic energy deep into the atmosphere (Kuipers \\& Pringle 1982). Irradiation is important primarily for flow rates sufficiently low for the shock to stand high above the surface. However, a large fraction of the reprocessed radiation appears in the UV rather than at soft X-rays, as shown quantitatively in the case of AM\\,Her (G\\\"ansicke et al. 1995). The ``blobby'' accretion mode is increasingly important at higher field strengths, because the blobs become more and more compressed when arriving at the surface of the white dwarf (e.g. Beuermann \\& Burwitz 1995; King 1995 and references therein).  While the soft X--ray component is in general adequately fitted by a black--body spectrum with a temperature of a few tens of eV (even if more sophisticated models are sometimes needed: see e.g. Van Teeseling et al. 1994), it has recently become clear that a simple thermal plasma model is no longer adequate in reproducing the hard X--ray component of Polars. Reflection from the white dwarf surface, complex absorption and multi--temperature emission may contribute significantly to the spectrum above a few tenth of keV (e.g. Cropper et al. 1997 and references therein).   \\begin{table*} \\centering \\caption{ASCA and BeppoSAX observations log} \\label{log} \\vspace{0.05in} \\begin{tabular}{cccccc} \\hline \\hline ~ & ~ & ~ & ~ & ~ & ~\\cr Instrument & En. range & Date of obs. & Exp. Time$^a$ & Count rate & 2-10 keV flux$^b$ \\cr ~ & (keV)  & ~ & (ks) & (cts/s)  & (erg/s/cm$^2$) \\cr \\noalign {\\hrule} ~ & ~ & ~ & ~ & ~ & ~\\cr ASCA/GIS2 & 0.8-10 & 1994 Oct 11-12 & 43 & 0.13 & 8.0$\\times$10$^{-12}$ \\cr ASCA/GIS3 & 0.8-10 & 1994 Oct 11-12 & 43 & 0.16 & ~ \\cr ASCA/SIS0 & 0.5-10 & 1994 Oct 11-12 & 42 & 0.21 & ~ \\cr ASCA/SIS1 & 0.5-10 & 1994 Oct 11-12 & 41 & 0.16 & ~ \\cr BeppoSAX/LECS & 0.1-10 & 1996 Sep 27 & 3.7 & 0.02 & ~ \\cr BeppoSAX/MECS & 1.5-10 & 1996 Sep 27 & 12.5 & 0.12 & 6.8$\\times$10$^{-12}$ \\cr ~ & ~ & ~ & ~ & ~ & ~\\cr \\hline \\hline $^a$ after applying the selection criteria\\cr $^b$ single--temperature plasma model \\end{tabular} \\end{table*}   X--ray emission from the polar system BL Hyi (H01319-68) was first detected by HEAO-1/A-2, and later on by Einstein (Agrawal et al. 1983), EXOSAT (Beuermann \\& Schwope 1989, hereafter B\\&S89) and ROSAT (Ramsay et al. 1996; Schwope et al. 1997). In particular, the HEAO-1 and EXOSAT observations showed that copious soft X-ray emission originates from the secondary pole in states of high accretion rate, while the EXOSAT observations showed this pole to become increasingly less active at lower accretion rates. This suggests a picture in which the source switches from one-- to two--pole accretion in going from low to high states (B\\&S89). Recent EUVE observations indicate that the soft X--ray emission is best reproduced by a 17 eV blackbody (Szkody et al. 1997).  Despite all the above observations, the hard X--ray spectrum of BL Hyi was still poorly known before the launch of ASCA and BeppoSAX satellites. ASCA observed this source in October 1994, while BeppoSAX observed it in September 1996, in the context  of a Core Program devoted to study flux and spectral variability of Polars on different timescales and on a wide energy range.  In this paper we present temporal and spectral analysis of both ASCA and BeppoSAX observations. Sec. 2 describes the observations and data reduction, while in Sec. 3 and Sec. 4 we present temporal and spectral results, respectively, which are then summarized and discussed in Sec. 5. ",
        "conclusions": "During the ASCA and BeppoSAX observations the source was significantly brighter than during the `intermediate state' of October 1985, when it was observed by EXOSAT (B\\&S89). Assuming a 10 keV thermal plasma spectrum, the 2--10 keV phase--averaged flux, estimated with the {\\sc pimms} package, was in fact $\\sim$3$\\times$10$^{-12}$ erg cm$^{-2}$ s$^{-1}$ during that EXOSAT observation, while it was $\\sim$2.3 and 2.7 times brighter during the BeppoSAX and ASCA observations, respectively. The ASCA flux is in turn slightly lower than that measured by Ariel V ($\\sim$10$^{-11}$  erg cm$^{-2}$ s$^{-1}$, Agrawal et al. 1983).  Comparing both ASCA and BeppoSAX folded light curves with the EXOSAT ones (and in particular with that of October 1985, for which the statistics were the best), one important difference is apparent: there is significant emission from the second pole, while this emission during the EXOSAT observation was less intense, even if still present. As clearly evident from the BeppoSAX light curve, the emission from the second pole is highly variable, changing from cycle to cycle, and then suggesting non--stationary accretion.  The ASCA and BeppoSAX hard X--ray spectra are the first of sufficient quality to reveal spectral complexities. A simple thermal plasma model, even if formally acceptable, is unsatisfactory due to both the presence of features (mostly at the lowest energies) in the residuals and the high temperature and metal (mainly iron) abundance obtained. The inclusion of complex absorption cures the low energy features. A partial covering model, with a column density of about 3$\\times$10$^{21}$ cm$^{-2}$ and a covering fraction of $\\sim$0.4 provides an acceptable description of the spectrum.  It is conceivable that the absorber in front of the hard X-ray component is the free-falling matter feeding the accretion spot. For example in V834 Cen, this material was identified by its Zeeman component shifted in wavelength as expected from the free-fall velocity (Schwope \\& Beuermann 1990). As the hard X-ray emitting post-shock plasma is probably located in front of the soft X-ray emitting sections of the photosphere, this absorbing material should affect the soft component as well. Due to substantial pre-ionization, however, it may be partially transparent to soft X-rays. On the other hand, the fact that Zeeman absorption from this material was seen in V834 Cen indicates that in that system neutral material is also present. Hence, the soft component may not emerge unhindered and its spectral properties may carry the signatures of the intervening absorber as well as those of the emitter. This has been taken into account when we have estimated the intensity of the soft component from the BeppoSAX data. The upper limit so obtained is about ten times that of the hard luminosity, if a  black--body temperature of 17 eV (Szkody et al. 1997) is assumed; if, instead, a temperature of 25 eV (Ramsay et al. 1995) is adopted the upper limit is much more tight, reducing to about the hard luminosity. We also computed the upper limit to the black--body to hot plasma flux ratio in the 0.1--2.4 keV, to make comparison with the theoretical expectation (see Fig.~8 in Beuermann 1997). The limits are 5 and 1.3 for the adopted black--body temperatures, consistent with the comparatively hard X-ray spectra encountered in low-field systems (Beuermann \\& Schwope 1994). This is especially true if the field strength of 12 MG measured by Schwope et al. (1995) is characteristic of the main accreting pole. Even a field strength of 23 MG reported by Ferrario et al. (1996) would still qualify BL Hyi as a low-field system, characterized by a strong hard X-ray component. Only at still higher field strengths is bremsstrahlung effectively suppressed by increasing cyclotron cooling (Woelk \\& Beuermann 1996, see also Beuermann 1997).  Finally, the inclusion in the spectral model of the reflection component from the white dwarf surface provides a significant improvement in the statistical quality of the fit and reduces the  abundance to roughly the solar value. This component has been already detected in a number of Polars (e.g. Done et al. 1995; Beardmore et al. 1995; Ishida et al. 1997; Cropper et al. 1997; Done \\& Magdziarz 1997) and is now recognized as an important ingredient of their X--ray spectrum. With all these components included, the best fit post--shock temperature results to be about 10 keV (with about 20\\% uncertainty) in both observations, consistent with the EXOSAT results (B\\&S89) for the same source and with the values usually found in Polars."
    },
    "9803/astro-ph9803258_arXiv.txt": {
        "abstract": "We derive a self similar solution for the propagation of an extreme relativistic (or Newtonian) radiative spherical blast wave into a surrounding cold medium. The solution is obtained under the assumption that the radiation process is fast, it takes place only in the vicinity of the shock and that it radiates away a fixed fraction of the energy generated by the shock. In the Newtonian regime these solutions generalize the Sedov-Taylor adiabatic solution and the pressure-driven fully radiative solution. In the extreme relativistic case these solutions generalize the Blandford-McKee adiabatic solution. They provide a new fully radiative extreme relativistic solution which is different from the Blandford-McKee fully radiative relativistic solution.  This new solution develops a hot interior which causes it to cool faster than previous estimates. We find that the energy of the blast wave behaves as a power law of the location of the shock. The power law index depends on the fraction of the energy emitted by the shock. We obtain an analytic solution for the interior of the blast wave. These new solutions might be applicable to the study of GRB afterglow or SNRs. ",
        "introduction": "Many astrophysical phenomena (SNRs, $\\gamma$-ray bursts - GRBs, AGN hot spots, etc.)  are believed to involve  radiative shock waves.  These shocks accelerate particles which emit the observed radiation.  In particular, it is widely accepted that the recently discovered GRBs afterglow results from an emission by relativistic shocks, created by the interaction between an initial ejecta and the interstellar medium. The recent observations of GRB afterglow have lead to  numerous attempts to model this phenomena.  If the radiative mechanisms are slow compared to hydrodynamics time scale, the blast wave evolution is adiabatic.  The propagation of such a  blast wave is described be a self similar solution:  The Sedov-Taylor solution (\\cite{sedov,taylor}) describes the Newtonian regime and the Blandford \\& McKee (1976) solution describes the extreme relativistic regime.  We call the solution radiative if the radiative mechanisms are fast compared to the hydrodynamic time scale.  A fully radiative solution is one in which all the energy generated by the shock is radiated away. \\cite{OM88} have shown that if a fully radiative blast wave emits its energy only in the vicinity of the shock, it can be described by one of two possible self-similar solutions:  The pressure driven snow-plow (PDS) or the momentum conserving snow-plow (MCS) solution. In the PDS solution a thin shell ``snow plows'' through the external medium, driven by the pressure of its hot, roughly isobaric, interior. In the MCS solution the interior has been cooled and a thin shell slows down while conserving momentum. \\cite{bm} have found an extreme relativistic fully radiative solution. This solution describes a thin shell with a cold interior and it can be considered as the relativistic generalization of the Newtonian MCS solution. However, momentum is not conserved in this solution, as in the relativistic case one has to consider the momentum carried by the emitted radiation.  These solutions are either adiabatic or fully radiative. It is likely that in some cases not all the energy would be radiated away, even though the cooling is fast.  Such is the situation if the shock distributes the internal energy between the electrons and protons and there is no coupling between the two afterwards.  Since only the electrons cool, only a fraction of the internal energy will be radiated. It is likely that at least in the initial phases of a GRB afterglow this would be the case. We consider here this ``partially radiative'' scenario.  The goal of our paper is to find an analytic solution for the evolution of Newtonian and relativistic partially radiative blast waves. This, we believe, will eventually lead to a physical description of the evolution of GRB afterglows in their non-adiabatic stage, and of other astrophysical phenomena. We find a self similar solution under the assumption that the radiative mechanism is fast (compared to the hydrodynamics time scale).  We parameterize the radiation by a dimensionless (the non dimensionality is essential for a self-similar solution) parameter, $\\epsilon$, which describes the fraction of the energy produced by the shock that is radiated away.  We recover the extreme relativistic adiabatic Blandford-McKee solution and the Newtonian Sedov-Taylor solution in the corresponding limit when $\\epsilon = 0$.  In the Newtonian case when $\\epsilon \\to 1$ we recover the PDS solution, but we do not reproduce the MCS solution. Similarly in the relativistic limit of $\\epsilon \\rightarrow 1$ the solution approaches a new relativistic PDS solution which is different from the  radiative Blandford-McKee solution.  We describe our model in Sec. \\ref{s:model}. It is composed of an adiabatic shock followed by a narrow radiative region and a wide self-similar adiabatic regime.  We proceed in Sec. \\ref{s:newtonian_cond} by calculating the radiative shock conditions for Newtonian blast waves. We describe the self similar equations and their solution in Sec. \\ref{s:newtonian_self}. Using the same method we turn in Sec. \\ref{s:rel} to the calculations of relativistic blast waves.  Finally, in Sec. \\ref{s:others_results} we discuss the limiting solutions obtained by other authors and compare them with our solution. ",
        "conclusions": "We have solved the hydrodynamical evolution of blast waves in which each shocked particle emits a fixed fraction of the energy it gains in the shock. We have divided  the blast wave into three regions: the adiabatic shock, the radiative layer, and the adiabatic interior. The adiabatic shock and the radiation layer were  combined to form a ``radiative shock'', which set the boundary conditions for the adiabatic interior.  For fast cooling cases we have obtained a solution for a planar radiative layer with arbitrary shock velocities, independent of the cooling process. We have found that in the shock frame the fluid cools, slows down and its pressure {\\it increases} during the cooling process.  We have obtained  self similar solutions for adiabatic interior for the Newtonian and the extreme relativistic cases. We find that radiation can change the hydrodynamics considerably. Because of this change in the blast wave evolution the luminosity  decays  slower (in time) than what was estimated earlier assuming that the self similar profiles are independent of the radiated energy.  We have also found that as a blast wave becomes more radiative, its matter concentrates near the shock and forms a dense shell. However, the internal pressure does not drop to zero. In the fully radiative limit of Newtonian blast waves we have reached the pressure-driven solution. We have obtained a new extreme relativistic solution for fully radiative blast wave, which resembles the Newtonian PDS solution.  Our self similar solutions reach this modified solution in the fully radiative limit.  This solution does not correspond to the Blandford-McKee radiative solution.  The pressure must be  continuous within the adiabatic interior and it cannot develop self a similar shock or a rarefraction wave (apart from the main strong shock with the ISM). Therefore, even in the Newtonian PDS limit with isobaric interior, the solution contains a self similar transition layer between the ``radiative shock'' conditions and the internal pressure.  The recently discovered X-ray, optical and radio emission following a GRB, so called ``Afterglow'' is widely believed to be the result of the deceleration of a relativistic material that collides with the surrounding matter.  According to the common model, the shock wave produced by the collision accelerates electrons to relativistic velocities. These electrons, that carry a fixed fraction of the internal energy produced by the shock emit synchrotron radiation which is the observed afterglow. For reasonable parameters, the electron cooling is fast (at least in the early stages of the afterglow evolution), so that the electrons loose most of their energy. Since the density behind the shock is low, the coupling between the electrons and the protons is negligible.  Therefore, the protons energy is not radiated away. This leads to a partially radiative rather than a fully radiative shock, followed by an adiabatic flow. This is {\\it exactly} the scenario that leads to the partially radiative blast waves derived in this paper. Our new self similar partially radiative blast wave can serve as the basic hydrodynamic solution from which spectra and light curves of the afterglow can be calculated.  In particular we derive a new   $\\Gamma(R)$ relation (Eqs. \\ref{r:m_of_eps}, \\ref{G_of_t} ) between the Lorentz factor and the radius. This is the critical relation that determines most afterglow observations.  This research was supported by a US-Israel BSF grant 95-238 and by a NASA grant NAG5-3516. Re'em Sari thanks the Clore Foundations for support."
    },
    "9803/hep-ph9803500_arXiv.txt": {
        "abstract": "The supersymmetric generalization of the limiting and Gaussian QCD spectra is obtained. These spectra are valid for $x \\ll 1$, when the main contribution to the parton cascade is given by gluons and gluinos. The derived spectra are applied to decaying superheavy particles with masses up to the GUT scale. These particles can be relics from the Big Bang or produced by topological defects and could give rise to the observed ultrahigh energy cosmic rays. General formulae for the fluxes of protons, photons and neutrinos due to decays of superheavy particles are obtained. ",
        "introduction": "The spectra of hadrons produced in deep-inelastic scattering and $e^+e^-$ annihilation are formed due to QCD cascading of the partons. In the Leading Logarithmic Approximation (LLA) which takes into account $\\ln(Q^2)$ terms this cascade is described by the Gribov-Lipatov-Altarelli-Parisi-Dokshitzer (GLAPD) equation \\cite{GLAPD}. The Modified Leading Logarithmic Approximation (MLLA) takes into account additionally $\\ln(x)$ terms, where $x=k_{\\parallel}/k_{\\parallel}^{max}$ and $k_{\\parallel}$ is the longitudinal momentum of the produced hadron. Color coherence effects are described in MLLA. Two approximate analytic solutions to the MLLA evolution equations have been obtained. % These are the {\\em limiting spectrum\\/} \\cite{MLLA} and the {\\em Gaussian spectrum\\/} \\cite{Mu,DFK}, in which we include the {\\em distorted Gaussian spectrum\\/} \\cite{distG,dkmt,esw}. The limiting spectrum is the most accurate one among them. In fact, it describes well the experimental data at large $x$, too, and this is natural, though accidental (for an explanation see Ref.~\\cite{dkmt}). The limiting spectrum has a free normalization constant, $K_{\\rm lim}$, which cannot be calculated theoretically and is found from comparison with experimental data. This constant has to be considered as a basic parameter of the theory, and it can be used at all energies, where the physical assumptions, under which the limiting spectrum is derived, are valid.  For detailed calculations of hadron spectra in $e^+e^-$ annihilation and comparison with experimental data see Ref.~\\cite{comparison}. Up to energies of existing $e^+e^-$ colliders, $\\sqrt{s} \\lsim 183$~GeV, the limiting spectrum and the distorted Gaussian spectrum describe well the available data. At large energies $\\sqrt{s}\\gsim 1$~TeV the production of supersymmetric particles might substantially change the QCD spectra. Apart from future experiments at LHC, supersymmetry (SUSY) might strongly reveal itself in the decays of superheavy particles. They can appear as relics of the Big Bang, or be produced by topological defects (TD), and can be the sources of the observed ultrahigh energy cosmic rays (UHECR) at $E \\gsim 1\\cdot 10^{10}$~GeV. The range of masses, $m_X$, of interest for UHECR goes from the GUT scale ($m_X \\sim 10^{16}$~GeV) or less down to $m_X \\sim 10^{12} - 10^{14}$~GeV.  In this paper we obtain the generalization  of the limiting and Gaussian QCD spectra for the SUSY case. Although the influence of SUSY on, e.g., the evolution of parton distributions or the running coupling constant was considered in many works in the 80's \\cite{split}, this is, to the best of our knowledge, the first time that fragmentation spectra of hadrons are examined for large $\\sqrt{s}$ up to the GUT scale in SUSY-QCD. As application we use these spectra for calculations of the fluxes of UHECR produced in the decays of superheavy particles. ",
        "conclusions": ""
    },
    "9803/astro-ph9803203_arXiv.txt": {
        "abstract": "The white dwarf stars WD\\,1614$+$136 and WD\\,1353$+$409 are not sufficiently massive to have formed through single star evolution. However, observations to date have not yet found any evidence for binarity. It has therefore been suggested that these stars are the result of a merger. In this paper we place an upper limit of $\\approx 50\\kms$ on the projected rotational velocities of both stars. This suggests that, if these stars are the results of a merger, efficient angular momentum loss with accompanying mass loss must have occurred. If the same process occurs following the merging of more massive white dwarf stars, the predicted rate of Type~Ia supernovae due to merging white dwarfs may have been greatly over-estimated. Further observations to determine binarity in WD\\,1614$+$136 and WD\\,1353$+$409 are therefore encouraged. ",
        "introduction": "The finite age of the universe limits the time available for  stars to evolve. The lowest mass stars have had insufficient time to evolve to the point where the remnant is seen as a white dwarf star. There is a monotonically increasing relationship between the initial mass of the star and the resulting white dwarf star and so this implies a minimum mass for white dwarf stars. The exact value is uncertain, but is around 0.55\\Msolar (Bragaglia, Renzini \\& Bergeron, 1995). Nevertheless, white dwarf stars are found with masses well below this limit. These white dwarfs are the consequence of binary star evolution in which the evolution of a star through the red giant phase is interrupted by a common--envelope phase in which a companion star is engulfed by the expanding envelope and then rapidly spirals in towards the core, ejecting the envelope. This arrests the formation of the degenerate red giant core resulting in an anomalously low mass white dwarf star.  Dramatic evidence for this scenario was provided by Marsh, Dhillon \\& Duck \\shortcite{Marsh95}. They observed  7 DA  white dwarfs selected for the low mass derived for them by Bergeron, Saffer \\& Liebert \\shortcite{Berg92} from their spectra. Marsh et~al.\\ were able to measure radial velocities with accuracies of a few \\kms\\ by using the narrow core of the \\halpha\\ absorption line. Periodic radial velocity variations showed at least 5 of the 7 stars to be binary stars. No evidence for binarity was found for WD\\,1614$+$136 or WD\\,1353$+$409.  It should be emphasized that the observations of Marsh et~al.\\ cannot rule out the possibility that these two white dwarf stars are binaries. Although a main--sequence companion more massive than 0.1\\Msolar\\ can be ruled out in both cases, the presence of another cool white dwarf, very low mass M dwarf or a brown dwarf, perhaps in a low inclination orbit, cannot be ruled out. With such strong evidence for the scenario outlined above it would seem to be inevitable that these white dwarfs were once members of binary systems. The nature of any companion, or its fate if it is no longer present, remain open questions.  The failure of Marsh et~al.\\ to detect binarity in WD\\,1614$+$136 and WD\\,1353$+$409 led Iben, Tutukov \\& Yungleson \\shortcite{Iben97} to suggest these stars are now single stars that are the result of a merger between  a white dwarf and the companion responsible for the common--envelope phase. In this paper we show that the detection of a narrow core to the \\halpha\\ makes this suggestion very unlikely unless angular momentum loss from the merger product is extremely efficient. ",
        "conclusions": "If WD\\,1614$+$136 and WD\\,1353$+$409  are now genuinely single stars, their low projected rotational velocity implies efficient angular momentum loss following the merging of these stars with the companion star responsible for the common envelope phase. If the accompanying mass loss occurs following the merging of more massive white dwarfs, their role as the progenitors of Type~Ia supernovae is put into doubt."
    },
    "9803/astro-ph9803035_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {\\bf Key words:~~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {We present a new distance determination to the Large and Small Magellanic Clouds using the newly developed red clump stars method (Paczy\\'nski and Stanek 1998). This new, single-step, Hipparcos calibrated  method seems to be one of the most precise techniques of distance determination  with very small statistical error due to large number of red clump stars  usually available.  The distances were determined independently along four lines-of-sight located at opposite sides of each Magellanic Cloud. The results for each line-of-sight are  very consistent. For the SMC we obtain the distance modulus: $m-M=18.56\\pm0.03\\pm0.06$~mag (statistical and systematic errors, respectively) and  for the LMC: ${m-M=18.08\\pm0.03\\pm0.12}$~mag where systematic errors are mostly due to uncertainty in reddening estimates. Both distances will be refined  and systematic errors reduced when accurate reddening maps for our fields are available.  Distance moduli to both Magellanic Clouds are ${\\approx0.4}$~mag smaller than  generally accepted values. The modulus to the LMC is in good agreement with  the recent determinations from RR~Lyrae type stars and upper limit resulting  from the SN1987A echo. We suspect that the distance to the LMC and SMC is  shorter by about 15\\% than previously assumed: 42~kpc and 52~kpc, respectively. Calibrations of the period-luminosity relation for Cepheids which give  overestimated distances to the LMC and SMC are probably incorrect and require urgent reanalysis.  We also present our color-magnitude diagrams around the red clump for the LMC  and SMC. We identify vertical red clump, first noted by Zaritsky and Lin (1997),  in the color-magnitude diagram of both Magellanic Clouds and we interpret it as an evolutionary feature rather than unknown stellar population between the  LMC and our Galaxy.}{Magellanic Clouds -- Galaxies: distances and redshifts -- Hertzsprung-Russel (HR) diagram} ",
        "introduction": "% Two of our closest neighboring galaxies from the Local Group -- the Magellanic  Clouds -- belong to the most important objects of the modern astrophysics.  Hosting many objects commonly used as standard candles at practically the same  distance they serve as ideal targets for calibrating the distance scale in the  Universe. Therefore it is crucial to have well established distances to both  Magellanic Clouds.  Until recently, that is in the pre-Hipparcos era, it was generally accepted  that the distance modulus to the Large Magellanic Cloud was in the range of  ${m{-}M{=}(18.5{\\div}18.6)}$~mag based mostly on the calibration of the period-luminosity  (P--L) relation for Cepheids (\\eg Laney and Stobie 1994, Gieren \\etal 1997 and  references therein) and measurements of the supernova SN1987A echo (\\cf Panagia \\etal 1991, Sonneborn \\etal 1997). In 1992 Walker (1992) noted  ${\\approx0.3}$~mag discrepancy between absolute magnitudes of RR~Lyrae type stars from the Galaxy and the LMC at distance of 18.5~mag. Based on the new  statistical parallax solution for RR~Lyrae, Layden \\etal (1997) recalibrated  the absolute magnitude-metallicity relation for these stars and found that the resulting distance modulus to the LMC might be significantly lower: ${m- M=18.28 \\pm0.13}$~mag.  The latest determinations of the distance to the LMC come from Hipparcos recalibration of the P--L relations for different type of pulsating stars. Both Cepheid and Mira recalibrations lead to larger distance moduli: Cepheids --  ${18.70\\pm0.1}$~mag (Feast and Catchpole 1997) or ${18.57\\pm0.11}$~mag (Madore and  Freedman 1998) and Miras -- ${18.54\\pm0.18}$~mag (van Leeuven \\etal 1997). The value ${m-M=18.50\\pm0.10}$~mag has been adopted by the HST Extragalactic Distance Scale Key  Project team (\\eg Ra\\-wson \\etal 1997 and other papers of the series) as the  distance to the LMC to which extragalactic, Cepheid-based distance scale is  tied.  On the other hand Hipparcos recalibration for RR~Lyrae based both on direct  parallax determination and statistical parallaxes gives results fully  confirming Layden \\etal (1997) determination: ${18.26\\pm0.15}$~mag (Fernley  \\etal 1998). Similar results from statistical parallaxes were recently  obtained by Popowski and Gould (1998). Thus it seems very likely that the  RR~Lyrae stars give ${\\approx0.25}$~mag shorter distance scale to the LMC.  Meanwhile, recent reanalysis of the supernova 1987A echo by Gould and Uza (1998) also suggests smaller distance modulus than previously derived ({\\it cf.}  Sonneborn \\etal 1997, Lundquist and Sonneborn 1997) with the upper limit as  low as ${m-M<18.37\\pm0.04}$~mag.  The distance to the Small Magellanic Cloud is poorly known. The distance modulus to the SMC is assumed to be about 18.9~mag (Westerlund 1990). Massey  \\etal (1995) obtained ${19.1\\pm0.3}$~mag from spectroscopic parallaxes. The most  recent determination based on the Cepheid period-luminosity relation gives  similar value: ${m-M=18.94}$~mag (Laney and Stobie 1994).  Summarizing this short review, the distance to the Large Magellanic Cloud is at  least controversial and the generally accepted value $m-M\\approx18.5$~mag cannot  be now treated as established beyond any doubt. The distance to the Small  Magellanic Cloud is assumed to be ${m-M\\approx18.9}$~mag, still poorly known, and  therefore any new, independent information about both distances is of the  greatest importance.  Recently Paczy\\'nski and Stanek (1998) proposed a new, single step method of  distance scale determination to the objects in the Galaxy and neighboring  galaxies. The method bases on the fact that the red clump giant stars seem to  be a very homogeneous group of objects and their mean magnitude in the {\\it I}-band is practically independent of their color. Therefore the red clump  stars might be considered as excellent standard candle candidates. The red  clump stars are very numerous compared with other objects used as standard candles so far (typically several orders of magnitude more numerous than \\eg  pulsating stars). Thus, their mean magnitudes can be obtained very precisely  with small statistical error making the method very attractive for precise determination of distances in the Universe.  The mean absolute magnitude of the red clump stars from the solar vicinity was  determined by Paczy\\'nski and Stanek (1998) who analyzed luminosities of about  600 such stars with precise parallaxes (accuracy better than 10\\%) measured by  Hipparcos (${M^{\\rm loc}_{I_0}=-0.185\\pm0.016}$~mag). Paczy\\'nski and Stanek (1998) applied the method for determination of the distance to the center of  the Galaxy comparing $M^{\\rm loc}_{I_0}$ with the mean {\\it I}-band magnitude of  the red clump stars from the Baade's Window obtained from photometry of the  first phase of the Optical Gravitational Lensing Experiment (OGLE). The new method was also applied by Stanek and Garnavich (1998) for determination of the distance to M31. Stanek and Garnavich (1998) also refined the maximum of the  local red clump stars luminosity function by limiting the Hipparcos sample of  the red clump stars to be volume rather than luminosity limited and derived  ${M^{\\rm loc}_{I_0}=-0.23\\pm 0.03}$~mag from 228 red clump stars located  within the distance ${d<70}$~pc. Very large number of red clump stars in the  Hipparcos sample makes the calibration very sound comparing with  calibrations of other standard candles. It should be noted that this  calibration might be even more precise in the future when good {\\it VI}  photometry is derived for more Hipparcos stars.  Although employing the red clump stars as distance indicator seems to be very attractive and can be possibly one of the most precise methods of distance determination one should be aware of potential systematic errors which can lead to errors in determined distance scale. The most important is proper determination of the interstellar extinction which can affect both target red clump luminosity as well as the local Hipparcos sample. However, the latter sample is very likely extinction free for ${d<70}$~pc or it is affected negligibly, \\eg that one analyzed by Stanek and Garnavich (1998). Nevertheless precise determination of  extinction to the target stars is very important to avoid large systematic  errors.  The second source of errors can be the differences in populations -- ages,  chemical composition etc., between the red clump stars of the target object  and those in the solar neighborhood. This is, however, a common problem of any  other method of distance determination when similar objects from different  locations of the Universe are compared. In case of the red clump stars  theoretical models predict that their luminosity is very weakly dependent on  chemical composition and age (Castellani, Chieffi and Straniero 1992, Jimenez,  Flynn and Kotoneva 1998). Nevertheless, the empirical fact of the very small  dispersion of {\\it I}-band magnitudes of the red clump stars as observed in the  solar vicinity, Baade's Window and M31 (Paczy\\'nski and Stanek 1998, Stanek and Garnavich 1998) should find theoretical explanation which could also answer the question about uncertainties introduced by comparison of different  populations of the red clump stars.  The natural consequence of the successful application of the red clump method  for determination of distances to the Galactic center (Paczy\\'nski and Stanek  1998) and M31 (Stanek and Garnavich 1998) is a distance determination to other  objects in the Local Group. The most natural candidates are the Magellanic  Clouds, which are targets of the microlensing surveys providing huge databases  of photometric data for millions of stars in both Magellanic Clouds.  The OGLE project is a long term microlensing survey which second phase --  OGLE-II started at the beginning of 1997 (Udalski, Kubiak and Szyma\\'nski  1997). The Magellanic Clouds are the new targets of the OGLE-II phase. Unlike  other microlensing surveys the OGLE project observations are carried out with  the standard {\\it BVI}-bands with majority (${\\approx75\\%}$) observations  obtained in the {\\it I}-band. Thus, the data can be precisely transformed to the  standard {\\it BVI} system and applied directly to many projects unrelated to  microlensing. In particular the OGLE observations are very well suited for the  above mentioned red-clump method as the most of data is collected in the {\\it I}-band.  In this paper we apply the red clump stars method for determination of distances to both Large and Small Magellanic Clouds. Although OGLE-II data  cover large fractions of both Magellanic Clouds we present here distance  determination for two most west- and eastward lines-of-sight toward both  Clouds where we assume extinction to be constant in the first approximation.  Thus, the results should be considered as preliminary and will be refined when  accurate maps of extinction based on the OGLE-II data are derived. ",
        "conclusions": "% Results of distance determination to the Magellanic Clouds with the red clump stars method (Table~3) lead to the surprising conclusion. All of our determinations are consistently ${\\approx0.4}$~mag smaller than generally  accepted distance moduli to the Magellanic Clouds.  There might be a few reasons of the discrepancy:  -- error of the zero point of the OGLE photometry. This is, however, extremely  unlikely. In Section~2 we discussed accuracy of the OGLE photometry and showed comparison with other reliable CCD photometries of the Magellanic Clouds (\\eg  Fig.~1). The OGLE mean magnitudes are the average of tens of individual observations  and were tied with the standard system based on observations collected on many  nights with hundreds of observations of standard stars. We believe that our photometry constitutes a huge set of secondary {\\it BVI} standards in both  Magellanic Clouds and large systematic errors are excluded.  -- errors in reddening estimates for our fields. Certainly overestimating of  the reddening leads to smaller distance moduli. However, one should note that  even with the zero reddening the distance moduli derived with the red clump method are  smaller than generally accepted for the Magellanic Clouds (by at least  0.1~mag). But extinction to the Magellanic Clouds does exist and for the LMC  the foreground extinction is at least ${E(B-V)=0.075}$~mag, ${A_I=0.15}$~mag, and to the  SMC ${E(B-V)=0.037}$~mag, ${A_I=0.07}$~mag, (Schlegel, Finkbeiner and Davis 1998). Reddening  to the SMC is much smaller and our distance determination is less affected by  the reddening estimate. But consistently the distance modulus to the SMC is  also smaller by ${\\approx0.4}$~mag.  -- the red clump stars method. We addressed this possibility in the Introduction. It seems unlikely from our present understanding of the red  clump stars, but in principle, the mean {\\it I}-band magnitude of the red clump stars might  be a function of stellar population (although it shows very small dispersion  in all objects investigated so far: local sample, Galactic bulge, M31 and  Magellanic Clouds). We stress here once again that this problem must be  carefully  investigated both theoretically and observationally to answer  potential questions.  As any of the above reasons does not seem to be satisfactory, we come to the  very tempting conclusion that the distance determination to the Magellanic  Clouds with the red clump method is correct, the distance moduli are ${\\approx  0.4}$~mag smaller than those generally accepted and both Magellanic Clouds are  located about 15\\% closer to our Galaxy than previously assumed: the LMC at 42~kpc, and the SMC at 52~kpc.  Our conclusion seems to be a strong argument in favor of the shorter distance to the LMC. Our determination is based on the independent, very straightforward, single-step method with minimum possibilities for the potential error sources. The new technique is calibrated with Hipparcos data with much better accuracy than calibration of any other standard candle (orders of magnitude more stars, very close objects with precise parallaxes). Independent determinations for four lines-of-sight located at different sides of each Magellanic Cloud yield very consistent results which makes them very sound. If we underestimated the reddening, the Magellanic Clouds are even closer to us. On the other hand if we assume only minimum, foreground reddening, which is certainly an underestimate, ($E(B-V)=0.075\\pm0.015$~mag for the LMC and $E(B-V)=0.037\\pm0.015$~mag for the  SMC, Schlegel, Finkbeiner and Davis 1998) the upper limits for distance moduli  are: $m- M<18.29\\pm0.03\\pm0.03$~mag for the LMC and $m-M<18.65\\pm0.03\\pm0.03$~mag for the SMC. Still a few sigma gap between distance moduli resulting from the red clump method and presently accepted values  remains.  Brocato \\etal (1996) presented the CCD photometry of selected globular  clusters in the LMC. We find there that one of the CMDs, showing the region around  NGC~1786, contains a neat red clump. Unfortunately Brocato \\etal (1996) do not  provide their photometry, but even visual inspection of their Fig.~1 indicates  that the mean magnitude of the red clump stars is ${I=18.2\\pm0.1}$~mag. NGC~1786 is located only one degree north from our field LMC$\\_$SC15N so it is  comparably affected by extinction (it falls almost into the extinction map of  Harris, Zaritsky and Thompson 1997). Thus resulting distance modulus from  NGC~1786 field and independent photometry is in perfect agreement with our red  clump determinations in Table~3.  As we mentioned in the Introduction, RR~Lyrae-based distance determination  also seems to favor smaller distance modulus to the LMC. All recent  determinations from statistical parallaxes and Hipparcos calibrations of  RR~Lyrae stars are within one sigma agreement with our determination. Also the  upper limit for the distance to the LMC from the SN1987A light echo by Gould and Uza (1998) is compatible with our determination.  The only arguments for the larger distance to the LMC come then from the Cepheid P--L calibration (and the Mira variables but because of much lower accuracy this determination is much less reliable). However, because of its complex nature, the calibration might be simply incorrect. An incorrect account of possible population effects, mainly metallicity, could result in overestimated distances obtained with that method (\\cf Sekiguchi and Fukugita 1998, Sasselov \\etal 1997). Our SMC distance determination also confirms this thesis: our distance modulus to the SMC is ${\\approx0.4}$~mag smaller than that derived from Cepheid P--L (Laney and Stobie 1994). It should be also noted that there exists ${\\approx0.3}$~mag discrepancy between Hipparcos calibrated Cepheid P--L distance to M31 (Feast and Catchpole 1997) and recent determination with the red clump stars (Stanek and Garnavich 1998) and globular clusters (Holland 1998). Thus we would like to stress at this point  the necessity of urgent reanalysis of the Cepheid P--L relation. Very  precise, {\\it BVI} light curves of hundreds of Cepheids from both Magellanic  Clouds will be provided by the OGLE project to the astronomical community later during this year and can be of great importance for such a project.  Concluding, it seems very likely that Cepheid P--L relation gives distance moduli overestimated by about ${\\approx0.4}$~mag (at least toward the Magellanic Clouds)  with all astrophysical consequences -- distances beyond M31, up to several Mpc  rely mainly on Cepheids and the LMC distance. Our independent distance measurement to the LMC  confirms the recent calibrations of the absolute magnitude--metallicity relation for RR~Lyrae stars giving fainter absolute magnitudes contrary to results of Chaboyer \\etal (1998) and therefore leading to older ages of globular clusters  (Layden \\etal 1996). On the other hand shorter distance to the LMC means  shorter distance to Cepheid-based galaxies and in turn larger Hubble constant  and shorter age of the Universe. The \"age problem\" gap increases.  We believe that there exists at least one crucial test which can verify our  present findings. Paczy\\'nski (1997) proposed detached eclipsing binaries as  an excellent method of distance determination. Having precise light curve and  spectroscopic data one can determine the distance to the eclipsing, well  detached system with accuracy of a few percent. The catalog of eclipsing stars  in the Magellanic Clouds suitable for such a distance determination with three  color light curves, periods etc.\\ will be released by the OGLE project in a few months. Faintness of the Magellanic Cloud stars will require the largest, 10-m  class telescopes to collect good quality spectroscopic data and a lot of  telescope time. However, importance of the problem makes such a project the one of the highest priority and we hope it will be undertaken soon.  \\Acknow{We would like to thank Prof.\\ Bohdan Paczy\\'nski for encouraging discussions and help at all stages of the OGLE project. We are indebted to Dr Krzysztof Stanek for useful coments and remarks. The paper  was partly supported by the Polish KBN grant 2P03D00814 to A.\\ Udalski.  Partial support for the OGLE project was provided with the NSF grant  AST-9530478 to B.\\ Paczy\\'nski.}  \\bigskip  {\\bf Note:} Based on theoretical isochrones fitting to the main sequence and the red clump in four fields in the LMC, Beaulieu and Sackett (1998, revised version of their paper) found significantly better fit assuming smaller distance modulus to the LMC (${m-M=18.3}$) providing thus, additional argument in favor of the shorter distance to the LMC.  \\newpage"
    },
    "9803/astro-ph9803345_arXiv.txt": {
        "abstract": "IRAS~06562$-$0337 has been the recent subject of a classic debate: proto-planetary nebula or young stellar object? We present the first 2$\\micron$ image of IRAS~06562$-$0337, which reveals an extended diffuse nebula containing approximately 70 stars inside a 30$\\arcsec$ radius around a bright, possibly resolved, central object.  The derived stellar luminosity function is consistent with that expected from a single coeval population, and the brightness of the nebulosity is consistent with the predicted flux of unresolved low-mass stars. The stars and nebulosity are spatially coincident with strong CO line emission. We therefore identify IRAS~06562$-$0337 as a new young star cluster embedded in its placental molecular cloud.  The central object is likely a Herbig Be star, $M \\approx 20 M_{\\odot}$, which may be seen in reflection.  We present medium resolution, high S/N, 1997 epoch optical spectra of the central object.  Comparison with previously published spectra shows new evidence for time variable permitted and forbidden line emission, including \\ion{Si}{2}, \\ion{Fe}{2}, [\\ion{Fe}{2}], and [\\ion{O}{1}].  We suggest the origin is a dynamic stellar wind in the extended, stratified atmosphere of the massive central star in IRAS~06562$-$0337. ",
        "introduction": "Garcia-Lario, Manchado, Sahu, and Pottasch (1993, hereafter GMSP) present the first detailed analysis of IRAS~06562$-$0337.  They argue that it is a proto-planetary nebula (PPN) undergoing final mass-loss episodes.  Their time-series of optical spectra, obtained over a 5 year period, show the onset of forbidden line emission and the possible evolution of the central star toward hotter temperatures. They derive a Zanstra temperature of 2$\\times$10$^{4}$ K, with a slight increase over a two year interval.  The effective temperature of the exciting star, T$_{eff}$ $\\approx$ 3.6$\\times$10$^{4}$ K, also showed a slight increase in two years. The H$\\alpha$ line profile changes in time, which GMSP interpret as variable high velocity winds associated with episodic mass-loss. The appearance of [\\ion{O}{3}] emission lines in 1990 and the resulting 4363/(4959+5007) line ratio requires an ionizing region of high electron density, log($n_{e}$) $\\approx$ 6.9. The absence of these lines in spectra obtained before and after 1990 is interpreted as collisional de-excitation due to changing densities in the ionized region effected by violent episodic mass-loss. From CO observations GMSP derive $V_{LSR}$ = 50 $\\pm$ 1 km sec$^{-1}$, which agrees with the velocity derived from their high resolution optical spectra. Adopting a model galactic rotation curve, they estimate a distance of 4 kpc, which compares to a distance of 2.4 kpc estimated from the equivalent width of Na D absorption seen in their spectra. The IRAS colors fit with blackbodies show a trend of decreasing temperature with increasing wavelength which implies a gradient of dust temperatures.  GMSP integrated the optical--IR spectral energy distribution of IRAS~06562$-$0337, yielding a luminosity of L = 7000 $L_{\\odot}$ for their preferred distance of 4 kpc.  Kerber, Lercher, and Roth (1996, hereafter KLR) describe an additional medium resolution, high S/N spectrum of IRAS~06562$-$0337 obtained in early 1996.  [\\ion{O}{3}] emission is still absent, but a wealth of \\ion{Fe}{2} and [\\ion{Fe}{2}] lines are found.  These lines confirm the high electron density derived by GMSP from the [\\ion{O}{3}] lines present in 1990.  KLR argue that the spectrum also implies a considerable density gradient in the object, as [\\ion{Fe}{2}] lines are collisionally suppressed at densities where \\ion{Fe}{2} lines exist.  They maintain the classification of IRAS~06562$-$0337 as a candidate PPN, designating it ``The Iron Clad Nebula''.  Bachiller, Gutierrez, and Garcia-Lario (1998, hereafter BGG) present new mm and sub-mm observations of IRAS~06562$-$0337. They derive $V_{LSR}$ = 54.0 $\\pm$ 0.2 km sec$^{-1}$ and adopt a different model Galactic rotation curve than GMSP to estimate a distance of 7 kpc. This distance yields a luminosity of 21000 $L_{\\odot}$ and a cloud mass M $>$ 1000 $M_{\\odot}$. From the strength of the CO emission and the presence of CS emission, BGG surmise that IRAS~06562$-$0337 is a ``young stellar object (or small cluster) still associated to its parent molecular cloud.'' BGG point out that IRAS~06562$-$0337 satisfies the three criteria for a Herbig Ae/Be star (Herbig, 1960) and the spectral energy distribution, which rises sharply in the far infrared (GMSP), is similar to Group II Herbig Ae/Be stars (Hillenbrand, 1992). They also note the presence of blue and redshifted wings in the CO emission indicating a bipolar outflow. The CO outflow may be driven by an eruptive ionized jet, which leads them to suggest the sporadic [\\ion{O}{3}] emission seen by GMSP originates in a Herbig-Haro object.  We present the first 2$\\micron$ image of IRAS~06562$-$0337. Our image reveals a compact cluster of stars surrounding a bright, central object.  We independently confirm the result also discovered by BGG that IRAS~06562$-$0337 is a young stellar object. In Section 2 of this paper, we describe our near-infrared observations and stellar census of the IRAS~06562$-$0337 cluster.  We also compare the CO(2$\\rightarrow$1) map of BGG with our image. In Section 3, we describe our new spectroscopic observations and summarize the resulting 1997 epoch emission line data.  We make a detailed comparison with the 1996 epoch emission line data of KLR. In Section 4, we present our conclusions. ",
        "conclusions": "The bright central object in IRAS 06562$-$0337 is probably a single Herbig Ae/Be class star, with $M \\approx 20 M_{\\odot}$. There is some evidence that it is seen at least partially in reflection. For example, the central star is poorly fit by the stellar point spread function derived from other stars in our K$^{\\prime}$ image. In addition, the centroids in the visible and at K$^{\\prime}$ are offset by $\\Delta\\alpha \\sim$0.6$\\arcsec$. A comparison of our 1997 epoch optical spectra to the 1996 epoch spectrum presented by KLR has provided new evidence for spectral line variability in IRAS~06562$-$0337 that is not easily interpreted as simple changes in temperature or density.  This argues against previous suggestions that the already-published spectral data reflect real-time evolutionary changes of the central star. Understanding the immediate environs of the central star is an essential first step to better understanding the emission line variability that has been observed.   We note that our analysis of the 1996--1997 spectral variability of IRAS 06562$-$0337 using simple zones and changes in one parameter (such as the electron temperature or electron density) is only intended as an interpretive tool and should be treated with some caution.  The inconsistencies of this analysis suggest the complexity of the actual emission line region.  If the central star is analagous to a Herbig Be star, then its ``fully ionized zone'' should have a very high electron density. Thus, the low density derived with the [\\ion{O}{1}] emission implies a stratification of densities, likely in the extended atmosphere of the central star in IRAS~06562$-$0337. The anti-correlation of [\\ion{O}{1}] and [\\ion{Fe}{2}] emission is not easily understood if they arise from a common partially ionized zone (perhaps located at the edge of a compact \\ion{H}{2} region around the central star), but is likely consistent with an origin in a dynamic stratified atmosphere. Spectral variability is also consistent with a stellar wind origin.  Stellar winds from Herbig Be type stars are highly variable, extremely complex, and poorly understood (e.g., they are possibly modulated by magnetic fields or non-radial pulsations; Catala et al. 1993). Wind line strengths and profiles can show nightly variations, some of which are correlated with temperature changes of the star (Scuderi et al. 1994).   The appearance of [\\ion{O}{3}] emission in IRAS 06562$-$0337 during 1990 (GMSP) can be understood as shocks in stellar jets from a variable stellar wind (Hartigan \\& Raymond 1993). Additionally, we emphasize that the size of the ionized region in IRAS 06562$-$0337 is small. Six cm radio observations yielded a surprising non-detection given the strong H$\\beta$ flux, and allow an upper limit angular diameter of $3.4\\times10^{-4}$ arcsec to be set for the ionized region (GMSP). If we assume the 7 kpc distance estimate of BGG, then the ionized region in IRAS 06562$-$0337 has an upper limit diameter of 2--3 AU. We conclude that stellar winds in the extended, stratified atmosphere of a very young, $M \\approx 20 M_{\\odot}$ star is the most plausible explanation for the puzzling spectral variability of IRAS~06562$-$0337, the ``Iron Clad Nebula.''    A valuable time series of spectra of the central object in IRAS 06562$-$0337, covering a 10 year interval, are now available in the literature.  A useful complementary study would obtain more densely sampled spectroscopic observations, to investigate the short timescale behavior of the emission.  In particular, correlated line variability and characteristic timescales could be used to constrain the dimensions of the ionized region, which we have suggested is likely to be an extended stellar atmosphere. We emphasize the need for high resolution and high S/N spectroscopic observations in the near-UV, and of the entire Balmer series, as these data can be used to study the extended atmospheres of Be stars (Burbidge \\& Burbidge 1953). It is important to build a consistent physical picture of the environs around the central star in order to answer questions such as ``is the central star seen in reflection?'' This problem in particular could be addressed with high resolution imaging and polarization data. Lastly, as a young star cluster, IRAS~06562$-$0337 is interesting and notable for its richness.  Multi-color near-infrared photometric observations would allow an accurate measurement of the IMF, and an investigation into problems such as mass segregation and star formation efficiencies (e.g., Barsony et al. 1991)."
    },
    "9803/astro-ph9803109_arXiv.txt": {
        "abstract": "Using time-resolved two-dimensional aperture photometry we have put upper limits on the pulsed emission from two proposed optical counterparts for PSR B1951+32.  Our pulsed upper limits of $m_{v,\\rm{pulsed}} > 23.3$, $m_{b,\\rm{pulsed}} > 24.4$, for the first candidate and $m_{v,\\rm{pulsed}} > 23.6$, $m_{b,\\rm{pulsed}} > 24.3$ for the second, make it unlikely that either of these is, in fact, the pulsar. We discuss three further candidates, but also reject these on the basis of timing results. A search of a $5\\farcs5 \\times 5\\farcs5$ area centred close to these stars failed to find any significant pulsations at the reported pulsar period. ",
        "introduction": "PSR B1951+32 is a  fast pulsar in the peculiar combination supernova remnant (SNR) CTB 80.  The characteristic spin-down age of the pulsar $(P/2\\dot P)$ and the dynamical age of the SNR (Koo et al.  1990) are both $\\sim10^5$ yrs.  Its low surface magnetic field strength, of $~5\\times10^7$ T, is unusual amongst the selection of young and middle-aged pulsars detected so far but it has been suggested that this is more representative of young pulsars in general (\\cite{lyne83}).  It is clearly important to test pulsar models with as wide a variety of pulsar parameters as possible.  Steady emission from the pulsar has so far been detected at radio (\\cite{strom87}; \\cite{kul88}) and X--ray (\\cite{beck82}; \\cite{wang84}) wavelengths.  Two possible optical counterparts have been proposed by Blair \\& Schild (1985) and Fesen \\& Gull (1985).  Since the initial discovery of the 39.5-ms pulsar (\\cite{kul88}), evidence for pulsed emission has also been found in X--rays (\\cite{safi95}) and $\\gamma$--rays (Ramanamurthy et al.  1995) with upper limits in the infrared (Clifton et al. 1988).  To-date no upper limits on pulsed emission in the optical have been published.  Optical observations are essential in determining the relative contributions of thermal and magnetospheric processes to pulsar emission. The continuing improvements in optical-detector sensitivities means that pulsed emission from 6 isolated neutron stars has already been detected. In this paper we present time-resolved optical observations of the central region of the supernova remnant CTB 80 and, in particular, of the two proposed pulsar candidates. ",
        "conclusions": "On the basis of luminosity, colour and timing we conclude that neither of the proposed optical counterparts are the pulsar PSR B1951+32.  We note that the slight blue extension to candidate 1 shows the expected characteristics of a hot object.  We would not expect to be sensitive to thermal radiation from the neutron star surface itself, as we estimate that a pulsar with T $\\approx 10^6$ K at the distance of CTB 80 would have a magnitude in the range $m_v$=30--31.  We may, however, be able to see emission from any hot circumstellar material.   From phenomenolgical models of pulsar optical emission (Pacini and Salvati 1983, 1987) we can derive estimated magnitudes for the emission.  For PSR 1951+32 this leads to values in the range 24--26.  However recent observations of other middle-aged pulsars, 0656+14 and Geminga (\\cite{car94}, \\cite{shear98}), have indicated a magnitude considerbly in excess of that predicted by the Pacini and Salvati model.  It is clearly very important to establish a positive optical identification for PSR B1951+32 in order to compare emission from pulsars of similar ages but with a range of other parameters such as $\\dot{P}$ and magnetic field. These observations will be crucial if we are to distinguish between the various models for high-energy emission. In particular, the pulse shape and fraction can be used to differentiate the polar cap and outer gap models.   We suggest observations with the VLT, Keck or the HST in order to definitively identify the optical counterpart of PSR B1951+32."
    },
    "9803/astro-ph9803279_arXiv.txt": {
        "abstract": "The measurement of shape parameters of sources in astronomical images is usually performed by assuming that the underlying noise is uncorrelated. Spatial noise correlation is however present in practice due to various observational effects and can affect source shape parameters. This effect is particularly important for measurements of weak gravitational lensing, for which the sought image distortions are typically of the order of only 1\\%.  We compute the effect of correlated noise on two-dimensional gaussian fits in full generality. The noise properties are naturally quantified by the noise autocorrelation function (ACF), which is easily measured in practice. We compute the resulting bias on the mean, variance and covariance of the source parameters, and the induced correlation between the shapes of neighboring sources. We show that these biases are of second order in the inverse signal-to-noise ratio of the source, and could thus be overlooked if bright stars are used to monitor systematic distortions. Radio interferometric surveys are particularly prone to this effect because of the long-range pixel correlations produced by the Fourier inversion involved in their image construction. As a concrete application, we consider the search for weak lensing by large-scale structure with the FIRST radio survey. We measure the noise ACF for a FIRST coadded field, and compute the resulting ellipticity correlation function induced by the noise.  In comparison with the weak-lensing signal expected in CDM models, the noise correlation effect is important on small angular scales, but is negligible for source separations greater than about $1^{\\prime}$. We also discuss how noise correlation can affect weak-lensing studies with optical surveys. ",
        "introduction": "The measurement of source morphologies in two-dimensional images is a fundamental problem in astronomy. The measurements of shape parameters are usually performed while assuming that the underlying noise is uncorrelated. However, spatial correlation of the noise is always present to some degree, and can significantly affect the derived parameters. In experimental situations, noise correlation can be produced by various effects such as convolution of background light with a beam (or point-spread function), interferometric imaging techniques, CCD readouts, etc.  This effect is particularly important for measurements of weak gravitational lensing. Weak lensing provides a unique opportunity to measure the gravitational potential of massive structures along the line-of-sight (for reviews see Schneider et al. 1992; Narayan \\& Bartelmann 1996). This technique is now routinely used to map the potential of clusters of galaxies (see Fort \\& Mellier 1994; Kaiser et al. 1994, for reviews). Detections of the more elusive effect of lensing by large-scale structure have been reported by Villumsen (1995) and Schneider et al. (1997) in small optical fields. The search for a strong detection of the effect on larger angular scales is currently being attempted with present and upcoming wide-field CCDs in the optical band (e.g. \\markcite{ste95}Stebbins et al. 1995; \\markcite{kai96}Kaiser 1996; \\markcite{ber97}Bernardeau et al. 1997), and with the FIRST radio survey (\\markcite{kam98}Kamionkowski et al. 1998; \\markcite{ref98}Refregier et al. 1998). The main challenge comes from the fact that weak lensing induces image distortions of only about 1\\%, and thus requires high-precision measurements of source-shape parameters. Correlated noise can produce spurious image distortions and correlations in the shapes of neighboring sources, and therefore must be carefully accounted for in weak lensing studies.  The effect of correlated noise is also particularly important for radio surveys performed with interferometric arrays. In such surveys, images are produced by Fourier inversion of the visibilities from a set of antenna pair. As a result of the incomplete visibility coverage, part of the noise on the image plane contains long range correlations which can extend across an entire field. The effect of these correlations is particularly relevant in the context of recent large radio surveys such as FIRST (\\markcite{bec95}Becker et al. 1995; \\markcite{whi96}White et al. 1996) and NVSS (Condon et al. 1997). While the present analysis was motivated by our attempt to measure weak lensing by large-scale structure with the FIRST radio survey, most of the following is general and can be applied to any wavelength and imaging technique.  \\markcite{con97}Condon (1997) computed the errors in gaussian fit parameters in the absence of noise correlation. He also gave a semi-quantitative treatment of the noise correlation which relies on simulations.  In the present paper, we generalize his approach to include a general analytic treatment of the noise correlation. We show how the spatial correlation of the noise can be naturally characterized by the noise Auto-Correlation Function (ACF). We explicitly derive the effect of noise correlation on source parameters for a general two-dimensional fit and for a two-dimensional gaussian fit. We also consider correlations induced in the parameters of nearby sources. In particular, we focus on the ellipticity correlation functions, which are used in searches for weak lensing by large-scale structure. We apply our formalism to the case of the FIRST radio survey. After measuring the noise ACF for one FIRST coadded field, we compare the noise-induced ellipticity correlations to those expected for weak lensing by large-scale structure. We show that, while they are important on small angular scales, noise correlation effects are negligible for source separations greater than about $1^{\\prime}$.  This paper is organized as follows. In \\S\\ref{noise_characterization}, we define the noise ACF and show how it can be measured in practice. In \\S\\ref{general_fit}, we consider a general two-dimensional least-square fit. We derive the bias produced by the noise correlation on the mean, variance and covariance of the fit parameters, and the correlation of the parameters for two neighboring sources. In \\S\\ref{fit_gaussian}, we apply these results to the case of two-dimensional gaussian fits and compute the error matrix. We explicitly derive the ellipticity correlation function induced by the noise correlation.  In \\S\\ref{first}, we consider the concrete case of the FIRST radio survey. We measure the noise ACF for a FIRST field, compute the induced ellipticity correlation function, and compare the latter to that expected for weak lensing. In \\S\\ref{optical}, we discuss the implications of this effect for optical surveys and related searches for weak lensing by large-scale structure. Finally, \\S\\ref{conclusion} summarizes our conclusions. ",
        "conclusions": "\\label{conclusion} We have studied the effect of noise correlation on the shape parameters in two-dimensional images. The noise correlation is conveniently described by the noise ACF which can easily be measured in practice. We derived the magnitude of the effect for a general two-dimensional least-square fit. The noise correlation can produce a bias, and affect the variance and the covariance of source parameters. In addition, it can produce systematic correlation in the parameters of pairs of sources. We find that the effect is of second order in the inverse SNR of the sources.  We applied these general results to the case of most practical interest, namely a two-dimensional gaussian fit. We computed the relevant matrices explicitly. In addition, we explicitly derived the systematic bias produced by this effect on the ellipticity correlation function of source pairs. This is particularly relevant for weak lensing studies.  As a concrete example, we studied the effect of correlated noise on the shape of sources in the FIRST radio survey. We measured the noise ACF and found long range features which extend beyond the central maximum. We computed the resulting systematic effect on the ellipticity correlation function. We find that, in the FIRST survey, noise-induced ellipticity correlations dominate over the expected weak lensing signal for $\\theta \\lesssim 0.5^{\\prime}$, but are negligible for $\\theta \\gtrsim 1^{\\prime}$.  We discussed the consequences of noise correlation for optical surveys. In optical images, noise correlation can arise from various effects such as the convolution of background light with the PSF, CCD read outs, shift-and-add preprocessing, drift-scanning, etc.  Because the effect is quadratic in the source SNR, the effect could be overlooked if systematic distortions are monitored solely with bright stars. The effect of noise correlation could thus be important for searches of galaxy-galaxy lensing and of weak lensing by large-scale structure in optical surveys."
    },
    "9803/astro-ph9803086_arXiv.txt": {
        "abstract": "A multi-domain spectral method for computing very high precision 3-D stellar models is presented. The boundary of each domain is chosen in order to coincide with a physical discontinuity (e.g. the star's surface).  In addition, a regularization procedure is introduced to deal with the infinite derivatives on the boundary that may appear in the density field when stiff equations of state are used.  Consequently all the physical fields are smooth functions on each domain and the spectral method is absolutely free of any Gibbs phenomenon, which yields to a very high precision.  The power of this method is demonstrated by direct comparison with analytical solutions such as MacLaurin spheroids and Roche ellipsoids.  The relative numerical error reveals to be of the order of $10^{-10}$. This approach has been developed for the study of relativistic inspiralling binaries. It may be applied to a wider class of astrophysical problems such as the study of relativistic rotating stars too. ",
        "introduction": "One of the most promising source of gravitational waves is the coalescence of inspiralling compact binaries. The recent development of interferometric gravitational waves detectors ({\\sl e.g. GEO600, LIGO, TAMA} and {\\sl VIRGO}) gives an important motivation for studying this problem. Such a study requires a relativistic formalism to derive the equations of motion and then an accurate and tricky method to solve the resulting system of partial differential equations. We have recently \\cite{BonGM97a} proposed a relativistic formalism able to tackle the problem of co-rotating {\\sl as well as} counter-rotating binaries system, these latter being more relevant from the astrophysical point of view. We present now a very accurate approach based on multi-domain spectral method that circumvents the Gibbs phenomenon to numerically solve this problem and which can be applied to a wide class of other astrophysical situations.  Various astrophysical applications of spectral methods have been developed in our group (for a review, see \\cite{BonGM97b}), including 3-D gravitational collapse of stellar core \\cite{BonM93}, neutron star collapse into a black hole \\cite{Gou91,GouH93,GouHG95,Nov98}, tidal disruption of a star near a massive black hole \\cite{Mar96}, rapidly rotating neutron stars \\cite{BonGSM93,SalBGH94a,SalBGH94b,HaeBS95}, magnetized neutron stars \\cite{BocBGN95,HaeB96} and their resulting gravitational radiation \\cite{BonG96}, spontaneous symmetry breaking of rapidly rotating neutron stars \\cite{BonFG96,BonFG98}, proto-neutron stars evolution \\cite{Gou97,GouHZ97,GouHZ98}.  In computational fluid dynamics, spectral methods are known for their very high accuracy \\cite{GotO77,CanHQZ88};  indeed for a ${\\cal C}^\\infty$ function, the numerical error decreases as $\\exp(-N)$ ({\\sl evanescent error}), where $N$ is the number of coefficients involved in the spectral expansion, or equivalently the number of grid points in the physical domain. This is much faster than the error decay of finite-difference methods, which behaves as $1/N^q$, with $q$ generally not larger than $3$.  For this reason, spectral methods are particularly interesting for treating 3-D problems --- such as binary configurations --- a situation in which the number of grid points is still severely limited by the capability of present and next generation computers.  Spectral methods lose much of their accuracy when non-smooth functions are treated because of the so-called {\\em Gibbs phenomenon}. This phenomenon is well known from the most familiar spectral method, namely the theory of Fourier series: the Fourier coefficients $(c_n)$ of a function $f$ which is of class ${\\cal C}^p$ but not ${\\cal C}^{p+1}$ decrease as $1/n^p$ only. In particular, if the function has some discontinuity, its approximation by a Fourier series does not converge towards $f$ at the discontinuity point: there remains a gap which is of the order 10\\%.  The multi-domain spectral method described in this paper circumvents the Gibbs phenomenon. The basic idea is to divide the space into domains chosen so that the physical discontinuities are located onto the boundaries between the domains (Sect.~\\ref{s:mapping}). The simplest example is the case of a perfect fluid star, where two domains may be distinguished:  the interior and the exterior of the star. The boundary is then simply the surface of the star. The second ingredient of the technique is a mapping between the domains defined in this way and some simple mathematical domains, which are cross products of intervals: $[a_1,a_2] \\times [b_1,b_2] \\times [c_1,c_2]$. The spectral expansion is then performed with respect to functions of the coordinates spanning these intervals (Sect.~\\ref{s:multi}). The method of resolution of a basic equation, namely the Poisson equation, is exposed in Sect.~\\ref{s:resol_Poisson}. For stiff equations of state, the above procedure is not sufficient to ensure the smoothness of all the functions.  Indeed, for a polytrope with an adiabatic index greater than 2, the density field has an infinite derivative on the surface of the star. We present in Sect.~\\ref{s:regul} a method for regularizing the density and recover the spectral precision. The power of the multi-domain spectral method is illustrated in Sect.~\\ref{s:illustr}, where comparisons are performed between numerical solutions obtained by an implementation of the method and analytical solutions (ellipsoidal configurations of incompressible fluids). Finally Sect.~\\ref{s:conclu} concludes the article by discussing the great advantages of the multi-domain spectral method for dealing with relativistic binary neutron stars. ",
        "conclusions": "\\label{s:conclu}  We have presented a new numerical approach capable of handling the surface discontinuities of stellar configurations, provided these discontinuities are star-like, which covers a wide range of astrophysically relevant situations. When used along with spectral methods this adaptive-domain technique ensures that no Gibbs phenomenon can appear. This results in a very high precision (evanescent error), as demonstrated in Sect.~\\ref{s:illustr} by comparison with exact analytical solutions. The relative error for 3-D configurations can reach $10^{-10}$ with a relatively small number of degrees of freedom ($N_r\\times N_\\theta \\times N_\\varphi = 37\\times 19\\times 18$ in each domain). Let us recall that very high accuracy is required for a lot of astrophysical problems such as numerical stability analysis. Among these problems let us mention the study of symmetry breaking of rapidly rotating stars and the determination of the orbital frequency of the last stable orbit of a neutron stars binary system.  The multi-domain spectral method is particularly well adapted to the computation of relativistic binary neutron star system. Three sets of domains can be used in this problem (see Fig.~\\ref{f:m-grille}): two sets of (three or more) domains centered on each star and a third set of (two or more) domains centered at the intersection between the rotation axis and the orbital plane. This latter set of domains which reaches spatial infinity is required to compute the gravitational field of relativistic configurations. When needed, the quantities computed on one of the three domain sets are evaluated at the collocation points of another set by means of the method presented in Sect.~\\ref{s:code}.  We are currently applying this numerical method to the computation of steady-state configurations of relativistic counter-rotating (i.e. irrotational with respect to an inertial frame) neutron star binaries, following the formulation developed in \\cite{BonGM97a}. We will report on the astrophysical results in a forthcoming paper.  An interesting  byproduct of the present technical paper is the following one.  In a previous work \\cite{BonGSM93}, we had been able to demonstrate that the virial error is representative of the true error (measured by direct comparison with analytical solutions) only in the spherically symmetric case. We had inferred that this remains valid in the axisymmetric and 3-D cases. In the present work, we have confirmed this conjecture, thanks to the ability of the present method to treat incompressible fluids, for which 3-D analytical solutions are available.    \\begin{figure} \\centerline{\\epsfig{figure=ng.eps,height=4cm}} \\caption[]{\\label{f:m-grille} Representation of the numerical domains that we use to compute relativistic steady-state configurations of binary neutron stars systems. The external domain extends to spatial infinity in order to compute the exact gravitational potentials. Due to the symmetry of the problem, only the $z > 0$ part of space is taken into account. } \\end{figure}"
    },
    "9803/astro-ph9803092_arXiv.txt": {
        "abstract": "X-ray emission from the distant lensing cluster CL2236-04 at $z$ = 0.552 was discovered by ASCA and ROSAT/HRI observations. If the spherical symmetric mass distribution model of the cluster is assumed, the lensing estimate of the cluster mass is a factor of two higher than that obtained from X-ray observations as reported for many distant clusters. However, the elliptical and clumpy lens model proposed by Kneib {\\it et al.}(1993) is surprisingly consistent with the X-ray observations assuming that the X-ray emitting hot gas is isothermal and in a hydrostatic equilibrium state. The existence of the cooling flow in the central region of  the cluster is indicated by the short central cooling time and the excess flux detected by ROSAT/HRI compared to the ASCA flux. However, it is shown that even if the AXJ2239-0429 has a cooling flow in the central region, the temperature measured by ASCA which is the mean emission-weighted cluster temperature in this case, should not be cooler than and different from the virial temperature of the cluster. Therefore, we conclude that the effect of the clumpiness and non-zero ellipticity in the mass distribution of the cluster are essential to explain the observed feature of the giant luminous arc, and there is no discrepancy between strong lensing and X-ray estimation of the mass of the cluster in this cluster. ",
        "introduction": "It is one of the prime objectives in observational cosmology to determine mass distributions and total masses of clusters of galaxies at various redshifts. It  provides a clue to constrain the unknown cosmological parameters:the species of dark matter, the nature of cosmological mass density fluctuations, the average density in the universe and the cosmological constant.\\\\ \\indent Recently, gravitational lensing has become a powerful tool to measure the mass distribution in clusters. Strong lensing events such as multiple images and/or giant luminous arcs make it possible to model the central region of clusters. However, many clusters show a mass discrepancy with a  total mass deduced from strong lensing observation being larger by a factor of 2--3 than that deduced from X-ray observations (Loeb \\& Mao 1994, Miralda-Escud\\'e \\& Babul 1995,  Kneib {\\it et al.} 1995, Schindler {\\it et al.} 1995, \\& 1997, Wu \\& Fang 1997, Ohta, Mitsuda, Fukazawa 1997). However, there are some clusters whose masses deduced in these two ways agree very well. In such cases, the multi-phase of the intra-cluster gas (Allen, Fabian \\& Kneib 1996, Allen 1997) or asymmetric mass distribution (Piere {\\it et al.} 1996, B\\\"ohringer {\\it et al.} 1997) has been taken into account in the lens modeling. This indicates that careful studies for a large number of clusters are necessary before reaching the conclusions that the lens modeling gives higher mass estimation than that from X-rays.  CL2236-04 was discovered by Melnick {\\it et al.} (1993) with a giant luminous arc. The redshift of the arc was measured to be $z_{arc}$=1.116 (Melnick {\\it et al.} 1993) confirming the gravitational lens picture very satisfactorily. The arc is rather straight and velocity structure across the arc is measured (Melnick {\\it et al.} 1993). These nature give strict constraints on the lensing mass estimation. Further, CL2236-04 is one of the best and rare clusters at high redshift of $z>0.5$ for which comparison of the mass distribution deduced by two individual methods is possible. We have performed the first deep ASCA and ROSAT/HRI pointing observations toward CL2236-04. In this paper, we report the discovery and the properties of X-ray emission from CL2236-04 and the mass distribution deduced from the X-ray observation is compared to the lensing mass. ",
        "conclusions": "A  new X-ray source in the direction of the distant lensing cluster, CL2236-04, is discovered by ASCA, named AXJ2239-0429. The source extent is resolved by the following ROSAT/HRI observation and it confirms that the X-ray source is the cluster of galaxies, CL2236-04.  The mass distribution of the cluster derived by  X-ray observations and by the observed nature of the giant luminous arc is compared. It is found that within the frame work of  a spherical isothermal $\\beta$ model, the required mass  to explain the observed location of the giant luminous arc is larger than that expected from the X-ray observations as reported for many giant luminous arc clusters (Wu \\& Fang 1997). An introduction of  a large cosmological constant reduces the lensing mass but the total reduction  is only 20-30\\%. Therefore, the cosmological constant can not be the main solution for the discrepancy found in many clusters.  Cooling flow can be one of the other possible solutions (Allen, Fabian \\& Kneib 1996, Allen 1997). According to the cooling flow model, the hot gas in the cluster central region is in the multi-temperature phase. The temperature of the hottest phase gas which represents a real gravitational potential depth, obtained by the multi-phase model fitting becomes much higher than the temperature obtained by the single phase model fitting because the last one is the average temperature of the multi-temperature gas. Allen (1997) is claiming that the most of the discrepancy reported before can be explained by the existence of the cooling flow. There are several evidences of the existence of the cooling flow in the central region of  AXJ2239-0429 as shown in Sec.2. It may be possible that the measured temperature of AXJ2239-0429 underestimates the cluster potential depth. If the ASCA temperature underestimates the cluster potential depth by a factor of two due to the cooling flow, the Einstein ring radius of the cluster for a source at $z=1.116$ becomes as large as the distance  between the giant luminous arc and the cluster center. However, the straight nature of the arc and a lack of the bright counter image can not be explained by a spherically symmetric lens. Further, the perturbation by the gravitational potential of galaxy A is not negligible. The absolute luminosity of galaxy A in visual band is $3.3\\times 10^{11}h_{50}^{-2}L_{\\odot}$ (Melnick et al. 1993).   This luminosity is translated into the the line of sight velocity dispersion of $\\sim 370{\\rm km/s}$ by using the Faber-Jackson law (Binney \\& Tremaine 1987). If the potential depth is about a factor of two deeper than that inferred from the ASCA temperature, it is impossible to construct the model which is able to explain the location and the observed feature of the giant luminous arc due to the large perturbation of the galaxy A. Therefore, we conclude that even if the AXJ2239-0429 has a cooling flow in the central region, the temperature measured by ASCA which is the mean emission-weighted cluster temperature in this case, should not be cooler than and different from the virial temperature of the cluster.  The most plausible solution is effect of the clumpiness of the gravitational potential due to galaxy A and non-zero ellipticity in the mass distribution of the cluster. KMG  proposed the cluster mass distribution model by taking into account these effects which explain the observed features of the giant luminous arc. The ellipticity, position angle and velocity dispersion of the  cluster predicted by KMG is surprisingly consistent with the results of the X-ray observations of the cluster as shown in the previous section. The velocity dispersion of galaxy A assumed by KMG is also consistent with that deduced by the Faber-Jackson law. Therefore, we conclude that the effect of the clumpiness of the gravitational potential due to galaxy A and non-zero ellipticity in the mass distribution of the cluster are essential to explain the observed feature of the giant luminous arc and the cluster mass distribution required by the strong lensing effect is totally consistent with the X-ray observations assuming that the X-ray emitting hot gas is isothermal and in a hydrostatic equilibrium state."
    },
    "9803/astro-ph9803167_arXiv.txt": {
        "abstract": "We present Washington system $C,T_1$ color-magnitude diagrams of 13 star clusters and their surrounding fields which lie in the outer parts of the LMC disk ($r>4\\arcdeg$), as well as a comparison inner cluster. The total area covered is large ($2/3^{\\Box^\\circ}$), allowing us to study the clusters and their fields individually and in the context of the entire galaxy. Ages  are determined by means of the magnitude difference $\\delta$T$_1$ between the giant branch clump and the turnoff, while metallicities are derived from the location of the giant and subgiant branches as compared to fiducial star  clusters. This yields a unique dataset in which ages and metallicities for both a significant sample of clusters and their fields are determined homogeneously. We find that in most cases the stellar population of each star cluster is quite similar to that of the field where it is embedded, thus sharing its mean  age and metallicity. The old population (t$\\geq$10 Gyr) is detected in most fields as a small  concentration of stars on the horizontal branch blueward and faintward of the prominent clump. Three particular fields present remarkable properties: (i) The so far unique cluster ESO121-SC03 at $\\approx$9 Gyr has a  surrounding field which shares  the same properties, which in turn is also unique in the sense that such a dominant old field component  is not present elsewhere, at least  not significantly in the fields as yet  studied. (ii) The field surrounding the far eastern intermediate age cluster OHSC\\,37   is noteworthy in the sense that we do not detect any evidence of LMC stars: it is essentially a Galactic foreground field. We can thus detect the LMC field out to $>11\\arcdeg$ (the deprojected distance of ESO121SC03), or $\\sim 11$ kpc, but not to $13\\arcdeg (\\sim 13$ kpc), despite the presence of \\cls at this distance. (iii) In the northern part of the LMC disk the fields of SL388 and SL509 present color-magnitude diagrams with a secondary  clump  $\\approx$0.45 mag fainter than the dominant intermediate age clump, suggesting a stellar population component    located behind the LMC disk at a distance comparable to that of the SMC. Possibly we are witnessing a depth effect in the LMC, and the size of the corresponding structure is comparable to the size of a dwarf galaxy.  The unusual spatial location of the cluster OHSC37 and the anomalous properties of the SL\\,388 and SL\\,509 fields  might be explained as  debris  from   previous LMC interactions with  the Galaxy and/or the SMC.  The mean    metallicity derived for the intermediate age outer disk clusters is $<$[Fe/H]$>$$=-0.7$ and for their surrounding fields $<$[Fe/H]$>$$=-0.6$. These values are significantly lower than found by Olszewski {\\it et al.} (1991, AJ, 101, 515) for a sample of clusters of similar age, but are in good agreement with several recent studies. A few clusters stand out in the age--metallicity relation in the sense that they are intermediate age clusters at relatively low \\met ([Fe/H]$\\approx -1$). ",
        "introduction": "Unveiling the   star formation  history and  chemical enrichment of galaxies is critical for understanding how they form and evolve. In this respect, Local Group galaxies play a fundamental r\\^ole (Hodge 1989). The  proximity of the Large Magellanic Cloud allows one to probe its stellar population properties with different techniques, namely photometry and spectroscopy of individual stars in clusters and the field, and integrated methods in the case of  star clusters. Such studies for nearby galaxies are important to better understand very distant galaxies, whose stellar populations can only be probed   by means of integrated properties.  Concerning field color-magnitude diagram (CMD) studies, ground-based observations have allowed the accurate study of the brighter evolutionary sequences, sampling relatively large fields throughout the LMC bar and disk (e.g. Butcher 1977, Hardy  {\\it et al.} 1984, Bertelli {\\it et al.} 1992, Westerlund, Linde \\&  Lyng\\aa\\ 1995, Vallenari {\\it et al.} 1996). A prominent intermediate age (1-3 Gyr) stellar population is universally present in  these fields, together with varying amounts of  young blue main sequence (MS) stars.   Hubble Space Telescope (HST) observations are limited to a small viewing area,  but in turn allow  deeper photometry, to well below the old MS turnoff. Several such fields have been studied with  the V and I bands. One, at $\\approx$4$^o$ north of the bar (Gallagher {\\it et al.} 1996) and another near the SE end of the bar (Elson, Gilmore \\& Santiago 1997) show evidence for the major star-forming event to have ocurred $\\approx$2 Gyr ago, in agreement with the ground-based studies. More recent HST studies (Holtzman {\\it et al.} 1997, Geha {\\it et al.} 1998) have investigated three LMC fields at 3$^o$ to 4$^o$ from the bar center, one located in the  north-east and two in the north-west. Surprisingly, they are finding many more faint MS stars than expected and suggest that there has been more star formation in the past than previously believed. Their models, assuming a standard IMF slope, suggest that fully one half of the stars in these fields were formed more than 4 Gyr ago. Although the latter studies refer to their fields as ``outer\", we point out that in the present work we deal with genuine outer disk fields, well beyond any of these HST studies.  Integrated photometry of large star cluster samples of all ages have shown differences in the spatial distribution of age groups  both in the bar region (Bica, Clari\\'a \\& Dottori 1992) and the entire LMC (Bica {\\it et al.} 1996). Differences in the  spatial distribution among  young groups have provided insight  on the formation process  and subsequent dynamical evolution of star cluster generations (Dottori {\\it et al.} 1996). This integrated photometry cluster sample has been compared  with integrated star cluster color models and has provided constraints on the cluster formation history (Girardi \\& Bica 1993, Girardi {\\it et al.} 1995).  CMDs of LMC star clusters have also revealed a large intermediate age population  (1-3 Gyr), which is separated  by a pronounced age gap  from the old stellar population as denoted by  a few genuine globular clusters (see Da Costa 1991, Suntzeff {\\it et al.} 1992, Olszewski, Suntzeff \\& Mateo 1996 for reviews). Recently, CMDs in the  Washington system of a sample of candidate old clusters selected from the Bica {\\it et al.} (1996) and Olszewski {\\it et al.} (1991) studies revealed them to instead be of intermediate age  (Geisler {\\it et al.} 1997, hereafter Paper I). This study increased  considerably the known sample of 1-3 Gyr old clusters with accurate age determinations,  and reinforced the conclusion that a major formation epoch was preceded by a quiescent  period of many Gyr, or that cluster dissipation has been more effective than generally believed (e.g. Olszewski 1993).  The objective of the present paper is to compare the properties of outer LMC clusters with those of their surrounding fields  by using  the same observational technique,  and to infer  the age-metallicity relation, and whether it depends  on the  spatial distribution throughout the LMC disk. In order to achieve this we employ Washington system  C, T$_1$ bands and construct CMDs using the data of Paper I. Ages are inferred from the difference $\\delta$T$_1$ between the giant branch clump and the turnoff (see also Paper I), and from  the occurrence  of particular stellar evolutionary sequences in the CMDs.   The giant and subgiant branches allow one to derive metallicities using a technique analogous to that of Da Costa and Armandroff (1990) for VI photometry. However, our combination of Washington system filters is three times   more \\met sensitive than the VI system (Geisler \\& Sarajedini 1996, 1998), allowing us to obtain accurate \\mets for both the clusters and their fields. The cluster/field sample and the observations are described in Section 2. The  cluster and field CMDs are described in Section 3. Ages and metallicities are derived in Section 4. In Section 5 we discuss  the chemical enrichment of the outer disk, and in Section 6 the presence of dual clumps in two fields is noted and the possibility of a depth effect in this portion of the LMC disk is discussed. Finally, the conclusions of this work are given in Section 7. ",
        "conclusions": "By using Washington system color-magnitude diagrams of a large sample of IACs and their surrounding fields in the LMC outer disk we were able to derive ages and metallicities. The cluster stellar population is a major component of the field where it is embedded, thus sharing its mean  age and metallicity properties, except in three  cases where the clusters are     $\\approx$0.3 dex more metal poor than the field, suggesting that the chemical enrichment was not globally homogeneous in the LMC. Our data is consistent with a scenario in which local   star formation events generated both the clusters and a significant part of their surrounding stellar fields. Thus  during the last 1-3 Gyr  the dynamical evolution of the disk has not significantly taken them apart, which provides information on the  diffusion time-scale  and mixing of stellar generations in the disk. The old population (t$\\geq$10 Gyr) is detected in most fields as a small  concentration of stars on the horizontal branch blueward and faintward of the prominent clump.  The unique cluster ESO121-SC03 at $\\approx$9 Gyr has a  surrounding field which shares  the same stellar \\pop. No other field so far studied is dominated by such an old population. One possibility would be that a field and cluster population coupling might last that long. Alternatively, we might be dealing with a building block recently accreted by the LMC in the form of a dwarf galaxy.  One IAC cluster (OHSC37) is so far from the LMC body that no surrounding LMC field is detected.  The present observations suggest that the LMC stellar disk extends out to between $\\sim 11-13$ kpc in deprojected radius. In the northern part of the LMC outer disk the fields of SL388 and SL509 present evidence of a depth effect with a secondary component    located behind the LMC disk at a distance comparable to that of the SMC. A background layer of stars  in the LMC was possibly detected, and its size is at least $\\approx$1 kpc, comparable to that of a dwarf galaxy.  The peculiar location of the cluster  OHSC37 and the depth effect in the SL\\,388 and SL\\,509 fields  might be explained as  debris  from   previous interactions of the LMC with  the Galaxy and/or the SMC.  The average metallicity derived for the present outer disk IACs is $<$[Fe/H]$>$$=-0.71\\pm0.17$ and for their surrounding fields $<$[Fe/H]$>$$=-0.61\\pm0.11$. A few clusters stand out in the age--metallicity relation in the sense that they are intermediate age clusters at [Fe/H]$\\approx -1$. The outer LMC disk (clusters and fields), at least of this age range (1-3 Gyrs) seem to be  more metal-poor than previously generally regarded.  The authors would like to thank Cristina Torres for providing data critical to the \\met calibration in advance of publication. E. Geisler, as always, was an inspiration. This research is supported in part by NASA through grant No. GO-06810.01-95A (to DG) from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. This work was partially supported by the Brazilian institutions CNPq and FINEP, the Argentine institutions CONICET and CONICOR, and the VITAE and Antorcha foundations."
    },
    "9803/astro-ph9803021_arXiv.txt": {
        "abstract": "The {\\it COBE} FIRAS data contain foreground emission from interplanetary, Galactic interstellar dust and extragalactic background emission. We use three different methods to separate the various emission components, and derive the spectrum of the extragalactic Far InfraRed Background (FIRB). Each method relies on a different set of assumptions, which affect the FIRB spectrum in different ways. Despite this, the FIRB spectra derived by these different methods are remarkably similar. The average spectrum that we derive in the $\\nu = 5 - 80~$cm$^{-1}$ (2000 -125 \\um) frequency interval is: $I(\\nu) = (1.3\\pm0.4)\\times10^{-5}\\ (\\nu /\\nu_0)^{0.64\\pm.12}\\ P_{\\nu}(18.5\\pm 1.2~$K), where $\\nu_0=100~$cm$^{-1}$ ($\\lambda_0=100\\ \\um$) and $P$ is the Planck function. The derived FIRB spectrum is consistent with the DIRBE 140 and 240 \\um\\ detections. The total intensity received in the 5 - 80~cm$^{-1}$ frequency interval is 14 nW m$^{-2}$ sr$^{-1}$, and comprises about 20\\% of the total intensity expected from the energy release from nucleosynthesis throughout the history of the universe. ",
        "introduction": "The FIRAS (Far~Infrared~Absolute~Spectrophotometer) instrument aboard the Cosmic Background Explorer ({\\it COBE}) satellite (Boggess \\etal\\ 1992 and references therein; Mather, Fixsen, \\& Shafer 1993) was designed to measure the spectrum of the Cosmic Microwave Background (CMB), the Galaxy, and the Far InfraRed Background (FIRB). The FIRAS instrument and calibration are discussed by Mather, Fixsen \\& Shafer (1993) and Fixsen \\etal\\ (1994b). The FIRAS observations of the CMB are discussed by Mather \\etal\\ (1990, 1994), and Fixsen \\etal\\ (1994a, 1996, 1997a). The FIRAS observations of the Galaxy are discussed in Wright \\etal\\ (1991), Bennett \\etal\\ (1994), and Reach \\etal\\ (1995). Puget \\etal\\ (1996) used the Pass 3 FIRAS observations to make a tentative background determination simular to one of the methods used here.  At frequencies below $\\sim$ 20 cm$^{-1}$ (500 \\um), the FIRB is overwhelmed by the CMB. However, the CMB can be subtracted from the data since its emission is spatially uniform, and it has a Planck spectrum. The dipole anisotropy of the CMB must be included but the other anisotropy is not large enough to be a problem. At frequencies above $\\sim$ 100 cm$^{-1}$ (100 \\um),beyond the FIRAS range, the observed spectrum is dominated by the zodiacal emission, which is much more complex, since its intensity and spectrum are spatially varying. In between these frequencies, the observed spectrum is dominated by Galactic emission, which can be identified by its distinct (from zodiacal) spectrum, and its spatial variation over the sky. After the subtraction of these foreground emissions, any isotropic residual emission may still contain a uniform component from the solar system or the Galaxy, which must be estimated and separated from the FIRB. In this paper we examine three methods, which give a consistent result for the FIRB. ",
        "conclusions": "The three methods just described have different assumptions, weaknesses and strengths. The first method assumes that the Galactic spectrum is fixed, and that only its intensity varies with position. Temperature variations are observed to occur on various scales (Reach et al. 1995; FIRAS Exp Sup 1997; Lagache et al. 1998). Based on the DIRBE analysis, we expect these variations to have a only a small effect on our results. The DIRBE high latitude data at 240 \\um\\ were analyzed using both a single- and a variable-spectrum model to describe the Galactic foreground emission (Arendt et al. 1998), with no significant difference in the derived background.  The second method does not assume a fixed Galactic spectrum, but almost all of the power in the fit is in a single component indicating that there is little systematic correlation of temperature with intensity of the H~I 21 cm line. This method solves directly for the dust emission from the dominant neutral and ionized phases of the ISM. It therefore avoids a potential problem of the third method (which relies on the DIRBE analysis) using backgrounds derived from the IR-to-H~I correlation only. Two problems with this method are, that at the FIRAS resolution, the correlation between the IR and the H~I line intensity is non-linear and that it ignores the molecular regions. The first was treated by using a quadratic fit to the relation. This choice is not unique, but it is supported by observations of the IR to H~I correlation in the Galaxy (Dall'Oglio et al. 1985; Reach, Koo, \\& Heiles 1994). Molecular clouds are less abundant at the high Galactic latitudes that are used here.  The third method relies on the accuracy of the DIRBE determination of the FIRB. The main uncertainty in the background determination are those associated with the determination of the foreground emission, and they are discussed in detail by Hauser et al. (1998). The use of a DIRBE template for the FIRAS background determination introduces another uncertainty, namely the consistency between the FIRAS and DIRBE calibration. Fixsen et al. (1997b) show that the most significant difference in the calibration is in the 240 \\um\\ DIRBE band, but using the FIRAS instead of the DIRBE calibration for that band introduces a very small change in  the background, from a value of 13.6 to 12.7 nW m$^{-2}$ sr$^{-1}$ (Hauser et al. 1998).  Overall, the three methods yield a consistent spectrum for the FIRB, an encouraging result considering the substantial differences in the three approaches (see fig 4a). The weaknesses of each method are compensated by the other methods. The color method assumes a single spectrum, but the other methods allow for color variation in the Galactic foreground. The line emission method uses a quadratic fit over 25\\% of the sky, but the other methods use linear fits which are insensitive to the fraction of sky.  The DIRBE method only uses the correlation between H~I and the foreground in small regions, but the line emission method includes ionized gas and the color method makes no assumptions about the gas.  In all three FIRB spectra the background peaks at $\\sim 50~$cm$^{-1}$ ($\\sim 200$ \\um), and exhibits a definite turnover at the higher frequencies. The higher frequencies are more affected by both noise and systematic effects. The average of the three spectra can be fit by:  \\begin{equation} I_\\nu=(1.3\\pm 0.4)\\times 10^{-5}(\\nu/\\nu_\\circ)^{.64\\pm.12}\\ P_\\nu(18.5\\pm1.2~K), \\end{equation} in the 5 to 80 cm$^{-1}$ frequency range ($\\lambda$ between 125 and 2000 \\um) where $\\nu_\\circ=100~$cm$^{-1}$, and $P_\\nu$ is the familiar Planck function. The uncertainties are highly correlated, with correlations of .98 for the intensity and index, $-$.99 for the intensity and temperature. and $-$.95 for the index and temperature.  Figure 4b compares the analytical fit to the FIRB derived here, to the tentative background derived by Puget et al. (1996). The spectra differ significantly at $\\nu \\sim 35 - 60~$cm$^{-1}$, probably the result of the difference in subtraction of dust emission related to H$^+$. The crosses in the Figure represent the nominal DIRBE detections at 140 and 240 \\um, while the diamonds represent the DIRBE detections using the FIRAS calibration. While the effect of the FIRAS calibration is larger at 140 \\um, the uncertainty in the calibration is larger at this wavelength. So, the FIRB derived here is consistent with that derived by the DIRBE, within the uncertainty of the DIRBE$-$FIRAS calibration. The range of values for the FIRB at 100 \\um\\ represent the upper limit derived by Kashlinsky, Mather, \\& Odenwald (1996) from a fluctuation analysis of the 100 \\um\\ DIRBE maps and the lower limit is derived by Dwek et al. (1998) from the 140 and 240 \\um\\ DIRBE detections.  Dwek \\etal\\ (1998) show that the uniform DIRBE residuals cannot be produced by any local (solar system or Galactic) emission sources. Hence, the uniform residual derived here is most likely of extragalactic origin. The total flux received in this wavelength region is 14 nW m$^{-2}$ sr$^{-1}$, or about 20\\% of the total expected flux of about 70 nW m$^{-2}$sr$^{-1}$ associated with the production of metals throughout the history of the universe in some models (Dwek \\etal\\ 1998)."
    },
    "9803/astro-ph9803217_arXiv.txt": {
        "abstract": "We study numerically the evolution of an adiabatic relativistic fireball expanding into a cold uniform medium. We follow the stages of initial free expansion and acceleration, coasting and then deceleration and slowing down to a non-relativistic velocity. We compare the numerical results with simplified analytical estimates. We show that the relativistic self similar Blandford-McKee solution describes well the relativistic deceleration epoch. It is an excellent approximation throughout the relativistic deceleration stage, down to $\\gamma \\sim 5$, and a reasonable approximation even down to $\\gamma \\sim 2$ though the solution is rigorous only for $\\gamma \\gg 1$. We examine the transition into the Blandford-McKee solution, and the transition from the solution to the non-relativistic self similar Sedov-Taylor solution. These simulations demonstrate the attractive nature of the Blandford-McKee solution and its stability to radial perturbations. ",
        "introduction": "Sedov (1946), Taylor (1950) and Von Neumann (1947) discovered, in the forties, a self similar solution of the strong explosion problem, in which a large amount of energy is released on a short time scale in a small volume. This solution is known today as the ``Sedov-Taylor'' self similar solution.  It describes a shock wave propagating into a uniform density surrounding. The shock wave and the matter behind it decelerate as more and more mass is collected. This solution describes well the adiabatic stage of a supernova remnant evolution.  Blandford and McKee (1977) have later established a self similar solution describing the extreme relativistic version of the strong explosion problem. In this solution the Lorentz factor of the shock and the fluid behind it is much larger than unity. Such high Lorentz factors arise if the rest mass contained within the region where the energy $E$ was released is much smaller than $E$. In other words when the region containing the energy is ``radiation dominated'' rather than matter dominated. Such a region was later termed a ``fireball''.  Cavallo and Rees (1978) have considered the physical processes relevant in a radiation dominated fireball as a model for gamma-ray bursts (GRBs). Goodman (1986) and Paczy\\'nski (1986) have then considered the evolution of such a fireball. They have shown that a initially the radiation-pair plasma in a purely radiative fireball behaves like a fluid and it expands and accelerated under its own pressure until the local temperature drops to $\\sim 20$keV, when the last pairs annihilate and the fireball becomes optically thin.  Later Shemi and Piran (1990) considered matter contaminated fireballs. They have shown that, quite generally, all the initial energy will be transfered to the baryons in such fireballs whose final outcome is a shell of relativistic freely expanding baryons.  Piran, Shemi and Narayan (1993) and M\\'esaz\\'aros Laguna and Rees (1993) have later carried these calculations in greater details. These works show that an initially homogeneous fireball will first accelerate while expanding and then coast freely as all its internal energy was transformed to kinetic energy.   The surrounding matter will eventually influence the fireball after enough external matter has been collected and most of the energy has been transfered from the shell to the ISM.  If the surrounding matter is diluted enough, then this will take place only after the initial free acceleration phase.  This influence was considered by M\\'esaz\\'aros and Rees (1992), and Katz (1993), who suggested that the GRB is produced during this stage. The detailed shock evolution was later studies by Sari and Piran (1995). It is only after these stages, that the fireball have given the ISM most of its energy and then the self similar deceleration solution of Blandford-McKee applies. When the shock decelerates enough so that it is no longer relativistic, it is described by the Sedov-Taylor solution.  Today it is widely accepted that GRBs involve relativistic expanding matter of this kind. While the GRB itself is produced, most likely, via internal shocks (Narayan, Paczy\\'nski \\& Piran, 1992; Rees \\& M\\'esaz\\'aros, 1994; Sari \\& Piran, 1997).  The observed GRB afterglow corresponds, on the other hand to the slowing down of this relativistic flow. This has led to an increasing interest in the fireball solution and in its various regimes. In this paper we study the evolution of a homogeneous fireball focusing on its interaction with the surrounding matter.  We do not consider here internal shocks, which arise due to interaction within the relativistic flow and require nonuniform velocity.  We have developed a spherically symmetric relativistic Lagrangian code based on a second order Gudnov method with an exact Riemann solver to solve the ultra-relativistic hydrodynamics problem. With this code, it is possible to track the full hydrodynamical evolution of the fireball within a single computation, from its initial ``radiation dominated'' stage at rest through its acceleration, coasting, and shock formation, up to the relativistic deceleration and finally to the Newtonian deceleration.  With typical parameters this computation spans more than eight orders of magnitudes in the size of the fireball and more than twenty orders of magnitudes in its density. We describe these computations here.  We show that, quite generically, the solution converges during the relativistic deceleration phase to the Blandford-McKee solution and then it transforms to the Sedov-Taylor solution. Even though the attractive nature of the Blandford-McKee solution suggests that it is stable we explore explicitly the stability of this solution and we show that it is stable to radial perturbations.  We review, first, in section \\ref{sec_Analytic} the current analytic understanding of the fireball evolution thorough the following stages: (i) free acceleration and coasting, (ii) energy transfer (iii) relativistic self similar solution and (iv) Newtonian Sedov-Taylor solution.  We discuss the numerical results in section \\ref{numerical}.  In section \\ref{perturbation} we examine the evolution of perturbations to the Blandford-McKee solution. We discuss the implications of these results in section \\ref{conclusions}. ",
        "conclusions": ""
    },
    "9803/astro-ph9803117_arXiv.txt": {
        "abstract": "We investigate the formation of virialized halos due to the gravitational collapse of collisionless matter using high-resolution N-body simulations. A variety of formation scenarios are studied, ranging from hierarchical clustering to monolithic radial collapse. The goal of these experiments was to study departures from the universal density profiles recently found to arise in cosmological settings. However, we found that even for models which exhibit quite a different formation history, the density and velocity dispersion profiles of the virialized halos are strikingly similar.  Power law density profiles do not result even in models with initial power law profiles and no initial substructure or non-radial motions. Such initial conditions give rise to a radial orbit instability which leads to curved velocity dispersion and density profiles. The shapes of the density profiles in all our models are well parameterized by the profiles of halos formed  in a generic cosmological setting. Our results show that the universality of dark halo density profiles does not depend crucially on hierarchical merging as has been suggested recently in the literature. Rather it arises because apparently different collapse histories produces a near universal angular momentum distribution among the halo particles. We conclude that the density and velocity dispersion profiles of virialized halos in an expanding universe are robust outcomes of gravitational collapse, nearly independent of the initial conditions and the formation history. ",
        "introduction": "The hierarchical collapse of dark matter into virialized halos is likely to have played a key role in the formation of galaxies and clusters of galaxies. Several aspects of this process have been studied in recent years, in particular the role of initial conditions and of ongoing mergers in shaping the final structure of the dark matter halos. However, it remains an open question whether violent relaxation erases all information about the initial conditions and the formation history of halos. If it does not, then the density  and velocity profiles of dark matter halos may bear imprints of the initial power spectrum as well as the cosmological density parameter $\\Omega$  and the cosmological constant $\\Lambda$.  Such a dependence between the initial power spectrum and the density profiles of virialized objects was pointed out by Hoffman \\& Shaham (1985) and Hoffman (1988). Their conclusions were based on the secondary infall model of Gunn \\& Gott (1972), and the self-similar solution of Fillmore \\& Goldreich (1984) and Bertschinger (1985).  These analytic calculations, however, were based on simplifying assumptions (\\eg spherical symmetry) and, therefore, could not consider all aspects of gravitational collapse. N-body simulations, which begin with generic initial conditions, provide an attractive alternative. Though the effects of gas dynamics are neglected, they are likely to represent a realistic description of the formation of galaxy clusters and the outer parts of galactic halos.  In recent years, multi-mass simulation techniques have been developed which allow one to investigate the formation of individual halos with high numerical resolution. Navarro, Frenk \\& White (1996, 1997) (hereafter NFW) have investigated the structure of dark matter halos which form in a cold dark matter (CDM) universe. They found that the density profiles of halos do not follow a power law, but tend to have a slope $\\alpha = d \\ln \\varrho/d \\ln r$ with $\\alpha=-1$ near the halo center  and $\\alpha = -3$ at large radii. Over more the four orders of magnitude in mass, the density profiles follow a universal law, which can be parameterized by \\begin{equation} \\frac{\\varrho(r)}{\\varrho_b} = \\frac{\\delta_n}{\\frac{r}{a_n} (1 + \\frac{r}{a_n})^2}. \\label{nfw} \\end{equation} The two parameters are the scale radius $a_n$ which defines the scale where the profile shape changes from slope $\\alpha >-2$ to $\\alpha <-2$, and the characteristic overdensity $\\delta_n$. They are related because the mean overdensity enclosed within the virial radius $r_{vir}$ is $180$. Equation~(\\ref{nfw}) differs slightly in its asymptotic  behavior at large radii from the profile which was proposed by Hernquist (1990) to describe the mass profiles of elliptical galaxies.  This profile, which goes like $r^{-4}$ at large radii, was used by Dubinsky \\& Carlberg (1991) to fit the density distribution of halos which were formed in their CDM-like simulation.  Cole \\& Lacey (1996), Tormen, Bouchet \\& White (1996), Navarro, Frenk \\& White (1997), Huss, Jain \\& Steinmetz 1997 (hereafter HJS) and Moore \\etal (1997), among others, have extended the original results of NFW to other initial spectra, to other cosmologies  and to higher spatial and mass resolution. Most of the above studies show that halos which form in a variety of cosmological models are well described by equation~(\\ref{nfw}). The scale radius and the central overdensity appear to be related to the formation time of the halo (Navarro \\etal 1997). These results suggest that the density profile found by NFW is quite generic for any scenario in which structures form due to hierarchical clustering. The power spectrum and cosmological parameters only enter by determining the typical formation epoch of a halo of a given mass and thereby the dependence of the characteristic radius $a$ on the total mass of the halo.  All these studies were based on initial particle distributions drawn from a Gaussian random field. Since cold dark matter scenarios have initial perturbations on all scales, the formation of halos generically proceeds by hierarchical merging of dark matter clumps. In this paper we instead use a set of artificial initial conditions, gradually increasing in complexity, to isolate the physical effects that determine the final properties of the halo. We control the degree of substructure by varying the velocity dispersion of the particles. The merging history of the halos varies between the two extremes of monolithic radial collapse and hierarchical merging. ",
        "conclusions": "We have simulated the collapse of dark matter halos with varying initial conditions. As shown in figure 1 the models are designed to vary the amount of sub-structure and merging in the collapsing halo. We found that even with a smooth initial particle distribution and minimal merging, the virialized density and velocity dispersion profiles of the halos are very similar to the halos formed in a cold dark matter cosmology. The density profile has logarithmic slope $-2$ at a characteristic radius and a slope shallower in the inner regions and steeper in the outer regions. It is well fit by the form proposed by NFW.  The only way we could get a power law density profile was to turn off the non-radial components of the gravitational force completely. Since this artificial condition  can only be satisfied in a simulation, we conclude that dark matter halos resulting from gravitational collapse do not have power law profiles. Instead the density and velocity dispersion profiles have a nearly universal characteristic shape, independent of the initial conditions or the formation history.  Our results suggest that the key to universal density profiles does not lie in the hierarchical merging history of cold dark matter halos, as recently argued by Syer \\& White (1996). It is a more generic feature of gravitational collapse, which can be derived from the near universal behavior of the velocity field of the halo particles. The velocity dispersion profile grows with time in nearly the same way for halos in all the models. The mechanism by which angular momentum is acquired can vary, e.g. model II has little merging and instead shows a strong bar instability, but the resulting velocity dispersion profiles have similar shape.  NFW used hierarchical merging as embodied in the Press-Schechter scenario to model the variation of the scale radius $a_n$ with halo mass and the cosmological model. To the extent that their results are empirically verified, they provide a useful guide to the variations of halo profiles. But our results suggest that the merging history is not a fundamental determinant of halo profiles, since the emergence of the characteristic scale radius occurs even in models with negligible merging. It would be interesting to pursue this issue further with detailed dynamical studies."
    },
    "9803/astro-ph9803267_arXiv.txt": {
        "abstract": "We report the discovery of a quasi-periodic oscillation (QPO) at 61.0$\\pm$1.7 Hz with the Rossi X-Ray Timing Explorer in the low-mass X-ray binary and persistently bright atoll source GX 13+1 (4U 1811--17).  The QPO had an rms amplitude of 1.7$\\pm$0.2\\% (2--13.0 keV) and a FWHM of 15.9$\\pm$4.2 Hz. Its frequency increased with count rate and its amplitude increased with photon energy.  In addition a peaked noise component was found with a cut-off frequency around 2 Hz, a power law index of around --4, and an rms amplitude of $\\sim$1.8\\%, probably the well known atoll source high frequency noise. It was only found when the QPO was detected. Very low frequency noise was present with a power law index of $\\sim$1, and an rms amplitude of $\\sim$4\\%. A second observation showed similar variability components.  In the X-ray color-color diagram the source did not trace out the usual banana branch, but showed a two branched structure.  This is the first detection of a QPO in one of the four persistently bright atoll sources in the galactic bulge. We argue that the QPO properties indicate that it is the same phenomenon as the horizontal branch oscillations (HBO) in Z sources. That HBO might turn up in the persistently bright atoll sources was previously suggested on the basis of the magnetospheric beat frequency model for HBO. We discuss the properties of the new phenomenon within the framework of this model. ",
        "introduction": "Based on their correlated X-ray timing and spectral behavior the brightest low-mass X-ray binaries (LMXBs) can be divided into two groups; the atoll sources and Z sources (Hasinger \\& van der Klis \\markcite{hk89}1989, hereafter HK89; van der Klis \\markcite{kl95}1995). Atoll sources show two states: the island and the banana state, after the tracks they produce in an X-ray color-color diagram (CD). Z sources on the other hand trace out a Z-like shape in a CD, with usually three branches: the horizontal, the normal, and the flaring branch.  The power spectra of atoll sources can be described by two noise components plus, sometimes, a Lorentzian component to describe quasi-periodic oscillations (QPOs). The first noise component, the very low frequency noise (VLFN) has a power law shape $P\\propto\\nu^{-\\alpha}$, with $1\\leq\\alpha\\leq1.5$. The other, the high frequency noise (HFN), can be described by a power law with an exponential cut-off $P\\propto\\nu^{-\\alpha} e^{-\\nu/\\nu_{cut}}$, usually with $0\\leq\\alpha\\leq0.8$ and $0.3\\leq\\nu_{cut}\\leq25$ Hz. The HFN sometimes has a local maximum (``peaked noise'') around 10-20 Hz (see van der Klis 1995). Yoshida et al. (1993) found peaked noise around 2 Hz. Several broad QPO(-like) peaks were found with the Rossi X-ray Timing Explorer (RXTE) around 20 Hz (Strohmayer et al. \\markcite{st96}1996; Ford et al. \\markcite{fo97}1997; Yu et al. \\markcite{yu97}1997; Wijnands \\& van der Klis \\markcite{wikl97}1997), and one at 67 Hz (simultaneously with one at 20 Hz, see Wijnands et al. \\markcite{wi97b}1997b;). In addition to QPOs below 100 Hz, QPOs between 300 and 1200 Hz, the so-called kHz QPOs, have been found.  The power spectra of Z sources show three broad noise components: VLFN with $1.5\\leq\\alpha\\leq2$, HFN with $\\alpha\\sim0$ and $30\\leq\\nu_{cut}\\leq100$ Hz, and low frequency noise (LFN), which has the same functional shape as the HFN with $\\alpha\\sim0$ and $2\\leq\\nu_{cut}\\leq20$ Hz. Note that despite having the same name, HFN in Z sources is not the same phenomenon as HFN in atoll sources. Z sources show three types of QPOs: the normal/flaring branch QPO (N/FBO) with centroid frequencies from 6 to 20 Hz, the horizontal branch QPO (HBO) from 15 to 60 Hz, and the kHz QPOs in the same range as observed in atoll sources. (For an extensive review on the power spectra of atoll and Z sources we refer to van der Klis \\markcite{kl95}1995. For kHz QPOs  we refer to van der Klis \\markcite{kl97}1997.)   GX 13+1 has been classified as an atoll source, although of all atoll sources it shows properties which are closest to that seen in the Z sources (HK89). Moreover, Schulz, Hasinger \\& Tr\\\"umper\\markcite{sc89} (1989) put GX 13+1 among the luminous sources that have been classified as Z sources. Together with GX 3+1, GX 9+1, and GX 9+9, GX 13+1 forms the subclass of the persistently bright atoll sources. In the CD they have only been seen to trace out banana branches and their power spectra can be described by relatively strong ($\\sim$3.5\\% rms) VLFN and weak ($\\sim$2.5\\% rms) HFN, as compared to other atoll sources. No QPOs have been found before in these sources, neither at frequencies below 100 Hz nor at kHz frequencies (Wijnands, van der Klis \\& van Paradijs\\markcite{wi97a} 1997a; Strohmayer et al.\\markcite{st97} 1997). For GX 13+1, so far no HFN has been observed. Its banana branch resembled a more or less straight strip in the CD, whereas the other three sources showed more curved banana branches (HK89). Stella, White \\& Taylor \\markcite{st85}(1985) have reported bimodal behavior of GX 13+1 in the hardness-intensity diagram (HID). In one state the spectral hardness was correlated with count rate, while in the other it was anticorrelated. The transition between the two states occurred within one hour.  The main difference between atoll and Z sources is the mass accretion rate, $\\dot{M}$. Atoll sources accrete at $\\dot{M}\\la0.5\\dot{M}_{Edd}$, whereas Z sources accrete at near Eddington rates.  It was proposed by HK89 that a second difference lies in the strength of the magnetic field strength $B$ of the accreting neutron star. Recent spectral modeling by Psaltis \\& Lamb \\markcite{ps96}(1996) suggests that indeed atoll sources have $B<10^9$ Gauss and Z sources  $B\\sim10^9$--$10^{10}$ Gauss. The bright atoll sources are found to have a higher inferred $B$ than the low luminosity atoll sources, making them the best candidates to show Z source HBO.  In this Letter we report the  discovery of a 57--69 Hz  QPO and of a two branched structure in  the CD and HID of  GX 13+1. We suggest that the QPO is the same phenomenon as Z source HBO. ",
        "conclusions": "We have discovered a QPO in GX 13+1 between 57 and 65 Hz and around 69 Hz. It is the first time a QPO has been found in a persistently bright atoll source. The rms amplitude of the QPO increased with photon energy. Together with the QPO, a peaked noise component was found, probably atoll source HFN, with a cut-off frequency of $\\sim$2 Hz. The HFN was only detected when the QPO was present. In the October data the QPO was only found in the upper branch, i.e. at high count rates. On the upper branch the QPO frequency increased with count rate from $\\sim$57 Hz to $\\sim$65 Hz. In the November data the QPO could only be detected at high count rates. Assuming that $\\dot{M}$ and count rate were positively correlated, the QPO frequency increased with $\\dot{M}$, at least on the upper branch in the October observation. However, during the November observation we found the QPO at $\\sim$69 Hz; the mean count rate was then lower than during the October observation of the $\\sim$65 Hz QPO.  The pattern traced out by GX 13+1 in the CD and HID in our observations does not resemble the patterns traced out by other atoll sources. Atoll sources trace out islands and/or a banana branch. During EXOSAT observations GX 13+1 traced out a banana branch (HK89) which is quite different from what we observed with RXTE. We also do not find evidence for the bimodal behavior found by Stella et al. \\markcite{st85}(1985) in the sense that both our branches in the HID show a positive correlation between hard color and count rate. The relatively sharp turn in the CD and HID may be related to one of the vertices in the patterns traced out by Z sources.  Comparison with EXOSAT observations, using our RXTE energy spectra folded with the EXOSAT response matrix, shows that the source was $\\sim$30\\% brighter during the RXTE observations. (Note that in the RXTE ASM lightcurve the source shows intrinsic variations of $\\sim$50\\%; during  our observations the ASM count rate was near average.) This might explain why the two branched structure has not been seen before in GX 13+1. In any case EXOSAT was not sensitive enough to have detected the QPO reported in this Letter.  Recently Stella \\& Vietri \\markcite{st98}(1998) proposed that at least some of the $<100$ Hz QPOs observed in atoll sources are due to Lense-Thirring precession of the inner part of the accretion disk. In order to test this model one needs the frequencies of the simultaneously observed kHz QPOs. Since no kHz QPOs have been observed in GX 13+1 it is not possible to test this model. The upper limits on kHz QPOs are comparable with those found in the other members of the subclass (Wijnands et al. \\markcite{wi97a}1997a; Strohmayer et al. \\markcite{st97}1997)   The QPO properties (frequency,  dependence on $\\dot{M}$, energy spectrum, and the simultaneous presence of a band limited noise component [the HFN]) are all similar to those found for HBO in Z sources. On the basis of the above described similarities we suggest that the QPO we found is the same phenomenon as the HBO, and that the HFN component is related to the QPO in a similar way as Z source LFN to HBO. The identification of atoll source HFN with Z source LFN was previously proposed by van der Klis \\markcite{kl94}(1994).  The HBO in Z sources can be explained by the magnetospheric beat frequency model (Alpar \\& Shaham\\markcite{al85}, 1985; Lamb et al.\\markcite{la85}, 1985) according to which the observed QPO frequency is the difference between the neutron star spin frequency and the frequency at which blobs of matter orbit the neutron star at the magnetospheric radius. In this model the frequency increases with $\\dot{M}$ and decreases with neutron star spin frequency and $B$.  The HBO is observed at frequencies between 15 and 60 Hz. LFN is found with a cut-off frequency between 2 and 20 Hz; it usually appears and disappears together with the HBO. In quantitative models the LFN is naturally produced as an extra component to the HBO, with a total power that is comparable to that in the HBO (Shibazaki \\& Lamb \\markcite{sh87}1987). Both components get stronger with photon energy.  According to HK89 the properties of atoll and Z sources are determined by $\\dot{M}$ and $B$. Spectral modeling based on this picture shows that of all atoll sources, the persistently bright subclass, especially GX 13+1, have $\\dot{M}$ and $B$ closest to those of the Z sources (Psaltis \\& Lamb \\markcite{ps96}1996). On the basis of the magnetospheric beat frequency model one might therefore expect to see HBO-like phenomena in these sources. The detection of a HBO-like phenomenon in GX 13+1 seems to confirm these expectations.  The fact that the QPO frequency in GX 13+1 is at the high end of the HBO frequency range in Z sources can be explained by a slower spinning neutron star or/and by a $B$ that is lower than in the Z sources, but still high enough to produce HBO. A higher $\\dot{M}$ than Z sources is unlikely in view of the luminosity."
    },
    "9803/astro-ph9803098_arXiv.txt": {
        "abstract": "Under the assumption that the gas mass fraction of galaxy clusters estimated out to an outer hydrostatic radius is constant, it is possible to constrain the cosmological parameters by using the angular diameter distance relation with redshift. We applied this to a sample of galaxy clusters from redshifts of 0.1 to 0.94, for which published gas and total masses are available from X-ray data. After scaling the gas fraction values to the $r_{500}$ radius (Evrard, Metzler, \\& Navarro 1996), we find an apparent decrease in gas fractions at high redshifts, which can be scaled back to the mean gas fraction value at z $\\sim$ 0.1 to 0.2 of (0.060 $\\pm$ 0.002) $h^{-3/2}$, when $\\Omega_m = 0.55^{+0.35}_{-0.23}$ ($\\Omega_m + \\Omega_{\\Lambda} =1$, 1-$\\sigma$ statistical error). However, various sources of systematic errors can contribute to the change in gas mass fraction from one cluster to another, and we discuss such potential problems in this method. ",
        "introduction": "In recent years, measurements of gas mass fraction in galaxy clusters together with the universal value for the cosmological baryon density as derived from nucleosynthesis arguments have been used to constrain the cosmological mass density of the universe. There are two basic assumptions used in such an analysis: (1) that the galaxy clusters are the largest virialized systems in the universe and based on hierarchical clustering models that clusters represent composition of the universe as whole, and (2) that the gas mass fraction when measured out to a standard (hydrostatic) radius is constant. Evrard (1997) applied this argument to a sample of galaxy clusters studied by David, Jones \\& Forman (1995) and White \\& Fabian (1995), with redshifts up to 0.2. The gas mass fraction values were scaled to the $r_{500}$ radius first defined by Evrard, Metzler \\& Navarro (1996), which has shown to be a conservative estimate of the outer hydrostatic radius based on numerical simulations. When both gas and total (virial) masses are estimated out to this radius, then the assumption that all clusters should have the same gas mass fraction is well justified. The mean gas mass fraction, as derived by Evrard (1997), for the low redshift clusters is $\\bar{f}^{\\rm X-ray}_{\\rm gas} (r_{500}) = (0.060 \\pm 0.003)\\, h^{-3/2}$. Recently, Myers {\\it et al.} (1997) observed the Sunyaev-Zel'dovich (SZ) effect  (Sunyaev \\& Zel'dovich 1980) in a sample of low redshift clusters, and derived a mean SZ gas mass fraction, $\\bar{f}^{\\rm SZ}_{\\rm gas}$, that ranges from (0.061 $\\pm$ 0.011) $h^{-1}$ to (0.087 $\\pm$ 0.030) $h^{-1}$ (we refer the reader to Birkinshaw 1998 for a recent review on the SZ effect). Using the mean cluster gas fraction and the nucleosynthesis derived value for the $\\Omega_b$, both Evrard (1997) and Myers {\\it et al.} (1997) put constraints on the cosmological mass density of the universe, $\\Omega_m$, with some dependence on the Hubble constant.  If cluster gas fractions are indeed constant in a sample of galaxy clusters when calculated out to the outer hydrostatic radius, then it is possible to constrain the cosmological parameters based on their dependence in the angular diameter distance relation with redshift. In Section 2 of this paper, we present the potential possibility of using cluster gas fraction as a {\\it standard candle} and apply this to a sample of $\\sim$ 53 clusters that range in redshifts from 0.1 to 0.94, for which published data on masses are available from literature. In Section 3, we discuss various systematics uncertainties and possible selection effects in this method. ",
        "conclusions": "Under the assumption that the baryonic (gas) fraction in galaxy clusters are constant, we have constrained the cosmological parameters using the angular diameter distance relation with redshift. However, the present X-ray based data from various studies may be providing a biased result due to unknown selection effects. When the angular diameter distance dependence on the traditional Hubble constant measurements using SZ and X-ray data are included with the present method involving cluster gas mass fraction, it is likely that tighter constraints on the cosmological parameters are possible. We hope to explore this possibility in a future paper. Given the biases that may go in to the presented gas mass fraction vs. redshift diagram, we suggest that results from a well defined and random sample of galaxy clusters, such as that would be available from AXAF in X-ray and from interferometric observations in SZ, be used to constrain the cosmological parameters using the angular diameter distance relation with redshift."
    },
    "9803/astro-ph9803051_arXiv.txt": {
        "abstract": "We present the results from {\\it ASCA} observations of NGC~3783 carried out during 1993 and 1996. Variability is observed both between the two epochs and within the individual observations, the latter due to a non-stationary process.  We find the spectra at both epochs to contain features due to absorption by ionized material, predominantly due to O{\\sc vii} and O{\\sc viii}. We find the opacity for such material to decrease by a factor $\\sim$2 in the $\\sim 0.7$--1.0~keV band between the epochs while the photoionizing flux increases by $\\sim$25\\%. By comparison with detailed photoionization models we show this behaviour is inconsistent with that expected from a single-zone of photoionized gas. A deficit in the data compared to such models, consistent with excess absorption by O{\\sc viii}, supports the suggestion that the material consists of two or more zones  Significant Fe $K$-shell emission is also observed during both epochs. We show this emission is dominated by the asymmetric line profile expected from the innermost regions of the accretion flow. We find no evidence that the Fe emission varied in either shape or equivalent width between the two epochs. ",
        "introduction": "\\label{Sec:intro}   The almost face-on spiral (SBa) galaxy NGC~3783 ($z = 0.0097$) is one of the brighter Seyfert galaxies in the sky in all wavebands (e.g. Alloin et al. 1995 and references therein). At X-ray wavelengths, NGC 3783 was first detected in the {\\it Ariel-V} sky-survey and has subsequently been observed by all major X-ray instruments (e.g. Malizia et al 1997 and references therein). Evidence for deep absorption features in the 0.5--1.5~keV band was first obtained in a {\\it ROSAT}\\ observation (Turner et al. 1993) and confirmed in a subsequent {\\it ASCA}\\ observation (George, Turner \\& Netzer 1995, hereafter G95). The latter workers demonstrated the opacity was dominated by $K$-shell absorption edges due to O{\\sc vii} and O{\\sc viii} within a screen of material (with an effective hydrogen column density $\\sim 10^{22}\\ {\\rm cm^{-2}}$) within the column--of--sight to the central source of X-ray emission. Subsequent {\\it ASCA}\\ observations have revealed such screens of photoionized material (so--called  \"warm absorbers'') to be a common features in Seyfert galaxies (Reynolds 1997; George et al 1998a, hereafter G98). G95 also suggested the tentative detection of the emission predicted to arise from within such a screen of ionized gas if it subtends a significant solid angle at the central source.  Here we present the results from the analysis of four new {\\it ASCA}\\ observations of NGC~3783 performed in 1996 July, along with a re-analysis of the two previous observations carried out in 1993 December (and as reported previously by G95, Reynolds 1997 and G98). The observations performed in 1996 were each separated by $\\sim$2 days, with the hope (based on the behaviour seen during 1993) of measuring the reaction of the ionized screen to changes in the illuminating continuum, and hence better constrain the location of and physical conditions within the screen. Unfortunately, the photoionizing source did not co-operate, exhibiting only relatively low-amplitude variations during the 1996 observations. However, we as we shall show, useful insights regarding the nature of the ionized, circumnuclear material can still be obtained by comparision of the 1996 and 1993 observations.  The observations and data-screening employed are described in \\S\\ref{Sec:data_anal}. In \\S\\ref{Sec:spatial} we describe the results of our spatial and temporal analysis. The results of our spectral analysis of the absorption due to the ionized gas are presented in \\S\\ref{Sec:spectral-cont}, confirming the presence of the ionized material during both 1993 and 1996. However we find the opacity of the ionized material to decrease by a factor $\\sim$2 between these epochs, and discuss physically realistic explanations. In \\S\\ref{Sec:nldl} we discuss the constraints which can be placed on the emission due to Fe $K$-shell processes in the 5--7~keV band. Finally, in \\S\\ref{Sec:Discuss} we review our findings in the context of the results from similar objects and future observations. ",
        "conclusions": "\\label{Sec:Discuss}  \\subsection{The ionized-absorber} \\label{Sec:disc-ionized}  Our observations confirm the presence of substantial opacity due to ionized material within the column--of--sight to the X-ray emitting region in NGC~3783. Furthermore we have found this opacity in the $\\sim 0.7$--1.0~keV band to decrease by a factor $\\sim$2 in the 18 months between the observations carried out in 1993 and 1996. This large change in opacity is inconsistent with that expected if the ionized material is simply responding to changes in the illuminating continuum. Such behaviour has been seen in at least two other objects (MCG-6-30-15, Fabian et al 1994; NGC~3227, Ptak et al 1994, George et al 1998b).  During both epochs we find the bulk of the opacity to be reasonably well modelled by a screen of ionized material. However, in all cases, we find a deficit is present in the data/model residuals. In \\S\\ref{Sec:93vs96} we briefly explored a number of possible explanations for such a deficit. We suggest that the most attractive explanation is in terms of there being two or more zones of photoionized material (\\S\\ref{Sec:2zone}). In this case the column variability can be attributed to a change in column of just one of the zones, which might then be inferred to exist as clouds moving into and out of the column--of--sight. Such an hypothesis has been suggested in other objects based on the differential variability in the depth of the O{\\sc vii} and O{\\sc viii} edges (e.g. NGC~3227, Ptak et al 1994; MCG-6-30-15, Otani et al 1996; NGC~4051, Guainazzi et al 1996), and/or, as here, spectroscopically based on detailed photoionization calculations (e.g. NGC~3516, Kriss et al 1996). Of course future observations will reveal whether the zone referred to as the 'ambient gas' in \\S\\ref{Sec:2zone} truly always lies within the column--of--sight. A viable alternative is that this zone is in fact simply a slower moving (or larger) structure traversing the the column--of--sight.  Absorption features due to C{\\sc iv}, N{\\sc v} and possibly H{\\sc i} have been observed in the ultraviolet (UV) spectrum of NGC~3783 (e.g. Shields \\& Hamann 1997, and references therein). Furthermore there is good evidence that the implied column density in these ions is variable (e.g. Maran et al 1996). Based on the column densities implied by the 1993 {\\it ASCA}\\ observations of the O{\\sc vii} and O{\\sc viii} edges as given in G95, Shields \\& Hamann demonstrated that the predicted column densities in the UV absorption lines to be in moderately good agreement with the observed range. Hence it was suggested that there might be a single zone of ionized material responsible for the absorption features in both the X-ray and UV regimes. Since the UV absorption due to C{\\sc iv} is seen imprinted on the corresponding broad emission line, this would clearly place the location of the ionized material at a radius larger than the region emitting the bulk of the broad C{\\sc iv} emission ($\\gtrsim $ several light-days). This is in agreement with our estimate in \\S\\ref{Sec:nonequilibrium} that the material had to be at a radius $\\lesssim 20$ light-days in order to be in ionization equilibrium. In Table~6 we list the column densities of the abundant Li-like ions predicted from our single-zone model of the ambient gas implied by the {\\it ASCA}\\ observations in 1996. We find column densities in the strong UV absorption lines within an order of magnitude of those observed, confirming the basic conclusion of Shields \\& Hamann. However, as acknowledged by Shields \\& Hamann and several other workers (e.g. Mathur 1997 and references therein), such comparisons are extremely sensitive to the form of the photoionizing continuum in the unobserved 13.6--300~eV band, and the assumed elemental abundances. Given these facts, and the observed variability, more detailed comparisons require simultaneous X-ray and UV observations of a wider species of ions such as may be possible with the launch of {\\it AXAF}, {\\it XMM} and {\\it FUSE}.  Unfortunately, due to the uncertainities in the calibration of the XRT/SIS instrument below 0.6~keV (and hence our exclusion of these data from our detailed spectral analysis) our data are unable to provide stringent constraints on the intensity of the emission features predicted to arise within the ionized material. However, as noted in \\S\\ref{Sec:ion_emis}, the limits provided by the current data are consistent with the material subtending a solid angle $\\lesssim2\\pi$ at the central source.  \\subsection{The Fe emission line} \\label{Sec:disc-line}  Our data confirm the presence of intense Fe $K$-shell emission within NGC~3783, with a total eqivalent width ($\\sim 300$~eV) and broad, asymetric profile fairly typical of other Seyfert 1 galaxies (e.g. Nandra et al 1997b, and references within). In \\S\\ref{Sec:nldl} we suggest the line can be deconvolved into two components. The first is a narrow component with a rest-frame energy of 6.4~keV (corresponding to an ionization state $\\lesssim$Fe{\\sc xii}) with an equivalent width ($\\sim 60$~eV) consistent with production within the distant structures within the circumnuclear environment of the central source (Ghisellini, Haadt, Matt 1994). The second component, referred to as the diskline component in \\S\\ref{Sec:nldl}, is that arising close to $6$ gravitational radii of a putative supermassive black hole. The extreme gravitiational effects suffered by the emission line photons at such radii give rise to an asymmetric profile. We find no compelling evidence that the Fe emission varied in either shape or equivalent width between the 1993 and 1996, indicating no significant changes in the distribution of Fe emissivity."
    },
    "9803/astro-ph9803045_arXiv.txt": {
        "abstract": "This paper presents a new derivation of the Generalized Poisson distribution.  The derivation is based on the barrier crossing statistics of random walks associated with the Poisson distribution.  A simple interpretation of this model in terms of a single server queue is also included.  In the astrophysical context, the Generalized Poisson distribution is interesting because it provides a good fit to the evolved, Eulerian counts-in-cells distribution measured in numerical simulations of hierarchical clustering from Poisson initial conditions.  The new derivation presented here can be used to construct a useful analytic model of the evolution of clustering measured in these simulations.  The model is consistent with the assumption that, as the universe expands and the comoving sizes of regions change as a result of gravitational instability, the number of such expanding and contracting regions is conserved. The model neglects the influence of external tides on the evolution of such regions.  Indeed, in the context of this model, the Generalized Poisson distribution can be thought of as arising from a simple variant of the well-studied spherical collapse model, in which tidal effects are also neglected.  This has the following implication:  Insofar as the Generalized Poisson distribution derived from this model is a reasonable fit to the numerical simulation results, the counts-in-cells statistic must be relatively insensitive to such effects.  This may be a consequence of the Poisson initial condition.  The model can be understood as a simple generalization of the excursion set model which has recently been used to estimate the number density of collapsed, virialized halos.  The generalization developed here allows one to estimate the evolution of the spatial distribution of these halos, as well as their number density.  For example, it provides a framework within which the halo--halo correlation functions, at any epoch, can be computed analytically.  In the model, when halos first virialize, they are uncorrelated with each other.  This is in good agreement with the simulations.  Since it allows one to describe the spatial distribution of the halos and the mass simultaneously, the model allows one to estimate the extent to which these halos are biased tracers of the underlying matter distribution. ",
        "introduction": "Consider an initially Poisson distribution of particles that clusters gravitationally as the universe expands.  In this paper, the initial Poisson distribution will also be called the initial Lagrangian distribution.  As time passes, the particle distribution evolves, as, for example, tightly bound virialized clusters (called halos, or dark matter halos, in this paper) form. Thus, the evolved distribution is different from the initial Lagrangian distribution.  In what follows, the evolved distribution will be called the Eulerian distribution.  The goal of this paper is to use the properties of the initial Lagrangian distribution to derive a reasonable approximation to the form of the evolved Eulerian distribution. In the absence of a model relating the two distributions, the only constraint is that required by mass conservation:  the number of particles in the initial and evolved distributions is the same, so the average density, $\\bar n$, in the two distributions must be the same.  In what follows, quantities measured in the Lagrangian space will have a subscript `0', while those in Eulerian space will not.  In this notation, mass conservation implies that $\\bar n_0=\\bar n$.  Studies of clustering from Poisson initial conditions (Itoh, Inagaki \\& Saslaw 1993 and references therein) show that when the initial, Lagrangian distribution is Poisson, then the evolved Eulerian distribution is Generalized Poisson.  This paper presents a model in which this is so.  The model is consistent with three general hypotheses about the evolution of clustering. The first is the hypothesis that, in comoving coordinates, initially denser regions contract more rapidly than less dense regions, and that sufficiently underdense regions expand. The second assumption is that, as the universe evolves, the number of such expanding and contracting regions is conserved---only their comoving size changes.  The third is that the influence of external tides on the evolution of such comoving regions can be neglected, if one is only interested in computing statistics such as the mass function of collapsed halos, or the distribution of counts in Eulerian cells. There are no compelling physical arguments for any of these assumptions, and initial particle configurations which violate some or all of these assumptions are relatively easy to construct. That the model predicts a counts-in-cells distribution which is a reasonable approximation to that measured in the numerical simulations suggests that, at least for clustering from Poisson initial conditions, these simple assumptions may also be reasonably accurate.  \\subsection{The Generalized Poisson distribution}\\label{gpdsec} Since it plays a central role in this paper, various known properties of the Generalized Poisson distribution are summarized below.  The Generalized Poisson distribution (Consul 1989) has the form \\begin{equation} p(N|V,b) = {\\bar N(1-b)\\over N!}\\, \\Bigl[\\bar N(1-b) + Nb\\Bigr]^{N-1} {\\rm e}^{-\\bar N(1-b) - Nb}\\!. \\label{gpdf} \\end{equation} Here $p(N|V,b)$ is the probability that a cell of size $V$ placed randomly within a particle distribution contains exactly $N$ particles.  If $\\bar n$ denotes the average density, then $\\bar N\\equiv \\bar nV$.  In this paper $0\\le b<1$, and, for reasons discussed below, it will be supposed that $b$ is not a function of $V$.  The case $b=0$ is the Poisson distribution.  Equation~(\\ref{gpdf}) is a Compound Poisson distribution (e.g. Saslaw 1989); it arises if point sized clusters, called halos in the following, have a Poisson spatial distribution, and the probability a randomly chosen halo contains exactly $n$ particles is \\begin{equation} \\eta(n,b) = {(nb)^{n-1}\\,{\\rm e}^{-nb}\\over n!}. \\label{borel} \\end{equation} This is the Borel$(b)$ distribution (Borel 1942).  In this paper, equation~(\\ref{borel}) will be called the halo mass function.  The Generalized Poisson distribution was first discovered in the astrophysical context by Saslaw \\& Hamilton (1984) (also see Sheth 1995a).  It provides a good fit to the distribution of particle counts in randomly placed cells, provided the particle distributions have evolved, as a result of gravitational clustering, from an initially Poisson distribution (Itoh, Inagaki \\& Saslaw 1993 and references therein).  In fact, the fits are significantly improved if $b$ is allowed to increase to an asymptotic value as $V$ increases.  This scale dependence is simply a consequence of relaxing the assumption that Borel clusters are point sized, but still requiring that they have some finite size.  The asymptotic value of $b$ is that which would have characterized the distribution, had the clusters been point sized (Sheth \\& Saslaw 1994). For this reason, the asymptotic value of $b$ is fundamental, and the point sized idealization useful.  This paper is mainly concerned with the point sized idealization, so that, in what follows, $b$ is independent of $V$.  The point sized idealization is also motivated by the following observation.  To a good approximation, the distribution of bound virialized halos in the numerical simulations is Borel$(b)$. Thus, to a good approximation, clustering from Poisson initial conditions evolves in such a way that, at all times, particles are bound up in Borel$(b)$ halos, and, at the time when they first virialize, these halos have a Poisson distribution.  The evolution of clustering is parameterized by the time dependence of $b$; it is zero initially, and it increases as the universe expands (Zhan 1989; Sheth 1995b and references therein).  Therefore, in the remainder of this paper, $b$ will be treated as a pseudo-time variable, and the Borel$(b)$ distribution will often be called the halo mass function at the epoch $b$.  As $V\\to 0$, most cells in the Lagrangian and Eulerian distributions will be empty.  Equation~(\\ref{gpdf}) shows that, in this limit, the probability that a cell is not empty is $\\bar N(1-b)$, and the probability that a non-empty cell contains exactly $N$ particles is given by the Borel$(b)$ distribution. In other words, at the epoch $b$, the halo mass function is the same as the vanishing-cell-size limit of the Eulerian counts in cells distribution (Sheth 1996a).  This fact will be useful later.  The Borel$(b)$ distribution can be derived from a number of different constructions, all of which are related to the Poisson distribution (Epstein 1983; Sheth 1995b; Sheth 1996b; Sheth \\& Pitman 1997).  In the context of this paper, all these constructions can be thought of as providing models that allow one to compute the Eulerian space distribution, in the limit of vanishing cell size, given that the Lagrangian space distribution is Poisson. One of these constructions, based on the statistics of random walk  barrier crossings associated with the Poisson distribution, is the excursion set model (Epstein 1983; Sheth 1995b). This paper shows how to derive the Generalized Poisson distribution from a simple generalization of this excursion set model. The generalization shows how to derive the Eulerian space Generalized Poisson distribution from the Lagrangian space Poisson distribution, for all cell sizes, and all times.  \\subsection{Outline of this paper} Section~\\ref{const} summarizes the random walk, excursion set model which leads to the Borel$(b)$ distribution. Sections~\\ref{shift} and~\\ref{expdf} describe a generalization of this model which leads to a new derivation of the Generalized Poisson distribution. Section~\\ref{halos} shows how to describe the spatial distribution of virialized halos within the context of this model.  It shows that the model is consistent with the Compound Poisson interpretation of the Generalized Poisson distribution -- in the model, Borel$(b)$ halos have a Poisson distribution at the time when they first virialize.  Moreover, in the model, the $V\\to 0$ limit of the counts-in-cells distribution is, indeed, the halo mass function. This shows explicitly that the excursion set approach developed here is able to reproduce the known properties of the Generalized Poisson distribution.  The relation between this model and the well-studied spherical collapse model (outlined in Appendix~\\ref{scoll}) is discussed in Section~\\ref{scm}.  Section~\\ref{queue} contains a brief digression which relates the excursion set model of the previous section to a simple single server queue system.  Section~\\ref{scale} discusses a scaling solution associated with the model that is analogous to the scaling solution found in Section~3.2 of Sheth (1995b). Section~\\ref{twobar} discusses the associated two barrier problem.  The solution of this problem may provide useful diagnostics in assessing the rate of evolution of the Eulerian statistics computed earlier in the paper.  Clustering from more general initial conditions, using the techniques developed here, is treated in a forthcoming paper. ",
        "conclusions": "This paper presents a new derivation of the Generalized Poisson distribution.  The derivation allows one to construct a useful model of hierarchical clustering from Poisson initial conditions. The resulting model is useful because the Poisson assumption allows one to solve many problems that, at present, have no solution if more realistic initial conditions are used.  The model is a simple generalization of the excursion set model developed by Bond et al. (1991).  Their approach allows one to estimate the mass function of collapsed halos; the generalization presented here allows one to describe the spatial distribution of these halos as well. The model can also be thought of as a simple variant of the spherical collapse model.  In the model, initially denser regions contract more rapidly than less dense regions, sufficiently underdense regions expand, the influence of external tides on the evolution of such regions is ignored, and the number of expanding and contracting regions is assumed to be conserved.  Strictly speaking, none of these assumptions can be justified physically.  However, these simplifications mean that the model can be worked out relatively easily. Moreover, the Generalized Poisson distribution, derived after making these assumptions, is a reasonably accurate fit to the Eulerian counts-in-cells distribution measured in numerical simulations of clustering from Poisson initial conditions. This suggests that, at least for estimating the evolution of the counts-in-cells statistic from such initial conditions, these simplifications are justified.  In the model, a collapsed halo occupies a vanishingly small volume.  In the simulations, collapsed halos have non-zero sizes---any given halo virializes at some fraction, typically about one half, of its turnaround radius. This means that on scales smaller than that of a typical halo, the counts-in-cells distribution computed here will cease to be a good approximation to that measured in the simulations. As discussed in the introduction, the fact that halos have non-trivial density profiles means that the $b$ parameter in equation~(\\ref{gpdf}) depends on scale.  A reasonable approximation to the effects of this scale dependence can be computed from models, such as those proposed by Navarro, Frenk \\& White (1996), of the density profiles of collapsed halos (see Sheth \\& Saslaw 1994 for details).  As Poisson initial conditions are not realistic anyway, this seems an unnecessary refinement to an already idealized model.  As the basic model has worked out so easily, as it allows one to estimate the extent to which halos are biased tracers of the mass, and, most importantly, as it provides a reasonably accurate description of the evolution of clustering measured in numerical simulations, it seems worth extending it to describe clustering from more general initial conditions. This extension is in progress."
    },
    "9803/astro-ph9803273_arXiv.txt": {
        "abstract": "We report on Australia Telescope Compact Array observations of the $\\sim$$10^5$~yr old pulsar PSR~B0906--49. In an image containing only off-pulse emission, we find a weak, slightly extended source coincident with the pulsar's position, which we argue is best interpreted as a pulsar wind nebula (PWN). A trail of emission extending behind the pulsar aligns with the major axis of the PWN, and implies that the pulsar is moving north-west with projected velocity $\\sim$60~\\kms, consistent with its scintillation speed. The consequent density we infer for the pulsar's environment is $>2$~cm$^{-3}$, so that the PWN around PSR~B0906--49 is confined mainly by the high density of its surroundings rather than by the pulsar's velocity.  Other properties of the system such as the PWN's low luminosity and apparent steep spectrum, and the pulsar's large characteristic age, lead us to suggest that this nebula is substantially different from other radio PWNe, and may represent a transition between young pulsars with prominent radio PWNe and older pulsars for which no radio PWN has been detected.  We recommend that further searches for radio PWNe should be made as here: at low frequencies and with the pulsed emission subtracted. ",
        "introduction": "The spin-down observed in most pulsars corresponds to a significant rate of energy loss, which manifests itself primarily in the form of a magnetized wind of relativistic particles (e.g. Rees \\& Gunn 1974\\nocite{rg74}).  Under certain conditions the interaction between this wind and its surroundings is observable, in the form of a pulsar wind nebula (PWN). At radio frequencies PWNe fall into two basic classes, plerions and bow-shock nebulae: plerions (e.g.\\ \\cite{hb87}) are the filled-center components of supernova remnants (SNRs) in which the pulsar wind is confined by the pressure of hot gas in the SNR interior, while bow-shock nebulae are confined by the ram pressure associated with their pulsar's high velocity (e.g.\\ \\cite{fk91}; \\cite{fggd96}).  Both are characterized by significant levels of linear polarization, a centrally-peaked morphology and a flat spectral index ($-0.3 < \\alpha < 0$; $S_\\nu \\propto \\nu^{\\alpha}$).  While only a tiny fraction of a pulsar's spin-down energy goes into producing the radio emission from such nebulae, radio PWNe are valuable diagnostics of the properties of both the wind and of the pulsar itself (e.g.\\ \\cite{fggd96}). Insight into the workings of such PWNe is limited by the fact that only six pulsars, all with ages $\\la10^5$~yr, have associated radio PWNe (\\cite{fs97}, hereafter FS97). Attempts to detect radio nebulae around older pulsars have so far been unsuccessful (e.g. \\cite{ccgm83}).  In an attempt to increase the sample of radio PWNe, FS97\\nocite{fs97} used the Very Large Array at 8.4~GHz to search for PWNe around 35 pulsars of high spin-down luminosity and/or velocity, but found no new PWNe down to a surface brightness $T_b \\sim 1.2$~K. They concluded that only young, energetic pulsars produce observable radio nebulae, and that the properties of a pulsar's wind may change as the pulsar ages.  However, a weak, compact or steep-spectrum nebula is difficult to detect using the approach of FS97, because the nebula is ``hidden'' by the emission from its associated pulsar. This has motivated us to attempt an alternative strategy for finding radio PWNe, whereby visibilities are recorded at high time resolution so that off-pulse images can be produced. The full results of this program are described elsewhere (\\cite{sgjf98}); here we report on the discovery of an unusual PWN associated with PSR~B0906--49.  PSR~B0906--49 ($l = 270\\fdg3$, $b = -1\\fdg0$) is a 107~ms pulsar (\\cite{dmd+88}) whose characteristic age $\\tau_c =$~112\\,000~yr and spin-down luminosity $\\dot{E} = 4.9 \\times 10^{35}$~erg~s$^{-1}$ place it amongst the 5\\% youngest and most energetic of all pulsars. \\HI\\ absorption has yielded distances to the pulsar in the ranges 2.4--6.7~kpc (\\cite{kjww95}) and 6.3--7.7~kpc (\\cite{sdw+96}), while a distance estimate using the pulsar's dispersion measure (\\cite{tc93}) is $6.6^{+1.3}_{-0.9}$~kpc. In future discussion we assign it a distance $7d_7$~kpc. The pulse profile shows a strong main pulse and weaker interpulse, both of which are $\\sim$90\\% linearly polarized (\\cite{wmlq93}; \\cite{qmlg95}).  Scintillation measurements imply a transverse velocity for the pulsar of $\\sim$50$d_7^{1/2}$~\\kms\\ (\\cite{jnk98}). ",
        "conclusions": "Using pulsar-gating observations at $\\lambda=20$~cm, we have found an off-pulse, slightly extended source coincident with PSR~B0906--49, which we interpet as a faint pulsar wind nebula. An associated trail implies a projected velocity for the pulsar of $\\sim$60~\\kms\\ along a position angle 315\\arcdeg.  The system is unusual in several ways:  \\begin{enumerate}  \\item the nebula has a lower luminosity and steeper spectrum than any other radio PWN yet discovered;  \\item PSR~B0906--49 is older than any other pulsar known to power a radio PWN, while only PSR~B1853+01 has a lower spin-down luminosity;  \\item the PWN appears to be generated by a slow ($\\sim$60~\\kms) pulsar moving through a dense ($>2$~cm$^{-3}$) medium. \\end{enumerate}  Thus PWN~B0906--49 appears to be very different from other radio PWNe. To account for this, we propose that either the low velocity and high density inhibits shock acceleration, or that the particle spectrum injected by a pulsar into its PWN steepens with age. In the latter case, PWN~B0906--49 may be in transition between the high luminosity, flat-spectrum PWNe seen around young pulsars, and the undetectable radio PWNe around older pulsars. We note that PWNe of low flux density and steep spectrum were heavily selected against in FS97's ungated, high frequency observations. We therefore recommend that future searches, particularly around intermediate-age pulsars, should be made as here: at low frequencies, and by imaging only the off-pulse emission.  A collection of PWNe produced under a variety of conditions still needs to be accumulated before we can fully understand how the radio emission from a PWN is produced and how it depends on the properties of its pulsar."
    },
    "9803/astro-ph9803335_arXiv.txt": {
        "abstract": "We report on the detection of four primordial galaxies candidates in the redshift range $4.55 \\le z \\le  4.76$, as well as 4 other candidates in the range $4.12 \\le z \\le 4.32$. These galaxies have been detected with the now common technique of the Lyman break, but with a very original instrumental setup, which allows the simultaneous detection on a single exposure of three narrow ($\\mathrm{FWHM} \\sim 200$~\\AA) bands. Depending on between which pair of bands the Lyman break is located, it thus allows to detect two ranges of redshifts. Two of the the higher redshift candidates are located only $50\\h50$~kpc apart, thus forming a possible primordial galaxy pair. ",
        "introduction": "Introduction} The knowledge of the epoch when galaxies formed, that is when they formed the bulk of their stars, is a major cosmological question. The detection of such primordial galaxies would give constraints on their formation, their chemical evolution or the formation of large scale structures. All the attempts to detect these galaxies at redshifts $z\\ge3$ by detecting the Ly$\\alpha$ emission line, either by visible and IR spectroscopy (Koo \\& Kron 1980, Schneider et al. 1991, Thompson \\& Djorgovski 1995), narrow band imaging (De Propris et al. 1993, Parkes et al. 1994) or Perot-Fabry (Thompson et al. 1995) have failed, while the Ly$\\alpha$ flux of a $z=4$ galaxy, $\\sim 10^{-16\\pm 1}$~erg~cm$^{-2}$~s$^{-1}$ (Thompson \\& Djorgovski 1995), should have been easily observed with 3.6m telescopes. Only recently, Hu (1998) reported the detection of \\lya\\ emitters at redshifts as high as 4.5. However, it is now known that even small amounts of dust can absorb the great majority of the Ly$\\alpha$ photons by resonant scattering (Charlot \\& Fall 1993).  The most promising method to observe these galaxies seems to use the Lyman break at 912~\\AA. This break is present in all kinds of galaxies, for all reasonable stellar formation scenarii (Bruzual \\& Charlot 1993), and is observed in the visible range for $z\\ge 3.5$.  The first very high redshift galaxies have been observed through multi band imaging, with broad filters located on each side of the Lyman break at $\\left<z\\right>=3.2$, selected using color criteria, and several of the candidates have recently been spectroscopically confirmed with the Keck telescope (Steidel \\& Hamilton 1992,1993, Steidel et al. 1995, Pettini et al. 1997). The average projected density of these objects is expected to be 0.5~gal~arcmin$^{-2}$ at $R \\sim 25$ and $\\left<z\\right> = 3.2$ (Steidel et al. 1995).  In this paper, we present the results of the search for very high redshift galaxies ($z\\sim4.5$ ), using the same technique, but with an original instrumental setup : we have designed multi band filters, with three narrow ($\\mathrm{FWHM} \\simeq 200$~\\AA) bands, placed so that they avoid the most prominent night sky emission lines. This disposition and the very high transmission (above 95\\%) of the filters compensate the narrowness of the bands, and reduces by a factor of three the amount of data to reduce, as well as telescope observing time.  Observations are described in Sect.~\\ref{data}, and our detections are discussed in Sect.~\\ref{res}. ",
        "conclusions": ""
    },
    "9803/astro-ph9803179_arXiv.txt": {
        "abstract": "The existence of bimodal disks is investigated.  Following a simple argument based on energetic considerations we show that stationary, bimodal accretion disk models in which a Shakura--Sunyaev disk (SSD) at large radii matches an advection dominated accretion flow (ADAF) at smaller radii are never possible using the standard slim disk approach, unless some extra energy flux is present. The same argument, however, predicts the possibility of a transition from an outer Shapiro--Lightman--Eardley (SLE) disk to an ADAF, and from a SLE disk to a SSD. Both types of solutions have been found. We show that for any given accretion rate a whole family of bimodal SLE--ADAF and SLE--SSD solutions can be constructed, each model being characterized by the location of the transition radius. ",
        "introduction": "Observations of black hole X--ray binaries (BHXBs, e.g.~V404 Cyg, Nova Muscae, A0620-00, GRO J1655-40, Cyg X-1) suggest that a bimodal accretion disk may be present in these sources (Thorne \\& Price \\cite{thorneprice:1975}; Shapiro, Lightman \\& Eardley \\cite{sle:1976}; Narayan, McClintock \\& Yi~\\cite{narclintyi:1996}; Esin, McClintock \\& Narayan \\cite{esinmcclnar:1997}). The inner part is a hot, optically thin, quasi--spherical accretion flow, producing a non--thermal spectrum via synchrotron radiation and Comptonization, with a cut--off around 100 keV.  The outer part is a geometrically thin, optically thick disk (Shakura \\& Sunyaev \\cite{shaksuny:1973}; Frank, King \\& Raine \\cite{frankkingraine:1992}) which emits a multi--color blackbody spectrum peaked around a few keV. By making plausible assumptions on how the transition radius $R_{\\trans}$ depends on the accretion rate $\\dot M$, one obtains a one--parameter family of spectra, once the black hole mass $M$ and the viscosity parameter $\\alpha$ are specified. The different spectral states of the BHXBs can then be explained by varying the accretion rate $\\dot M= \\dot m\\dot M_{Edd}$, where $\\dot M_{Edd}$ is the Eddington accretion rate.  Historically the hot, inner region was first thought to be a cooling--dominated, optically thin accretion disk, the SLE disk (Shapiro, Lightman \\& Eardley \\cite{sle:1976}, henceforth SLE). However, such disks are subject to a violent thermal instability, which makes them useless for constructing stationary bimodal models. An alternative to the SLE disk could be the recently discovered hot, advection dominated accretion flows (ADAFs: Narayan \\& Yi \\cite{narayanyi:1994}, \\cite{narayanyi:1995}; Abramowicz et al. \\cite{abrchen:1995}), which are stable against both thermal and viscous perturbations. ADAF models have been quite successful in reproducing the spectral properties of BHXBs in quiescence and of the Galactic center source Sgr A$^{*}$ (Narayan, Yi \\& Mahadevan \\cite{naryimah:1995}). A bimodal disk structure, in which an outer Shakura--Sunyaev disk (SSD) connects smoothly to the inner ADAF, seems very promising in explaining the whole range of spectral states of BHXBs, from the quiescent to the high state (Esin, McClintock \\& Narayan \\cite{esinmcclnar:1997}).  Advection dominated accretion flows are dynamically very stable.  The robustness of the ADAF solutions led Narayan \\& Yi (\\cite{narayanyi:1995}) to suggest that perhaps nature will always select this option whenever it becomes available. Yet, the precise mechanism through which the SSD material is converted into an ADAF remains a matter of debate. Observations of Cyg X-1 in the low state suggested that the transition from a cold to a hot disk might result from the secular instability of the radiation--pressure dominated inner part of the SSD (SLE; Ichimaru \\cite{ichimaru:1977}). The even more violent thermal instability of the radiation--pressure dominated region (Pringle, Rees \\& Pacholczyk \\cite{pringreespac:1973}; Piran \\cite{piran:1978}) might also be held responsible for such a transition. The drawback is that the transition radius for these models is usually quite close to the black hole, $\\rad_{tr}\\simeq 45\\alpha^{2/21}\\dot m^{16/21}m_{*}^{2/21}\\rad_g$, where $m_{*}\\equiv M/M_{\\odot}$ and $\\rad_g = 2GM/c^2$ is the Schwarzschild radius (SLE; Frank, King \\& Raine~\\cite{frankkingraine:1992}), whereas observations of many BHXBs (e.g.~Narayan, Barret \\& McClintock \\cite{narbarclint:1997}; Hameury et al.~\\cite{hamlasclnar:1997}) seem to imply a transition radius $\\rad_{tr}\\sim 10^4 \\rad_g$. Moreover, no self--consistent global solution for such a viscously/thermally driven transition has been found so far, and, more probably, the thermal instability results in a genuinely time--dependent behavior (e.g.~a limit cycle, see Meyer \\& Meyer--Hofmeister \\cite{meyer2hofm:1981}, Mineshige \\& Wheeler \\cite{minwheel:1989}; Honma, Matsumoto \\& Kato \\cite{honmatskat:1993}; Szuszkiewicz \\& Miller \\cite{szuszkmiller1:1997}, \\cite{szuszkmiller2:1997}), rather than in a stationary transition to a hot flow. More promising in this respect are the evaporation models in which matter evaporates from the SSD forming an advection dominated corona (Meyer \\& Meyer--Hofmeister \\cite{meyermeyhof:1994}; Narayan \\& Yi \\cite{narayanyi:1995}).  The question remains whether we understand why it seems not to be possible to convert an SSD directly into an ADAF at a certain radius. This is the question we address in this paper. We use a simple argument based on energetics to predict which kind of bimodal disk structures should be possible. We conclude that only bimodal accretion flows which are formed by a SLE disk outside and by an ADAF or SSD inside are allowed. We confirm our qualitative prediction by actually constructing global models for both bimodal SLE--ADAF and SLE--SSD solutions. ",
        "conclusions": "The goal of this investigation has been to shed some light on the physics of disk transitions. Using a simple argument based on energetics, we have shown that global, stationary bimodal solutions are possible if the outer disk satisfies a simple condition. Our main result is that only an SLE disk can match an inner ADAF (or SSD), while a bimodal accretion configuration formed by an outer SSD and an inner ADAF is not permitted, at least within the standard slim disk physics.  The two allowed bimodal structures, an outer SLE disk glued to an inner ADAF or SSD, have been computed numerically. We found that, in both cases, the transition radius turns out to be arbitrary, so a whole family of solutions exists. This may sound odd, since there is no physical reason to ask the sound speed to take a precise value at an intermediate radius, as we did in finding our numerical solutions. In principle the condition should be placed at the outer edge. In doing so, however, one should be extremely accurate in fixing $c_s(\\rad_{out})$ to find the transition at the desired radius since most bimodal solutions correspond to very nearly identical values of $c_s$ (or any other variable) at the outer or inner edge.  It remains nevertheless true that the structure of bimodal disks is determined once a complete set of conditions is prescribed at the edges and at the sonic radius, without the need of any extra piece of physics.  The existence of the SLE--ADAF and SLE--SSD bimodal solutions agrees with the energy argument of section \\ref{sec-therm-trans}, and strengthens its plausibility.  The same argument indicates that a SSD--ADAF model is not allowed. There is a caveat here, however. The essence of the energy argument lies in recognizing that the non--locality of the advective cooling term in the energy equation can be removed by performing a transformation to the fluid frame, and doing the thermal analysis in the frame comoving with the fluid. If, on the other hand, thermal conduction is introduced, the energy equation becomes a second order diffusion type equation and there is no way of removing the non--locality by changing the observer's frame. In this case our argument definitely does not apply.  Honma (\\cite{honma:1996}) has shown that SSD--ADAF structures are possible when the thermal heat flux from the hot ADAF into the SSD is accounted for. In such accretion flows the location of the transition follows self--consistently from the model, contrary to what we have found in this paper.  This is because the heat flux will evaporate per unit time an amount of matter from the SSD which needs to be in balance with the actual accretion rate.  The transition radius will therefore automatically adjust itself until this balance is reached.  In concluding, we would like to state clearly that, because of the thermal instability of the SLE disk, the models presented here most probably do not exist in nature."
    },
    "9803/astro-ph9803209_arXiv.txt": {
        "abstract": "We derive an analytic expression for the intensity of resonance-line radiation ``trapped'' in a semi-infinite medium. Given a source function and destruction probability per scattering, the radiation pressure due to trapped photons can be calculated by numerically integrating over analytic functions. We apply this formalism to a plane-parallel model stellar atmosphere to calculate the radiation pressure due to Lyman-$\\alpha$ photons produced following absorption of UV and X-rays from an AGN. For low surface gravity stars near the AGN ($g \\sim 10\\:\\cmss,\\ r \\sim 0.25 {\\rm\\:pc}$), we find that the pressure due to Lyman-$\\alpha$ photons becomes an appreciable fraction of that required for hydrostatic support. If the broad emission line emitting gas in AGNs and QSOs consists of stellar outflows, it may be driven, in part, by Lyman-$\\alpha$ pressure. ",
        "introduction": "\\label{intro}  Interest in the effect of hard UV and X-ray radiation on stellar atmospheres and stellar evolution began in earnest with work on X-ray binaries (e.g., \\cite{DOst73}, \\cite{BSun73}, and \\cite{Arons73}), but the subject has since been revisited in the context of active galactic nuclei (AGN).  Consider a star of effective temperature $T_{\\rm eff}$ and an AGN of luminosity $L = 10^{46}L_{46}\\:\\ergs$. The incident power per area from the AGN is equal to $\\sigma T_{\\rm eff}^4$ at the ``heating'' distance \\begin{equation} d_h = \\left(\\frac{L}{4\\pi\\sigma T_{\\rm eff}^4}\\right)^{\\onehalf} = 1.5\\times 10^{17} L_{46}^{\\onehalf}\\left( \\frac{5000 {\\rm K}}{T_{\\rm eff}}\\right)^2 \\:{\\rm cm}; \\label{eq:rheat} \\end{equation} heating of the stellar surface by the AGN is important at distances $d \\lesssim d_h$.  Fabian (1979) was apparently the first to suggest that heating by AGN radiation might substantially affect the envelopes of nearby stars, resulting in enhanced mass loss. Edwards (1980) argued that when a star approached within $d \\lesssim d_h$, the irradiated photosphere would develop supersonic horizontal winds carrying heat from the illuminated hemisphere to the shadowed side.  Edwards conjectured that enhanced mass loss could then occur, resulting in a ``cometary star'' as the stellar outflow is accelerated radially away from the AGN\\@. Matthews (1983) noted the importance of direct radiation pressure from the AGN, arguing that this could ablate matter tangentially from the stellar photospheres of giant or supergiant stars.  Voit \\& Shull (1988) discussed the effects of X-rays on the upper atmosphere of a star near the AGN\\@.  They argued that X-ray heating of a $\\sim 1 M_{\\sun},$ $R_{*} \\approx 100 R_{\\sun}$ star ($g\\approx 3\\:\\cmss$) would result in a hot, ionized wind, with temperature at the critical point $T \\approx 3\\times 10^5 {\\rm K}$; radiation pressure in C IV, N V, and O VI resonance lines would contribute to the acceleration of the wind.  They concluded that stars which are already red supergiants could develop winds with $\\dot M \\approx 10^{-7}L_{46}^{0.9}d_{18}^{-2}R_{*,100}^2 M_{\\sun} {\\rm yr^{-1}}$, where the distance from the AGN is $d = 10^{18} d_{18}{\\:\\rm cm}$, and $R_{*,100} \\equiv R_*/100R_{\\sun}$. These mass loss rates exceed normal mass loss rates for red supergiants only within a distance of $d_{18} \\lesssim 0.3 L_{46}^{0.45} R_{*,100}\\:{\\rm cm}$. They estimated the mass loss due to ablation by radiation pressure (\\cite{Matthews83}) for a $\\sim 1 M_\\odot$ star to be $\\dot M \\approx 10^{-8}L_{46}d_{18}^{-2}R_{*,100}^{5/2}$ provided $d_{18} < 0.6L_{46}^{\\onehalf} R_{*,100}$.  There has also been some exploration of how AGN radiation might affect the evolution of stars (\\cite{Matthews83}; \\cite{Verbuntetal84}; \\cite{Tout89}).  Tout \\etal discussed the evolution of stars immersed in a blackbody background with temperature $T = 10^{3.75 - 4} {\\rm K}$, concluding that the main sequence evolution is largely unaffected, but that the star will expand to a radius $\\sim 10$ times larger than usual during the ``red giant'' phase of evolution. Note, however, that the energy density of AGN radiation equals a $10^4{\\rm K}$ blackbody at a distance of only $2\\times 10^{16} L_{46}^{0.5}\\:{\\rm cm}$ from the AGN\\@.  Furthermore, it is not clear how their conclusions would have to be modified for the anisotropic and nonthermal radiation field of the AGN\\@.  These discussions of the effects of AGN radiation on stellar atmospheres and evolution have led some to consider a ``bloated stars scenario'' (BSS) to explain the origin of the AGN broad line region (BLR). The BSS proposes that the BLR consists of the stellar winds and/or expanded envelopes of stars irradiated with AGN radiation (e.g., \\cite{Pen88}, \\cite{Kaz89}).  The attractiveness of this scenario stems from its efficient use of known resources: stars are believed to be present near the ``central engines'' of AGN, their gravity provides a possible ``containment'' mechanism for BLR clouds, there is ample mass with which to replenish the clouds, and the mass shed provides material to fuel the AGN\\@.  So far, however, no specific model has successfully explained how all these mechanisms might function, though pieces of the puzzle have been explored (\\cite{AN94}).  High signal/noise measurements of line profiles indicate that the number of discrete clouds or bloated stars must be large, perhaps exceeding $\\sim 3\\times 10^6$ for Mrk~335 (Arav \\etal 1997) and $\\sim 3\\times 10^7$ for NGC~4151 (Arav \\etal 1998).  Such large numbers appear to challenge the BSS, but the difficulties may not be insurmountable.  Recent observations and photoionization models of Baldwin \\etal (1996) and Ferland \\etal (1996) offer some concrete evidence for the BSS\\@.  They find that the BLR seems to contain several distinct components.  One (component ``A'') has sharp (FWHM $\\approx 1000\\:\\kms$), symmetric line profiles centered on zero velocity, while another (component ``B'') appears to be outward-flowing, peaked at zero velocity but with a long blue tail (down to $-11,000\\:\\kms$). They interpret ``A'' as the expanded envelopes of stars, and ``B'' as their radiatively accelerated, outflowing winds.  The aim of the present paper is to call attention to the fact that irradiation of a star by hard X-rays from the AGN will result in the generation of Lyman-$\\alpha$ (\\lya) photons within the stellar atmosphere, and the pressure of these trapped resonance-line photons may contribute to mass loss.  The possible importance of \\lya\\ photons was suggested previously by Puetter [unpublished, cited in Penston (1988) and Voit \\& Shull (1988)].  Voit \\& Shull rejected Puetter's suggestion, arguing that the \\lya\\ pressure could be estimated to be \\begin{equation} P_{\\rm Ly\\alpha} \\approx \\left(\\frac{aT^4}{3}\\right)\\frac{W_\\lambda}{\\lambda} \\label{eq:PVoitShull} \\end{equation} where $W_\\lambda$ is the equivalent width of the \\lya\\ absorption profile of the gas between the point of interest and the surface, and $T$ is the gas temperature.  Voit \\& Shull then used $W_\\lambda/\\lambda < 1$ to obtain an upper limit $P_{\\rm Ly\\alpha} \\lesssim aT^4/3$ which, for stars of interest, is insufficient to ``bloat'' the atmosphere. However equation (\\ref{eq:PVoitShull}) presumes the radiation field within the gas to be close to a blackbody at the gas temperature $T$ --- an assumption which need not be valid when \\lya\\ photons are being generated within the gas by nonthermal processes, such as H(1s$\\rightarrow$2p) excitation by photoelectrons and secondary electrons in a non-Maxwellian ``tail'' to the electron energy distribution function.  It is therefore essential to estimate the rate of production of \\lya\\ photons, and to examine their diffusion in physical space and frequency space as well as the possibility of photon ``destruction'' by, for instance, collisional de-excitation of electronically excited H atoms.  In \\S\\ref{sec-linetrans}, we give a brief treatment of resonance-line transfer, extending the detailed calculations of Neufeld (1990).  In particular, we derive an expression for the case of a semi-infinite, absorbing slab of material.  Since resonance-line photons created in a stellar atmosphere will almost certainly not be able to scatter through the entire star without being absorbed, this limit is an excellent approximation.  In \\S\\ref{sec-model}, we apply our calculational method to a simple, plane parallel, irradiated model atmosphere. In \\S\\ref{sec-results}, we present the results. In \\S\\ref{sec-summ} we summarize and discuss implications and further directions for work. ",
        "conclusions": "\\label{sec-summ}  We have developed a formalism for computing the \\lya\\ pressure in a plane-parallel stellar atmosphere. The treatment of resonance-line trapping given in \\S 2 can be incorporated into already developed stellar atmosphere codes. For conditions believed representative of the BLR (ionizing flux of $\\sim 4\\times 10^{10}\\:\\ergcmcms$), stars with $g\\lesssim 10\\:\\cmss$ will develop an appreciable pressure due to trapped \\lya\\ photons in their atmospheres.  We have neglected numerous complicating factors, all of which appear likely to enhance the rate of mass loss. First of all, the gravitational acceleration $g\\propto r^{-2}$, rather than the assumed $g=\\mbox{constant}$.  As a result, the outer layers of the atmosphere will have a larger scale height and must ultimately go over into a thermal wind, even if the lower layers are essentially hydrostatic and stable.  With the outer layers of the atmosphere photoionized and heated to $T_4 \\gtrsim 1$, a fluid element will have positive enthalpy at $r \\gtrsim 600\\, T_4^{-1} (M/M_{\\sun}) R_{\\sun}$.  Furthermore, we have neglected the radiation pressure due to resonance lines other than \\lya, as well as the radiation from the star itself.  Ferland (1997) reports that including radiation pressure from all resonance lines rather than just a select few sometimes increases the total radiation by an order of magnitude (using CLOUDY's escape probability approximation, of course). Since no part of our derivation in \\S 2 depends explicitly on the resonance line being \\lya, it is straight forward to extend our calculation to include other resonance lines.  Most important, however, is the fact that except along the ``axis'' --- the line from the center of the star to the AGN --- the atmosphere will be subject to substantial tangential stresses due to three effects related to anisotropic irradiation by the AGN: \\begin{enumerate} \\item The momentum deposited by absorbed photons --- at anywhere other than the axis, the momentum will have a tangential component; \\item Transverse gradients in the \\lya\\ pressure --- at any particular radial column density $N_{\\rm H}$, the \\lya\\ production rate will be greater along the axis, where X-rays are entering radially; and \\item Transverse gradients in the gas temperature --- the rate of X-ray heating will be largest along the axis, where the zenith angle is zero. \\end{enumerate} These tangential stresses must result in significant tangential flows, with flow velocities that might approach escape speeds along the ``terminator,'' perhaps resulting in a ``cometary star'' as envisioned by Edwards (1980) and Matthews (1983).  It would obviously be of great interest to calculate axisymmetric fluid-dynamical models of red giant atmospheres subject to intense X-ray irradiation in order to investigate the effects of tangential stresses. It should be noted, of course, that we lack a quantitative theory for mass loss even from unperturbed evolved stars, so we are hardly poised to develop a definitive treatment including this new complication.  Nevertheless, it does appear likely that trapped \\lya\\ may help drive mass loss in low-$g$ stars subject to intense X-ray irradiation. The extra \\lya\\ radiation pressure and associated tangential stresses might induce such evolving stars to ``bloat'' further and lose mass prematurely, so that the population of stars near the AGN core could be skewed from the average throughout the galaxy.  Finally, we note that none of our calculations are dependent on the gravitating object being a star.  Recently, Walker and Wardle (1998) have suggested the existence of a population of cool, self-gravitating clouds in the galactic halo to account for ``Extreme Scattering Events'' which occur in monitoring the flux of compact radio sources. They propose clouds of approximately a Jovian mass with a radius of about an AU, which would correspond to a gravitational acceleration $g \\sim 6\\times 10^{-4}\\:\\cmss$ at the ``surface.''  With a total column density of $\\sim 10^{26}\\:{\\rm cm}^{-2}$, the majority of the AGN energy flux would be absorbed well before reaching (and disrupting) the center of the cloud. The proposed particle number density of $\\sim 10^{12}\\:{\\rm cm}^{-3}$ corresponds well to that inferred by Ferland \\etal (1996) and Baldwin \\etal (1996).  Thus if these self-gravitating, sub-stellar clouds exist in host galaxies of AGN, they seem to be potential candidates for the origin of the BLR\\@. It is unclear, though, how they could exist with a high metal content, as required by BLR observations, without rapidly cooling and collapsing (in much less than a Hubble time) prior to exposure to AGN irradiation.  Although we conclude that X-ray induced \\lya\\ pressure would not be significant for main sequence stars near AGN, it could be important either for lower gravity stars which have evolved off the main sequence up the giant branch, or for substellar, self-gravitating clouds. It remains to be seen whether the resulting cometary structures can account for the required BLR covering factor and cloud population, and the AGN mass supply."
    },
    "9803/astro-ph9803252_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "We have demonstrated the techniques which we have developed in the RVSAO suite to create a system for the accurate, automated reduction of spectra for galaxy redshifts and stellar radial velocities.  More than half of all published redshifts have been measured using these techniques, as well as a large number of stellar radial velocities.  The correlation method for obtaining redshifts can be successfully extended from absorption line spectra to emission line spectra, with a substantial improvement in effectiveness over the previous method for obtaining emission line redshifts, automated line fitting.  The reduction of emission-line spectra requires different reduction steps than absorption line correlations.  Emission line correlation redshifts are susceptible to blunders due to the presence of cosmic rays. However, using automated line fitting ({\\bf emsao}) and absorption line correlation velocities the blunder rate can be kept near zero, with the degree of automation kept high.  We have developed new techniques for calibrating and characterizing the blunder rate and the individual errors in redshift measurements. The blunder rate for RVSAO reductions can be kept near zero by the use of some simple heuristics to identify possible mistakes.  For typical redshift survey data from the FAST spectrograph the automation rate is 95\\%.  Our self calibrating internal error estimator is accurate to $\\sim$ 20\\%.  Large, stable surveys enable development of more accurate and stable error estimators.  We have developed new methods for creating, calibrating, and using galaxy redshift templates.  We have created an emission line template, femtemp97, having the median properties of a large set of strong emission line spectra.  We have created an absorption line template, fabtemp97, having the mean properties of a large set of absorption line spectra showing no sign of emission.  These spectra arise from physically distinct processes, and can be used to form a pair of basis vectors to perform a 2-D spectral classification.  We have developed a new method for establishing the zero point for redshift observations, a method which minimizes the systematic differences between emission and absorption line redshifts.  This zero point is determined as accurately as we can establish the wavelength calibration using standard HeNeArFe lamps.  We have shown a technique for measuring and eliminating differences between the instrumental zero point and the true zero point.  We have shown improved techniques for a number of the substeps necessary to obtain accurate redshifts, including: removal of emission(absorption) lines when correlating against an absorption(emission) line template; suppression of the night sky lines; supression of the continuum; design of the Fourier filter; and zero padding of spectra.  The rapid development of large aperture telescopes with multi-object spectrographs presents substantial challenges for redshift and radial velocity reductions.  Reducing one or two orders of magnitude more spectra of objects one or two orders of magnitude fainter while maintaining high quality control standards and minimal personnel costs is clearly a difficult problem. RVSAO provides a solid methodological and software basis to meet these new challenges."
    },
    "9803/astro-ph9803314_arXiv.txt": {
        "abstract": "We report the detection of faint emission in the high-excitation [\\ion{O}{iv}] 25.90$\\mu$m line in a number of starburst galaxies, from observations obtained with the Short Wavelength Spectrometer (SWS) on board ISO. Further observations of \\object{M~82} spatially resolve the [\\ion{O}{iv}] emitting region. Detection of this line in starbursts is surprising since it is not produced in measurable quantities in \\ion{H}{ii} regions around hot main-sequence stars, the dominant energy source of starburst galaxies. We discuss various models for the formation of this line. [\\ion{O}{iv}] that is spatially resolved by ISO cannot originate in a weak AGN and must be due to very hot stars or ionizing shocks related to the starburst activity. For low-excitation starbursts like \\object{M~82}, shocks are the most plausible source of [\\ion{O}{iv}] emission. ",
        "introduction": "Mid-infrared fine structure lines are powerful probes of dusty and obscured galactic nuclei, being able to penetrate extinctions up to the equivalent of  A$_V\\sim 50$. Using the Short Wavelength Spectrometer (SWS) on board the Infrared Space Observatory (ISO), it is possible to detect faint lines and sources. The rich observed spectra can be used for a detailed modelling of the ionizing spectra of starbursts (e.g. Rigopoulou et al. \\cite{rigo96}, Kunze et al. \\cite{kunze96}) and AGNs (Moorwood et al. \\cite{moor96}). Clear differences between their spectra make these lines a valuable new tool for discriminating between AGN and starburst activity in visually obscured galaxies. AGN spectra include emission from highly ionized species and the so-called coronal lines, requiring photons up to $\\sim$300\\,eV for their creation. In contrast, starburst spectra are dominated by lines of low excitation species, because even hot, massive stars emit few ionizing photons beyond the \\ion{He}{ii} edge at 54\\,eV.  Line ratios like [\\ion{O}{iv}]\\,25.9$\\mu$m / [\\ion{Ne}{ii}]\\,12.8$\\mu$m and [\\ion{Ne}{v}]\\,14.3$\\mu$m / [\\ion{Ne}{ii}]\\,12.8$\\mu$m have been used by Lutz et al. (\\cite{lutz96a}) and Genzel et al. (\\cite{genzel98}) to establish the dominant source of luminosity in ultraluminous infrared galaxies (ULIRGs). In some of the starburst templates studied, very faint [\\ion{O}{iv}] emission was found, about two orders of magnitude weaker than in typical AGNs. Faint [\\ion{O}{iv}] emission in starbursts is not relevant for establishing the power source of ULIRGs, but its origin poses an interesting problem because its creation ionization energy is slightly above the \\ion{He}{ii} edge. In this letter we examine possible mechanisms for its production. ",
        "conclusions": "We have discussed various excitation mechanisms for faint [\\ion{O}{iv}] emission from starburst galaxies. In general, starburst-related sources and in particular ionizing shocks provide the most plausible explanation. Weak buried AGNs may be plausible for individual sources but can be ruled out for the best studied case of \\object{M~82} whose [\\ion{O}{iv}] emitting region has been spatially resolved. In addition, the fairly small scatter in [\\ion{O}{iv}] versus starburst luminosity favours a starburst-related origin, since no finetuning of two independent mechanisms is required."
    },
    "9803/astro-ph9803122_arXiv.txt": {
        "abstract": "In the $z=1.94$ {\\CIV} absorption line system in the spectrum of quasar Q1222+228 ($z_{em}=2.04$), we find two clouds which have contrasting physical conditions, although they are only at a 17~{\\kms} velocity separation. In the first cloud {\\SiII}, {\\SiIV}, and {\\CII} are detected, and {\\AlII} and {\\AlIII} column density limits in conjunction with photoionization models allow us to infer that this cloud has a large Si abundance and a small Al abundance relative to a solar abundance pattern. This pattern resembles that of Galactic metal--poor halo stars, which must have formed from such high redshift gas. The second cloud, in contrast, has detected {\\AlII} and {\\AlIII} (also {\\SiIV} and {\\CII}), but no detected {\\SiII}. We demonstrate, using photoionization models, that Al/Si must be greater than (Al/Si)$_\\odot$ in this unusual cloud. Such a ratio is not found in absorption profiles looking through Milky Way gas. It cannot be explained by dust depletion since Al depletes more severely than Si. Comparing to other Al--rich environments, we speculate about the processes and conditions that could give rise to this abundance pattern. ",
        "introduction": "\\label{sec:intro}  One of the motives for studying quasar (QSO) absorption line systems is to document the photoionization structure, chemical composition, and kinematics of the ISM and halos of high redshift galaxies at a level of detail on par with studies of the Milky Way ISM and Halo. QSO lines of sight that pass through high redshift galaxies sample multiple gaseous structures having a variety of physical conditions such as found in Galactic {\\HI} and {\\HII} regions, supershells, infalling Halo gas, and material being processed at the Galaxy/Halo interface. Recognizing the abundance patterns, photoionization conditions, and kinematics associated with these various types of structures is a prerequisite for a detailed understanding of the evolution of galactic gas.  Based upon low resolution spectra, early QSO absorption line efforts have been limited to simple curve of growth analyses of equivalent widths, providing only a weighted--average of the absorbing gas properties in each system. A great deal has been learned from these studies and a global (statistically based) picture of chemical and ionization evolution has been suggested [see Steidel (1993\\nocite{steidel93}) and references therein]. The next logical step is to examine {\\it variations\\/} in the chemical and ionization conditions within single galaxies and to incorporate kinematic information. Using high resolution UV spectra, several researchers studied the Milky Way Disk and Halo at a detailed level (\\cite{wel97}; \\cite{savaraa}; \\cite{snow96}; \\cite{sf95}; \\cite{fs94}; \\cite{sf93}). Some recent efforts have focused on the cloud--by--cloud conditions in high redshift galaxies (cf.~\\cite{tripp97}; \\cite{pet94}). It is hoped that, ultimately, a statistical picture of the processes that give rise to the observed absorbing gas properties will improve our understanding of the present--epoch Milky Way and galactic evolution. In this {\\it Letter\\/} we study the cloud--to--cloud properties in a {\\CIV} system at $z \\sim 1.94$ along the line of sight toward the quasar Q$1222+228$ in order to: 1) demonstrate a large variation in ionization and/or abundance conditions between two kinematically adjacent clouds in the same absorber; and 2) present an unusual cloud that has an Al/Si abundance ratio enhanced by a factor of several relative to the solar ratio. ",
        "conclusions": "\\label{sec:summary}  Two kinematically adjacent, absorbing clouds in the same galaxy at $v\\sim0$~{\\kms} and $\\sim 17$~{\\kms} have very different abundance patterns, if they are in photoionization equilibrium. Based upon CLOUDY models, Cloud A is inferred to have Si enhanced relative to a solar abundance pattern, and Al underabundant relative to solar. This abundance pattern is characteristic of Milky Way Halo stars (\\cite{lau96}; \\cite{savaraa}) which could have formed from gas like that observed in this $z=1.94$ galaxy. We would expect such a pattern for many clouds in high redshift galaxies. Independent of the details, Cloud B has Al several times enhanced relative to Si (compared to the solar abundance pattern). This Al enhancement is highly unusual. If depletion onto dust grains is important we would expect a smaller Al/Si since Al depletes more readily than Si. To date, such an Al enhanced pattern has not been seen in absorption profiles looking through Milky Way interstellar gas (\\cite{savaraa}; \\cite{snow96}; \\cite{sf95}; \\cite{fs94}; \\cite{sf93}).  How could an enhancement of Al originate in a $z=1.94$ cloud? In an attempt to find clues we note three astrophysical environments in which Al enhancement is observed: 1) in the stellar photospheres of the most metal poor globular clusters (\\cite{shet96}; \\cite{smith96}); 2) in the broad line regions of some AGNs\\footnote{We note that the absorber lies 30,000 {\\kms} from the emission redshift of the QSO, which is not a BAL AGN.} (\\cite{shields97}); and 3) in the photospheres of Milky Way Bulge stars (\\cite{mr94}). The common theme for the enhancement of Al in these three seemingly different environments is novae. The novae could produce the Al either directly in their ejecta (\\cite{smith96}) or indirectly by providing magnesium isotopes to the ISM that later deposit onto stellar photospheres (\\cite{langer}). The $^{26}$Mg and $^{27}$Mg isotopes would then be converted to Al through proton capture in deep CNO convective mixing layers in metal poor stars (\\cite{langer}). More generally, does this suggest that a concentration of novae contributed to enhancing the Al in this cloud? In the globular cluster environment one key is to have a large number of stars that were formed coevally. Another key, which may apply to all three environments, is to have a potential well large enough to retain the gas. These may be prerequisites in our case also. Another possibility for excess Al production is a particular class of supernova for progenitors over a narrow mass range. The predicted amount of enhancement relative to other elements depends on the specific supernovae model adopted (\\cite{nom97}), however only a small subclass of models would give Al/Si in agreement with that inferred for this cloud.  Just how unusual is this cloud with a large Al/Si ratio? Such a pattern has not been seen in absorption along dozens of lines of sight toward Milky Way disk and halo stars. Most high redshift clouds do not have large {\\AlIII} and {\\AlII} (relative to {\\SiII}). However, in the $z = 2.14$ damped {\\Lya} absorber in Q$0528-251$ (\\cite{lu96}), we have identified a single outlying cloud (at $\\sim -100$~{\\kms}) for which the equivalent width of the {\\AlII} $\\lambda 1671$ transition is larger than that of the {\\SiII} $\\lambda 1527$ transition. The other clouds in the same system clearly have the opposite ratio. Finding more examples and establishing similarities between Al--rich environments are logical next steps in diagnosing the origin of this abundance pattern at high redshift and perhaps understanding the anomalous enhancements seen in metal--poor globular cluster and Bulge stars."
    },
    "9803/astro-ph9803158_arXiv.txt": {
        "abstract": "s{ We estimate the   dust polarized emission in our galaxy at high galactic latitudes, which is the dominant foreground for measuring CMB polarization using the high frequency instrument (HFI) aboard Planck surveyor.  We compare it with the level of CMB polarization and conclude that,  for angular scales $\\le 1^{\\circ}$, the  scalar-induced CMB polarization and temperature-polarization cross-correlation are  much larger than  the foreground level at $\\nu \\simeq 100 \\, \\rm GHz$.  The tensor-induced signals seem to be at best comparable to the foreground level.} ",
        "introduction": "The forthcoming satellite CMB projects MAP and Planck surveyor hold great promise for detecting the CMB polarization. A major stumbling block this  detection  is the unknown level of galactic polarized foregrounds. In this paper, we attempt to estimate the level of dust polarized emission at high galactic latitudes. We model this  emission  using the three-dimensional HI maps of the  Leiden/Dwingeloo survey at high galactic latitudes and  the fact that the dust emission, for a wide range of wavelengths, has a tight correlation with the HI emission maps of this survey~\\cite{boul}. Assuming the dust grains to be oblate with axis ratio $\\simeq 2/3$, which recent studies support~\\cite{hil}, we determine the intrinsic dust polarized emissivity. The distribution of magnetic field with respect to the dust grain distribution is quite uncertain, we thus  consider three extreme  cases (to be described below). ",
        "conclusions": ""
    },
    "9803/astro-ph9803228_arXiv.txt": {
        "abstract": "We investigate the implications of a very thick (scale height 1.5 - 3.0 kpc) disk population of MACHOs. Such a population represents a reasonable alternative to standard halo configurations of a lensing population.  We find that very thick disk distributions can lower the lens mass estimate derived from the microlensing data toward the LMC, although an average lens mass substantially below $0.3\\Msol$ is unlikely.  Constraints from direct searches for such lenses imply very low luminosity objects: thus thick disks do not solve the microlensing lens problem. We discuss further microlensing consequences of very thick disk populations, including an increased probability for parallax events. ",
        "introduction": "Current data from the MACHO collaboration \\cite{MACHOmass} indicate that in the context of a spherical isothermal model with a Maxwellian velocity distribution, some significant fraction of the Galactic halo is composed of MACHOs with masses roughly in the range 0.1 to 1.0 $\\Msol$. Such masses are consistent with several astrophysical candidates for MACHOs -- white dwarfs, neutron stars, and black holes -- each of which presents serious challenges for stellar formation and evolution theories. However, the MACHO component of the halo, if it is not the major component, as in Cold Dark Matter scenarios, may have a very different distribution from the typically assumed spherical isothermal model. The MACHO distribution may be in a significantly flattened halo and/or, due to dissipation, more centrally condensed. In addition, such a distribution might have a significant rotational component. These possibilities have strong implications not only for the MACHO fraction of the halo, but also for the mass estimates derived from the event durations.  The velocity dispersion and mass density distribution which describe the halo model are crucial input parameters in extracting an estimate of the lens mass from the data. The event duration is given by the radius of the Einstein ring divided by the MACHO velocity transverse to the line of sight, where the Einstein ring radius is a function of the position of the MACHO along the line of sight and the lens mass. The masses of the lenses can only be determined statistically, in the context of an a priori assumption about the distribution and velocity dispersion of the lenses. Such an analysis, using a spherical isothermal model yields a central mass estimate of $\\approx 0.4\\Msol$.  However, as discussed above, the MACHO distribution and velocity dispersion may be very different from that assumed in the standard halo model.  Earlier work \\cite{us_rotate} in exploring models with a highly flattened halo and a bulk rotational component has suggested that a very highly condensed model for the MACHO distribution, such as a thick disk, might reduce the MACHO mass estimate from the current data to a level consistent with brown dwarf candidates.  Previous explorations \\cite{us_nomacho} have indicated that such models may be able to reproduce the observed optical depths toward the Large Magellanic Cloud (LMC) and the Galactic bulge.  The overall shape of the Galactic halo is unknown.  Attempts to determine the shape of galactic halo potentials from flaring of the outer Galactic gas layer \\cite{Olling,Sackett} point to a flattened halo; flattening of the potential is also supported by simulations of the cold dark matter halo formation \\cite{nbody}.  While it is highly unlikely that the entire halo is in a disk-like configuration, it is not unreasonable to assume that the MACHO component of the halo is significantly more flattened than the dark matter halo.  The condensation of a gaseous halo component to form a very thick disk is likely to result in significant star formation and thus the production of a population of lens candidates.  In this paper we explore in detail the consequences of such a lens population.  We first define the density and velocity structure, and outline the constraints on these distributions.  We then present predictions and observational consequences of such thick disks; in particular we determine the expected frequency of parallax events. ",
        "conclusions": "Microlensing studies have yielded much exciting data in the past few years and are continuing to survey different lines of sight through the Galaxy in order to probe the Galactic halo.  However, the conclusions that can be drawn from the data to date are very model dependent -- assumptions about the distribution of the lenses and their velocity structure have a strong impact on their interpretation. Thus we need to examine a wide range of reasonable lens distributions.  Very thick disks present a reasonable alternative to a halo population of lenses.  If the lenses are stellar remnants, it seems likely that their configuration will be more condensed than that of a standard non-baryonic halo.  While we have found that very thick disks cannot lower the lens mass estimate to the brown dwarf regime, they have the advantage that their total mass in MACHOs is somewhat less than that for a standard halo that is truncated at $50\\kpc$ (and much less than a MACHO halo which traces the extended dark halo out to at least $100\\kpc$.)  A thick disk distribution cannot produce an optical depth toward the LMC of more than about $1.5 \\times 10^{-7}$.  However, it can explain (within the experimental uncertainties) all or a significant fraction of the current optical depth estimates.  As more events are detected, it may be possible to distinguish between very thick disk and halo lens distributions.  The most promising avenue for such a discriminant is the observation of parallax events.  Because disk lenses would be both closer and on average slower than halo lenses we expect a higher rate of such events for a disk population. Although the survey experiment light curve measurements can only marginally discriminate between disk and halo distributions for reasonable numbers of events, a modest expenditure of telescope time to obtain one percent photometry on all the Magellanic cloud events should be capable of making the distinction very clearly."
    },
    "9803/hep-ph9803479_arXiv.txt": {
        "abstract": "Contrary to naive cosmological expectations, all evidence suggests that the universe contains an abundance of matter over antimatter. This article reviews the currently popular scenario in which testable physics, present in the standard model of electroweak interactions and its modest extensions, is responsible for this fundamental cosmological datum. A pedagogical explanation of the motivations and physics behind electroweak baryogenesis is provided, and analytical approaches, numerical studies, up to date developments and open questions in the field are also discussed. ",
        "introduction": "\\label{Intro}  The most basic distinction drawn between particles found in nature is that between particle and antiparticle. Since antiparticles were first predicted \\cite{Dirac1,Dirac2} and observed \\cite{Anderson1,Anderson2}, it has been clear that there is a high degree of symmetry between particle and antiparticle. This means, among other things, that a world composed of antimatter would behave in a similar manner to our world. This basic tenet of particle physics, the symmetry between matter and antimatter, is in stark contradiction to the wealth of everyday and cosmological evidence that the universe is composed almost entirely of matter with little or no primordial antimatter.  The evidence that the universe is devoid of antimatter comes from a variety of different observations. On the very small scale, the absence of proton-antiproton annihilations in our everyday actions, constitutes strong evidence that our world is composed only of matter and no antimatter. Moving up in scale, the success of satellite launches, lunar landings, and planetary probes indicate that our solar system is made up of the same type of matter that we are, and that there is negligible antimatter on that scale. To determine whether antimatter exists in our galaxy, we turn to cosmic rays. Here we see the first detection of antimatter outside particle accelerators. Mixed in with the many protons present in cosmic rays are a few antiprotons, present at a level of around $10^{-4}$ in comparison with the number of protons (see for example \\citeasnoun{SA 88}). However, this number of antiprotons is consistent with their secondary production through accelerator-like processes, $p+p\\rightarrow 3p + {\\bar p}$, as the cosmic rays stream towards us. Thus there is no evidence for primordial antimatter in our galaxy. Finally, if matter and antimatter galaxies were to coexist in clusters of galaxies, then we would expect there to be a detectable background of $\\gamma$-radiation from nucleon-antinucleon annihilations within the clusters. This background is not observed and so we conclude that there is negligible antimatter on the scale of clusters (For a review of the evidence for a baryon asymmetry see \\citeasnoun{GS 76}).  A natural question to ask is, what is the largest scale on which we can say that there is no antimatter? This question was addressed by \\citeasnoun{GS 76},and by \\citeasnoun{FS 85}, and in particular has been the subject of a careful recent analysis by \\citeasnoun{CRG 97}. If large domains of matter and antimatter exist, then annihilations would take place at the interfaces between them. If the typical size of such a domain was small enough, then the energy released by these annihilations would result in a diffuse $\\gamma$-ray background and a distortion of the cosmic microwave radiation, neither of which is observed. Quantitatively, the result obtained by the latter authors is that we may safely conclude that the universe consists entirely of matter on all scales up to the Hubble size. It therefore seems that the universe is fundamentally matter-antimatter asymmetric.  The above observations put an experimental upper bound on the amount of antimatter in the universe. However, strict quantitative estimates of the relative abundances of baryonic matter and antimatter may also be obtained from the standard cosmology. Primordial nucleosynthesis (for a review see \\citeasnoun{CST 95}) is one of the most powerful tools of the standard cosmological model. The theory allows accurate predictions of the cosmological abundances of all the light elements, H, $^3$He, $^4$He, D, B and $^7$Li, while requiring only a single input parameter. Define $n_b$ to be the number density of baryons in the universe. Similarly define $n_{\\bar b}$ to be the number density of antibaryons, and the difference between the two to be $n_B$. Then, if the entropy density in the universe is given by $s$, the single parameter required by nucleosynthesis is the baryon to entropy ratio  \\be \\eta \\equiv \\frac{n_B}{s} = \\frac{n_b-n_{\\bar b}}{s} \\ , \\ee and one may conservatively say that calculations of the primordial light element abundances are correct if  \\be 1.5\\times 10^{-10} < \\eta < 7\\times 10^{-10} \\ . \\label{nucleo} \\ee Although the range of $\\eta$ within which all light element abundances agree with observations is quite narrow (see figure \\ref{bbnfig}), its existence at all is remarkable, and constitutes a strong confirmation of the standard cosmology. For recent progress in nucleosynthesis see \\citeasnoun{TFB 96} and \\citeasnoun{CH 97}.  The standard cosmological model provides a complete and accurate description of the evolution of the universe from extremely early times (a few minutes) to the present day ($10-20$ billion years) given a host of initial conditions, one of which is the value of $\\eta$. This standard picture is based on classical, fluid sources for the Einstein equations of General Relativity (GR). While the success of the standard cosmology is encouraging, there remains the question of the initial conditions. One approach is just to consider the initial values of cosmological parameters as given. However, the values required for many parameters are extremely unnatural in the sense that the tiniest deviation from them leads to a universe entirely different from the one that we observe. One well known example of this is the initial value of the mass density of the universe, the naturalness of which is at the root of the {\\it flatness problem} of the standard cosmology.  The philosophy of {\\it modern} cosmology, developed over the last thirty years, is to attempt to explain the required initial conditions on the basis of quantum field theories of elementary particles in the early universe. This approach has allowed us to push our understanding of early universe cosmology back to much earlier times, conservatively as early as $10^{-10}$ seconds, and perhaps much earlier.  The generation of the observed value of $\\eta$ in this context is referred to as {\\it baryogenesis}. A first step is to outline the necessary properties a particle physics theory must possess. These conditions were first identified by \\citeasnoun{Sakharov} and are now referred to as the three {\\it Sakharov Criteria}. They are:  \\begin{itemize}  \\item Violation of the baryon number ($B$) symmetry.  \\item Violation of the discrete symmetries $C$ (charge conjugation) and $CP$ (the composition of parity and $C$) \\item A departure from thermal equilibrium.  \\end{itemize}  The first of these is rather obvious. If no processes ever occur in which $B$ is violated, then the total number of baryons in the universe must remain constant, and therefore no asymmetry can be generated from symmetric initial conditions. The second Sakharov criterion is required because, if $C$ and $CP$ are exact symmetries, then one can prove that the total rate for any process which produces an excess of baryons is equal to the rate of the complementary process which produces an excess of antibaryons and so no net baryon number can be created. That is to say that the thermal average of $B$, which is odd under both $C$ and $CP$, is zero unless those discrete symmetries are violated. Finally, there are many ways to explain the third criterion. One way is to calculate the equilibrium average of $B$:  \\bea \\langle B\\rangle_T & = & \\hbox{Tr}(e^{-\\beta H}B) \\nonumber \\\\ & = & \\hbox{Tr}[(CPT)(CPT)^{-1}e^{-\\beta H}B)] \\nonumber \\\\ & = & \\hbox{Tr}(e^{-\\beta H}(CPT)^{-1}B(CPT)] \\nonumber \\\\ & = & -\\hbox{Tr}(e^{-\\beta H}B) \\ , \\eea where in the third step I have used the requirement that the Hamiltonian $H$ commutes with $CPT$, and in the last step used the properties of $B$ that it is odd under $C$ and even under $P$ and $T$ symmetries. Thus $\\langle B\\rangle_T = 0$ in equilibrium and there is no generation of net baryon number. This may be loosely described in the following way. In quantum field theories in thermal equilibrium, the number density of any particle species, $X$ say, depends only on the energy of that species, through  \\be n_{eq}(X)=\\frac{1}{e^{(E-\\mu)/T} \\pm 1}\\ , \\ee where $\\mu$ is the chemical potential corresponding to baryon number. Since the masses of particle and antiparticle are equal by virtue of the CPT theorem, and $\\mu=0$ if baryon number is violated, we have that  \\be N_{eq}(X)=\\int\\frac{d^3p}{(2\\pi)^3}n_{eq}=N_{eq}({\\bar X})\\ , \\ee and again there is no net asymmetry produced.  The focus of this article is to review one popular scenario for generating the baryon asymmetry of the universe (BAU), as quantified in equation~(\\ref{nucleo}), within the context of modern cosmology. In general, such scenarios involve calculating $n_B$, and then dividing by the entropy density  \\be s=\\frac{2\\pi^2}{45}g_* T^3 \\ , \\ee where $g_*$ is the effective number of massless degrees of freedom at temperature $T$. While there exist many attempts in the literature to explain the BAU (for a review see \\citeasnoun{AD 92}), I will concentrate on those scenarios which involve anomalous electroweak physics, when the universe was at a temperature of $10^2\\,$GeV (for earlier reviews see \\citeasnoun{NTreview}, \\citeasnoun{CKNreview}, and \\citeasnoun{review}). The production of the BAU through these models is referred to as {\\it electroweak baryogenesis}.  In the next section I will describe baryon number violation in the electroweak theory both at zero and at nonzero temperature. In section \\ref{CP} I shall move on to the subject of CP violation, explaining how this arises in the standard model and how it is achieved in some popular extensions. Section \\ref{EWPT} contains an account of the electroweak phase transition, including a discussion of both analytic and numerical approaches. Having set up the framework for electroweak baryogenesis, I turn in section \\ref{local} to the dynamics in the case where baryon production occurs close to a phase boundary during a phase transition. In section \\ref{nonlocal}, I then extend these ideas to include the effects of particle transport, or diffusion. Section \\ref{MSSM} contains a description of how baryogenesis is implemented in a popular extension of the standard model, the minimal supersymmetric standard model (MSSM). In section \\ref{defects}, I explain how, in some extensions of the electroweak theory, baryogenesis may be mediated by topological defects, alleviating the constraints on the order of the phase transition. Finally, in section \\ref{conclusions} I summarize the results and comment on open questions and future directions in the field.  It is my hope that this article will fulfill its intended role as both a review of the background and basic material for beginners in the field, and as a summary of and commentary on the most recent results and directions in the subject. However, the focus of this article, as in any such endeavor, is quite idiosyncratic, and I apologize to any of my colleagues whose work has been omitted or incorrectly detailed. A different focus can be found in other accounts of the subject and, in particular, for a comprehensive modern review of numerical approaches I recommend \\citeasnoun{review}.  A note about conventions. Throughout I use a metric with signature $+2$ and, unless explicitly stated otherwise, I employ units such that ${\\hbar}=c=k=1$ so that Newton's constant is related to the Planck mass through $G=M_{pl}^{-2}$. ",
        "conclusions": "\\label{conclusions} Modern particle cosmology consists of the union of the hot big bang model with quantum field theories of elementary particles. Under mild assumptions, it is a consequence of this structure that when the universe was extremely young and hot, the net baryon number of the universe was zero. That is, the number of particles carrying a given baryon number in any region was equal on average to the number of the appropriate antiparticles carrying the opposite baryon number. However, on the other hand, it is a clear observational fact that the universe is maximally baryon - antibaryon asymmetric. This fact is quantified from the considerations of primordial nucleosynthesis. These calculations are perhaps the most impressive success of the standard cosmology, accurately predicting the abundances of the light elements from the single input parameter of the baryon to entropy ratio, which is constrained as in equation~(\\ref{nucleo}). Until recently, there were two possible explanations for this. First, the universe as a whole could be baryon number symmetric, but baryon-antibaryon separation could have resulted in an apparent baryon asymmetry in the local universe. the second possibility is that some dynamical process took place as the universe evolved, causing baryons to be preferentially produced over antibaryons. A recent analysis \\cite{CRG 97} has ruled out the former possibility over scales up to the size of the observable universe. It therefore appears that the latter option, {\\it baryogenesis}, must have taken place.  In this article I have tried to describe how a number of different physical effects, all present in the standard electroweak theory at nonzero temperature, can come together in the context of the expanding universe to implement baryogenesis. It is an amazing fact about the Glashow-Salam-Weinberg model and its modest extensions that they satisfy all three Sakharov criteria for producing a baryon excess. As a result, over the last decade, {\\it electroweak baryogenesis} has been a very popular scenario for the generation of the BAU.  There are four technical issues to be investigated when considering models of electroweak baryogenesis. these are  \\begin{enumerate} \\item How is the departure from equilibrium realized? Is the electroweak phase transition strongly first order or are topological defects necessary?  \\item How is sufficient CP violation obtained?  \\item What is the rate of baryon number violating processes? Is this rate high enough in the unbroken phase and can washout of any asymmetry be avoided?  \\item What are the actual dynamics of baryon number production?  \\end{enumerate}  I hope I have described how each of these issues has been addressed in the literature on electroweak baryogenesis. Detailed analyses of the phase transition, coupled with the smallness of the CP violation due to phases in the CKM matrix, have made it clear that an extension of the standard model is required to make the scenarios viable. While general two-Higgs doublet theories have been considered, perhaps the most appealing candidate from particle physics motivations is the minimal supersymmetric standard model. The implementations of electroweak baryogenesis in this model have recently been investigated and the range of parameters of the model for which an appreciable BAU can be generated have been calculated. This range should be accessible to future particle colliders and thus the scenario of EWBG in the MSSM should be experimentally testable. If the MSSM or two-Higgs models are not chosen by nature, then other models of electroweak baryogenesis may be relevant. If the topological structure of the relevant theory admits gauge solitons, then defect mediated electroweak baryogenesis may be important, irrespective of the order of the phase transition.  Although I have argued that there exist viable scenarios of electroweak baryogenesis, particularly in the context of supersymmetric models, there remain a number or open questions and directions for future research. I shall list the ones that I feel are most important.  \\begin{enumerate} \\item At present, the best quantitative understanding of the electroweak phase transition comes from lattice Monte Carlo simulations. While the results from these are impressive, they do not provide an intuitive understanding of the microphysics of the phase transition. The numerical results are partially supported by some analytical approaches but these often cannot be trusted in the physical range of Higgs masses. An analytic understanding of the nonperturbative dynamics of the phase transition would be an important step forward.  \\item The chemical potential and ``fluid'' approaches to nonlocal baryogenesis provide good analytical tools for understanding the nonlocal production of baryons. However, in the case of local baryogenesis, analytical methods which yield believable quantitative results have yet to be found. Again, it is encouraging that numerical calculations, for example of Chern-Simons number diffusion, are providing quantitative predictions but an appropriate analytical model is very desirable.  \\item While it will strongly support electroweak baryogenesis if supersymmetry is verified with parameters in the correct range, it is important to ensure that there is sufficient CP violation. To this end, phenomenological predictions of CP violation in SUSY models and the corresponding experimental tests are crucial hurdles that the scenario must pass.  \\end{enumerate}  If electroweak baryogenesis is correct, then no matter what the relevant electroweak model is, the physics involved is extremely beautiful and diverse. The topology of gauge field theories, the physics of phase transitions, CP violation, plasma dynamics and thermal field theory all play a part in generating the BAU. However, perhaps the most attractive feature of electroweak baryogenesis scenarios is that they should be testable at the next generation of particle colliders. It is an exciting possibility that we may soon understand the origin of the matter that makes up our universe."
    },
    "9803/gr-qc9803084_arXiv.txt": {
        "abstract": "We develop the general theory of stars in Saa's model of gravity with propagating torsion and study the basic stationary state of neutron star. Our numerical results show that the torsion force decreases the role of the gravity in the star configuration leading to significant changes in the neutron star masses depending on the equation  of state of star matter. The inconsistency of the Saa's model with Roll-Krotkov-Dicke and Braginsky-Panov experiments is discussed.  \\noindent{PACS number(s): 04.40.Dg,04.40.-b,04.50.+h} ",
        "introduction": "In recent years the interest in scalar-tensor theories of gravity has been renewed. One reason for this is the important role which these theories play in the understanding of inflantionary epoch. On the other hand the scalar-tensor gravitation (the so called \"dilaton gravity\") arises naturally from the low-energy limit of the super-string theory \\cite{Witten}, \\cite{DT}.  The predictions of scalar-tensor theories may differ drastically from these of general relativity. For example such a phenomenon -- \"spontaneous scalarization\" was recently discovered by Damour and Esposito-Farese as a non-perturbative strong field effect in a massive neutron star \\cite{Damour}. Other interesting phenomenon is \"gravitational memory\" of black holes proposed in \\cite{Barrow}. The \"gravitational memory\" in the case of boson stars was investigated in \\cite{TLS} (see also \\cite{CS}).Their stability through cosmic history using catastrophe theory was investigated in \\cite{TSL}.  Many theories of gravity with propagating torsion involving a scalar field have been proposed in the last decades, too \\cite{HRRS}, \\cite{HRR}, \\cite{SG}. In such theories contrary to the usual Einstein-Cartan gravity \\cite{Hehl1}-\\cite{Hehl3}, there are long-range torsion mediated interactions. Carrol and Field \\cite{Carrol} have examined some observational consequences of propagating torsion in a wide class of models involving a scalar field. They conclude that for reasonable models the torsion could be detected experimentally.  Recently a new interesting model with propagating torsion was proposed by Saa \\cite{Saa1}-\\cite{Saa5}. This model involves a non-minimally coupled scalar field as a potential of the torsion of space-time. As one can see Saa's model is very close to the dilaton gravity.  In the present article we investigate both analytically and numerically a neutron star in the Saa's model and compare obtained results with these in the general relativity. We also discuss new predictions of the theory under consideration.  The paper is organized as follows. In section 2 we consider briefly Saa's model. In section 3 we give the necessary information for the vacuum solutions of the field equations. The equations determining static equilibrium solutions for a neutron star are discussed in section 4. Numerical results for the neutron star are discussed in section 5. The stability of the neutron star is discussed via catastrophe theory in section 6. The inconsistency of the Saa's model with Roll-Krotkov-Dicke and Braginsky-Panov experiments is discussed in section 7. ",
        "conclusions": "We have solved the system of equations (\\ref{NF}) coupled with the state equations (\\ref{EOS1}) and (\\ref{EOS2}) numerically  using the method due to Runge-Kutta-Merson with automatic error control. The results are shown in the corresponding figures.\\\\ Hereafter all masses are measured in units $M_{\\bigodot}$. \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz01a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz01b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{a) $M -log(\\varepsilon_{c})$ dependence. \\hskip 1.8truecm b) $M_{T} - M_{R}$ dependence.\\hskip .7truecm} \\vspace{.5truecm} \\label{Fig1} \\end{figure} First we concentrate our attention on the case of non-interacting neutron gas. In Fig. 1a) the dependence of the three masses $M_{T}$,$M_{K}$,$M_{R}$ on the central density $\\varepsilon_{c}$ is shown. The appearance of a cusp in Fig. 1b) , where the dependence of the tensor mass $M_{T}$ on the rest mass $M_{R}$ is presented, shows that their maxima lie at the same point. Although, the maxima of the rest mass $M_{R}$ and Keplerian mass $M_{K}$ are too close in Fig. 1a) they don't lie at the same point, as it may be seen from Fig. 2a) which shows the dependence of $M_{K}$ on $M_{R}$. We see also that Keplerian mass is considerably greater than the tensor one -- about three times.   In Fig. 3a) the $M_{T} - R$ dependence is represented. It's seen that the $M_{T} - R$  curve in our case is fairly similar to the one of general relativity, but there are significant differences, too. The maximum mass ${M_{T}}_{max} $ in our case is $ \\approx 0.35M_{\\bigodot}$, while in general relativity the Oppenheimer-Volkoff's mass is $M_{OV} = 0.7M_{\\bigodot}$. The radius corresponding to the mass $M_{\\bigodot}$ is  $R = 4.2 km $, while in the case of general relativity  $ R = 9.6 km $. If we look at Fig. 2b) where the dependence of $M_{T}$ on the central density ${\\varepsilon}_{c}$ is shown, we note that ${M_{T}}_{max}$ lies at ${\\varepsilon}_{c} \\approx 4.5 * 10^{16} g/{cm}^3$, while  $M_{OV}$ lies at ${\\varepsilon}_{c} \\approx  5 * 10^{15} g/{cm}^3$ in general relativity. The average  density in our case is about $ 4 $ times greater than the one in general relativity. Hence, in the model under consideration the neutron star is more compact and has a mass about $1/2$ - times smaller than $M_{OV}$. In Fig. 3b) the dependence of Keplerian  mass on the star radius is presented. It's seen that the Keplerian mass is about $1.5$ times greater than $M_{OV}$. \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz02a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz02b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{a) $M_{Kepler}-M_{R}$ dependence. \\hskip 2truecm b) $M_{T} -log(\\varepsilon_{c})$ dependence.\\hskip .5truecm} \\vspace{.5truecm} \\label{Fig2} \\end{figure}  \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz03a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz03b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{a) $M_{T}-R$ dependence. \\hskip 2.4truecm b) $M_{Kepler}-R$ dependence.\\hskip  1.1truecm } \\vspace{.5truecm} \\label{Fig3} \\end{figure}  \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz04a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz04b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{a) $k- r$ dependence. \\hskip 2truecm b) $K -log(\\varepsilon_{c})$ dependence.\\hskip  .5truecm} \\vspace{.5truecm} \\label{Fig2_} \\end{figure}  In the  Fig. 4a) the dependence $k(r)$ is shown inside the star (for central density $7.5*10^{15}g/{cm}^3 $). In accordance to the general considerations $k$ increases from the center of the star to the surface, where $k$ takes a value $K=k(R) \\approx 0.45 - 0.46$, which is close to $0.5$. The dependence $K(\\varepsilon_{c})$ of $K$ on the star central density $\\varepsilon_{c}$ is shown in Fig. 4b). It's seen that $K$ decreases when density increases, which is similar to the previous case. So, the ratio of the torsion force to the gravitational one takes its minimum value at the center of the star and is the greatest at the surface. \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz05a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz05b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{a) $K-M_{T}$ dependence. \\hskip 2.4truecm b) $K-R$ dependence.\\hskip  1.1truecm } \\vspace{.5truecm} \\label{Fig3_} \\end{figure}  As it may be seen from Fig. 5a) expressing the dependence $K(M_{T})$, the torsion-urged effects are relatively strongest in the case of small masses -- with increasing of the star mass (up to the point where the star loses its stability) $K$ decreases. It's seen from Fig. 5b), where the dependence of $K$ on the star radius $R$ is shown, that in the area of stability  $K$ decreases when $R$ decreases too -- the more compact stars are, the smaller $K$ they have.  Fig. 6a) presents the dependencies $\\theta(r)$ and $\\nu(r)$ inside the star. One may see that ${1 \\over 2}\\nu - 3\\Theta < 0$  everywhere. The dependencies $m_{T}(r)$,$m_{Kepler}(r)$ and $m_{\\theta}(r)$ are shown in Fig. 6b)  for central density $7.5*10^{15}$. As it has already mentioned all masses  increase with $r$. \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz06a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz06b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{a) $\\Theta, \\nu -r$ dependence. \\hskip 1.9truecm b)$m_{T},m_{Kepler},m_{\\theta} -r$ dependence.\\hskip .7truecm } \\vspace{.5truecm} \\label{Fig4} \\end{figure}  The following figures illustrate the case of Tsuruta-Cameron equation of state (TCES). \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz07a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz07b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{TCES. a) $M -log(\\varepsilon_{c})$ dependence . \\hskip 1.2truecm b) $M_{T} -M_{R}$ dependence.\\hskip  1.2truecm } \\vspace{.5truecm} \\label{Fig5} \\end{figure}  We see from the figures that the maximum tensor mass in this case is about $2M_{\\bigodot}$ and the corresponding radius is about $7.5km$ - the same quantities in general relativity are correspondingly $\\approx 1.6M_{\\bigodot}$ and $\\approx 11.5km$. Hence, the interaction between the nucleons leads to an increase in the maximum mass, as in general relativity.  Note the differences between the Fig. 5b) and Fig. 6b) (for the case of non-interacting neutron gas), and the corresponding  Fig. 11b) and Fig. 12b) (for the case of Tsuruta-Cameron equation of state). There one can see the strong dependence of some results in the Saa's model of gravity with propagating torsion on the equation of state of star's matter.  \\vskip 2truecm  \\begin{figure}[htbp] \\vspace{3.5truecm} \\special{psfile=tfiz08a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz08b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{TCES. a)$M_{Kepler}-M_{R}$ dependence. \\hskip .4truecm b) $M_{T}-log(\\varepsilon_{c}) $ dependence.\\hskip 1truecm} \\vspace{.5truecm} \\label{Fig6} \\end{figure}  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz09a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz09b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{TCES. a) $M_{T}-R$ dependence. \\hskip 1.6truecm b) $M_{Kepler}-R$ dependence.\\hskip  1.7truecm} \\vspace{.5truecm} \\label{Fig7} \\end{figure}   \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz10a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz10b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{TCES. a) $k-r$ dependence. \\hskip 1.3truecm b) $K -log(\\varepsilon_{c}) $ dependence.\\hskip 3truecm} \\vspace{.5truecm} \\label{Fig8} \\end{figure}  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz11a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz11b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{TCES. a) $K-M_{T}$ dependence. \\hskip 1.3truecm b) $K-R $ dependence.\\hskip 3truecm} \\vspace{.5truecm} \\label{Fig8_} \\end{figure}  \\begin{figure}[htbp] \\vspace{4truecm} \\special{psfile=tfiz12a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \\special{psfile=tfiz12b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \\vskip 0.5truecm \\caption{TCES. a) $\\Theta, \\nu-r$ dependence. \\hskip 1.3truecm b) $m_{T},m_{Kepler}, m_{\\theta}-r$ dependence.\\hskip 3truecm} \\vspace{.5truecm} \\label{Fig8+} \\end{figure} We have also examined the Harrison-Wheeler's equation of state \\cite{HTWW}. As in general relativity the numerical results are very close to these for the noninteracting neutron gas. For example the maximum tensor and  Keplerian mass is correspondingly $\\approx 0.35M_{\\bigodot}$ and $\\approx 1M_{\\bigodot}$, and the corresponding radius is $3.8 km$.  Other equations of state (of politropic type) have been examined, too. The corresponding maximum tensor mass of a neutron star reaches a value about $2.5-2.6M_{\\bigodot}$, while the corresponding Keplerian mass is about $6-6.5M_{\\bigodot}$.  As it is seen from numerical calculations the tensor mass and Keplerian mass differ very significantly from each other. This behaviour of the  masses is qualitatively the same as in the case of a boson star in Brans-Dicke theory with $\\omega=-1$  \\cite{Whinnett}. Saa's model corresponds to the value of the Brans-Dicke parameter $\\omega=-{4\\over 3}$ which is close to $-1$. That's why the observed qualitative agreement is natural.  Hence, including a scalar field with an approximately the same $\\omega$ in physically different kinds of stars  we find similar departures of the corresponding predictions of general relativity.  This conclusion holds also in the case of large $\\omega$. For example, in \\cite{TLS} and \\cite{TSL} Brans-Dicke boson star has been considered with $\\omega = 400$. There the results are quite similar to corresponding ones in general relativity which is just the limit $\\omega \\to \\infty$. The star is a bit lighter but has a higher density in future.    In this article we have examined the basic spherically symmetric stationary state of stars in the Saa's model of gravity with propagating torsion.  In the model under investigation there is no need to consider unknown charges creating the torsion-dilaton field. Its source is the very spinless matter. The whole geometry of the space-time (including metric and torsion) is determined by the familiar properties of this matter.  The parameters of the vacuum solution are determined only by the spinless matter without adopting an existence of new properties, too. In contrast to the corresponding models in the general relativity here we have two parameters $K$ and $a$ of the vacuum solutions. The values of these parameters depend on the mass distribution in the star which is related with the equation of state of the star's matter. For a fixed equation of state both parameters become functions only of the star mass, but these functions are not the same for the different equations of state. The first parameter $K$ being the ratio of the magnitude of torsion-dilaton force and of the magnitude of gravitational force for realistic equations of matter state takes values in the interval $[{1\\over 3},{1\\over 2}]$, depending on the star's mass. The second one -- $a$ is analogous to the gravitational radius in general relativity and takes positive values depending on the value of the parameter $K$ and on the value of the star's mass.  To be specific in the present article we restrict our attention to the model of neutron stars where the effects of nonlinearity are essential as in general relativity. Numerical results and analytical considerations show that the space-time torsion may have a significant role in their structure. The new torsion force decrease in some extent the role of the gravity in the star configuration and may lead to an increasing or decreasing of the maximum neutron star mass depending on the equation of state.   The complete investigation of the consistence of the whole Saa's model of gravity with propagating torsion (including all type of physical fields) with the reality is still an open problem. The results of the present article may have not only independent value, but are necessary for reaching the solution of this critical problem. For example, after the first version of the present article was send for publication a new results based on it which show the inconsistency of Saa's model with solar system gravitational experiments were found and published independently \\cite{FY}.  Saa's model is a simple model based on pure geometrical reasons which allows to overcome a basic difficulties of the old models of gravity with torsion: the inconsistency of the application of the minimal coupling principle in action principle and directly in the equations of motion \\cite{Saa1}. Moreover, it leads to propagating torsion which is another important physical property of this model. Unfortunately, it is not compatible with basic physical experiments. This means that one has to modify this model preserving its important new physical properties in a proper way to comply with real physics. A new investigations in this direction are in progress (see for example \\cite{F2}, \\cite{F3}).  Nevertheless the model under consideration contradicts to basic experiments its detailed investigation is quite instructive because Saa's model is a special case of scalar-tensor theory of gravity with a nonminimal coupling of the scalar torsion-dilaton field with matter. A similar, but more general coupling of string-dilaton can be expected in the string theory and it leads different physical effects. For example, in the article by Damour and Polyakov \\cite{DP} as a consequences of nonminimal coupling between matter and dilaton was derived a violation of the equivalence principle, as far as some string-dilaton effects in the early universe \\cite{DP}. Unfortunately, the present days string theory is not able to predict definitely the form of the coupling between string-dilaton and the real matter. In contrast, Saa's model is the only one we know with complete determined interaction of the (torsion) dilaton with all kinds of matter. This interaction is similar to the one expected in other models. In the present article it is shown at first that the nonminimal coupling of the dilaton field with the matter will emerge in extreme conditions in a neutron star and that it leads to clear new physical phenomena of different kind, some details of which may depend on the equation of matter state. Most probably similar effects will appear in other possible modifications of the theory of torsion-dilaton, as far as in the other types of theories of dilaton. For example the neutron stars may give us a way to a real physics in string theories if the string-dilaton interactions with real matter will be established. Hence, the most important conclusion of present article is that looking for physical manifestation of the dilaton one has to investigate in details its influence on the neutron star structure.    \\bigskip \\bigskip \\bigskip  \\noindent{\\Large\\bf Acknowledgments}  \\bigskip \\bigskip \\bigskip We are deeply grateful to the unknown referees who suggested to consider different types of masses of the neutron star in Saa's model and stability analysis of the neutron star via catastrophe theory, and to add to the present paper results about the consistence of this model with physical reality, as far as for pointing out the references \\cite{DT}, \\cite{TLS}, \\cite{CS}, \\cite{TSL}, \\cite{HRR}, \\cite{SG}, \\cite{Whinnett}, \\cite{KMS1} and \\cite{KMS2}. The work on this article has been partially supported by the Sofia University Foundation for Scientific Researches, Contracts~No.No.~245/97,~257/97, and by the Bulgarian National Foundation for Scientific Researches, Contract~No.~F610/97. One of us (PF) is grateful to the leadership of the Bogoliubov Laboratory of Theoretical Physics, JINR, Dubna, Russia for hospitality and working conditions during his stay there in the summer of 1998 when a part of this investigation has been completed.   \\bigskip \\bigskip \\bigskip"
    },
    "9803/astro-ph9803185_arXiv.txt": {
        "abstract": "We report the results of a systematic study of the vacuum- ultraviolet ($\\lambda \\simeq$ 1150 to 2000 \\AA) spectra of a sample of 45 starburst and related galaxies observed with the IUE satellite. These span broad ranges in metallicity (from 0.02 to 3 times solar), bolometric luminosity ($\\sim 10^{7}$ to $4 \\times 10^{11} L_{\\odot}$), and galaxy properties (e.g. including low-mass dwarf galaxies, normal disk galaxies, and massive galactic mergers). The projected size of the IUE spectroscopic aperture is typically one to several kpc and therefore usually encompasses the entire starburst and is similar to the aperture-sizes used for spectroscopy of high-redshift galaxies. Our principal conclusion is that local starbursts occupy a very small fractional volume in the multi- dimensional manifold defined by such fundamental parameters as the extinction, metallicity, and vacuum-UV line strengths (both stellar and interstellar) of the starburst and the rotation speed (mass) and absolute magnitude of the starburst's `host' galaxy. More metal-rich starbursts are redder and more heavily extinguished in the UV, more luminous, have stronger vacuum- UV lines, and occur in more massive and optically-brighter host galaxies. We advocate using these local starbursts as a `training set' to learn how to better interpret the rest-frame UV spectra of star-forming galaxies at high-redshift, and stress that the degree of similarity between local starbursts and high-redshift galaxies in this multi-dimensional parameter space can already be tested empirically. The results on local starbursts suggest that the high- redshift `Lyman Drop-Out' galaxies are typically highly reddened and extinguished by dust (average factor of 5 to 10 in the UV), may have moderately high metallicities (0.1 to 1 times solar?), are probably building galaxies with stellar surface-mass-densities similar to present-day ellipticals, and may be suffering substantial losses of metal-enriched gas that can `pollute' the inter-galactic medium. ",
        "introduction": "Starbursts are sites of intense star-formation that occur in the `circum-nuclear' (kpc-scale) regions of galaxies, and dominate the integrated emission from the `host' galaxy (cf. Leitherer et al 1991). The implied star-formation rates are so high that the existing gas supply may sustain the starburst for only a small fraction of a Hubble time (in agreement with detailed models of the observed properties of starbursts, which imply typical burst ages of-order 10$^{8}$ years). Both optical objective prism searches and the IRAS survey have shown that starbursts are major components of the local universe (cf. Huchra 1977; Gallego et al et al 1995; Soifer et al 1987). Indeed, integrated over the local universe, the total rate of (high-mass) star-formation in circumnuclear starbursts is comparable to the rate in the disks of spiral galaxies (Heckman 1997; Gallego et al 1995; Tresse \\& Maddux 1998). Thus, starbursts deserve to be understood in their own right.  Starbursts are even more important when placed in the broader context of contemporary stellar and extragalactic astrophysics. The cosmological relevance of starbursts has been dramatically underscored by one of the most spectacular discoveries in years: the existence of a population of high-redshift (z $>$ 2) star-forming field galaxies (cf. Steidel et al 1996; Lowenthal et al 1997). The sheer number density of these galaxies implies that they almost certainly represent precursors of typical present-day galaxies in an early actively-star-forming phase. This discovery therefore moves the study of the star-forming history of the universe into the arena of direct observations (Madau et al 1996), and gives added impetus to the quest to understand local starbursts.  Observations in the vacuum-UV spectral regime are crucial for both understanding local starbursts, and for relating them to galaxies at high-redshift. This is the spectral regime where we can most clearly observe the direct spectroscopic signatures of the hot stars that provide most of the bolometric luminosity of starbursts (e.g. Sekiguchi \\& Anderson 1987; Fanelli, O'Connell, \\& Thuan 1988; Leitherer, Robert, \\& Heckman 1995). Moreover, the vacuum-UV contains a wealth of spectral features, including the resonance transitions of most cosmically-abundant ionic species (cf. Morton 1991). These give UV spectroscopy a unique capability for diagnosing the (hot) stellar population and the physical, chemical, and dynamical state of gas in starbursts.  Since ground-based optical observations of galaxies at high-redshifts sample the vacuum-UV portion of their rest-frame spectrum, we cannot understand how galaxies evolved without documenting the vacuum-UV properties of galaxies in the present epoch. In particular, a thorough understanding of how to exploit the diagnostic power of the rest-frame UV spectral properties of local starbursts will give astronomers powerful tools with which to study star-formation and galaxy-evolution in the early universe.  Accordingly, we have undertaken an analysis of the vacuum-UV spectroscopic properties of a sample of 45 starburst and related galaxies in the local universe using the IUE archives. In section 2, we describe our analysis of these data. In section 3 we use the data to document the empirical vacuum-UV spectroscopic properties of local starbursts and to assess the dependence of these UV properties on crucial parameters like the dust content, metallicity, and luminosity of the starburst, and the mass and luminosity of the `host galaxy'. We will also describe how published analyses of HST UV spectra allow us to better understand the empirical results from the IUE data, and in particular to ascertain the origin of the absorption features seen in the IUE data (stellar vs. interstellar). In section 4, we will interpret the results from section 3, use these results to elucidate some of the properties of starbursts at high redshift, and point-out some potentially far-reaching implications. ",
        "conclusions": "The principal conclusion of this paper is that local starbursts occupy a very small fractional volume in the multi-dimensional manifold defined by such fundamental parameters as the extinction, metallicity, and vacuum-UV line strengths (both stellar and interstellar) of the starburst and the rotation speed (mass) and absolute magnitude of the starburst's `host' galaxy. More metal-rich starbursts are redder, more heavily extinguished (IR-dominated), have stronger vacuum-UV lines, occur in more massive and brighter host galaxies, and can be more luminous.  Starbursts with solar or higher metallicity are heavily extinguished indeed: only 1 to 10\\% of the vacuum-UV radiation escapes and the resulting UV spectral energy distribution is nearly flat in F$_{\\lambda}$ (Fig. 2). Only the most metal-poor starbursts ($<$ 10\\% solar) are relatively unaffected by dust (Fig. 2). These correlations are not surprising, since we expect the dust-to-gas ratio in the starburst's ISM to scale roughly with metallicity.  In principle, relatively red vacuum-UV spectral energy distributions could be due to the effects of the age of the stellar population or the form of the initial mass function (IMF) rather than dust. That is, the redder spectra might correspond to older starbursts, starbursts whose UV spectrum was more heavily contaminated by an older underlying stellar population, or starbursts with a steeper (less O-star-enriched) IMF. However, the strong correlation between the UV color ($\\beta$) and the strengths of the CIV and SiIV (predominantly) stellar wind lines (Fig. 5b) shows that this cannot be the case: an age or IMF effect would require the stellar wind features due to O stars to be {\\it weaker} in the redder starbursts - exactly the opposite of what we observe.  The metal-poor starbursts have systematically weaker vacuum-UV absorption-lines (Fig. 1). In the case of the stellar wind lines (Fig. 3), this result agrees with expectations based on both theory and observations of metal-poor, hot, high-mass stars. In the case of the interstellar lines (Figure 4), this probably reflects both higher average ionic column densities and larger velocity dispersions in the more metal-rich starbursts. That is, even though the strong interstellar lines are optically-thick (and their strength will thus be set to first-order by the velocity dispersion), there will still be a weak dependence of equivalent width on column density. Moreover, the more metal-rich starbursts can also be more powerful (Fig. 8) and are situated in more massive galaxies (Fig. 11). We expect them to have higher average ISM velocity dispersions, since both gravity and the mechanical `stirring' provided by supernovae and stellar winds will contribute to the gas dynamics.  The trend for the more metal-rich starbursts to occur in brighter and more massive host galaxies (Fig. 11 and 12) presumably reflects the well-known mass-metallicity relation defined by normal galaxies (e.g. Zaritsky, Kennicutt, \\& Huchra 1994). That is, the ISM in massive galaxies will already be pre-enriched to relatively high metallicity before the starburst is even initiated. During the starburst itself, the fraction of the newly-synthesized metals that are retained (rather than blown out) should also be a strong function of the depth of the gravitational potential well (e.g. Dekel \\& Silk 1986). The correlation between the metallicity and the luminosity of the starburst (even though it is due in part to selection effects - see section 3.4) may reflect other fundamental trends: more massive galaxies can host more powerful starbursts (Fig. 11) and have more metal-rich interstellar gas with which to fuel the starburst. The luminosity {\\it vs.} rotation-speed connection in starbursts has been previously discussed by Heckman (1993), Lehnert \\& Heckman (1996b), and Meurer et al (1997): basic considerations of causality in a self-gravitating system demand that the maximum star-formation rate will scale as the cube of the rotation speed.  The above results are not only very illuminating regarding the nature of the starburst phenomenon, they also have a variety of interesting implications for the interpretation of the rest-frame-UV properties of galaxies at high-redshift.  First and foremost, starbursts in the present universe emit most of their light in the far-infrared, not in the ultraviolet (Fig. 9). Thus, an ultraviolet census of the local universe would significantly underestimate the true high-mass star-formation-rate and would systematically under-represent the most powerful, most metal- rich starbursts (Fig. 2), occuring in the most massive galaxies (Fig. 11). This {\\it may} also be true at high-redshift, where the current estimates of star-formation rely almost exclusively on data pertaining to the rest-frame vacuum-UV (e.g. Madau et al 1996, 1998). If there {\\it is} a dependence of dust extinction on luminosity at high-redshift, this will affect the apparent shape of the luminosity function. For example, current samples at high-redshift might under-represent young/forming massive elliptical galaxies. These would have high metallicities (and thus a high dust-to-gas ratio) coupled with very large average gas surface mass densities ($\\sim 10^{22} - 10^{23}$ cm$^{-2}$), and would hence have large dust opacities.  Using the strong correlation between the vacuum-UV color of local starbursts ($\\beta$) and the ratio of far-IR to vacuum-UV light emitted by local starbursts, Meurer et al (1998) estimate that an average vacuum-UV-selected galaxy at high-redshift (e.g. Steidel et al 1996; Lowenthal et al 1997) suffers about 2 to 3 magnitudes of UV extinction (in agreement with modeling of the rest-frame UV-through-visible spectral energy distributions by Sawicki \\& Yee 1998). The `correct' prescription for de-extincting the high- redshift galaxies in order to correctly obtain the bolometric luminosity and star-formation rate is a matter of on-going debate (see Pettini 1997 and Madau et al 1996, 1998 for other viewpoints).  For example, one uncertainty at high-redshift is the unknown relative importance of the age or the IMF of the stellar population {\\it vs.} dust in producing the observed UV spectral energy distribution. As we have argued above, the trend for the {\\it redder} starbursts to have {\\it stronger} absorption features due to massive stars (Fig. 5b) rules out age or the IMF as the primary determinant of UV color in our sample. Thus, it will be extremely interesting to see if the same correlation holds in the high-redshift galaxies. In any case, it seems fair to conclude that the history of high-mass star-formation in the universe at early times will remain uncertain until the effects of dust extinction are better understood.  The extinction corrections advocated by Meurer et al (1997) imply large bolometric luminosities for the high-redshift galaxies ($\\sim$ 10$^{11}$ to 10$^{13}$ L$_{\\odot}$ for H$_{0}$ = 75 km s$^{-1}$ Mpc$^{-1}$ and q$_{0}$ = 0.1). Interestingly, the bolometric surface-brightnesses of the extinction-corrected high-redshift galaxies are then very similar to the values seen in local starbursts: $\\sim$ 10$^{10}$ to 10$^{11}$ L$_{\\odot}$ kpc$^{-2}$ (Meurer et al 1997). The high-redshift galaxies thus appear to be `scaled-up' (larger and more luminous) versions of the local starbursts. The physics behind this `characteristic' surface-brightness is unclear (cf. Meurer et al 1997; Lehnert \\& Heckman 1996a). However, it is intruiging that the implied average surface-mass-density of the stars within the half- light radius ($\\sim$ 10$^{2}$ to 10$^{3}$ M$_{\\odot}$ pc$^{-2}$) is quite similar to the values in present-day elliptical galaxies. Are we witnessing the formation of elliptical galaxies and/or bulges, or only the formation of precursor fragments thereof (e.g. Sawicki \\& Yee 1998; Trager et al 1997)?  If the high-redshift galaxies are indeed scaled-up versions of local starbursts, we might expect that they would follow the same trends summarized above. For example, do the high- redshift galaxies should show the same strong correlation between the strength of the UV absorption-lines (stellar and interstellar) and $\\beta$ (Fig. 5), whereby the more metal-rich local starbursts are both redder and stronger-lined? As more near-IR (rest-frame visible) data become available, will the high-redshift galaxies obey the starburst correlations (Fig. 11 and 12) of both vacuum-UV color and absorption-lines strengths with M$_{B}$ and rotation speed (do they follow a mass-metallicity relation?). When the high-redshift galaxies are extinction-corrected, are the more luminous systems redder and stronger-lined in the UV (as in the case of the local starbursts - Fig. 9 and 10 respectively)? Finally, do the strengths of the low-ionization (primarily interstellar) and high-ionization (significant stellar wind contribution) absorption-lines correlate well with one-another, as in the local starbursts (Fig. 7)?  If the answers to these questions are `yes' (that is, if the high-redshift galaxies and local starbursts do occupy the same parts of the multi-dimensional parameter space summarized above), we might be able to use the UV properties of the high- redshift galaxies to `guesstimate' their metallicities. For example, the strong correlation between vacuum-UV color ($\\beta$) and metallicity in local starbursts (Fig. 2) - if applied naively to high- redshift galaxies - would suggest a broad range in metallicity from substantially subsolar to solar or higher and a median value of perhaps 0.3 solar. This is somewhat higher than the mean metallicity in the damped Ly$\\alpha$ systems (Pettini et al 1997; Vladilo 1997; Lu et al 1997), but this may be due to selection effects: the UV-selected galaxies are the most actively star-forming regions of galaxies, while the damped Ly$\\alpha$ systems tend to sample the outer, less-chemically-enriched parts of galaxies (e.g. Prochaska \\& Wolfe 1997) or perhaps proto-galactic fragments or dwarf galaxies (e.g. Vladilo 1997). Our `guesstimated' metallicities of the high-redshift galaxies are much higher than those suggested by Trager et al (1997), who argue that we are witnessing the formation of Population II spheroidal systems.  It is also important to sound a cautionary note: most of the strong absorption-lines in the spectra of local starbursts are of {\\it interstellar} rather than {\\it stellar} origin. This is probably generically true for the low-ionization resonance lines (although B stars may make a non-negligible contribution). Even the high-ionization resonance lines like CIV and SiIV can contain strong interstellar contributions. In fact, in extreme cases (e.g. NGC 1705) even these lines are dominated by interstellar gas. This highlights the difficulty in using the UV spectra of galaxies to deduce their stellar content: data of relatively high spectral resolution and signal-to-noise are needed to reliably isolate the stellar and interstellar components. Simply measuring the equivalent widths of the lines is not enough: it is the {\\it profile shapes} that contain the key information. The stellar photospheric absorption-lines identified in Fig. 1 may be the most straightforward features for the detection of stars in high-redshift galaxies, since they are not resonance lines, and so are uncontaminated by interstellar absorption. Unfortunately, most of these lines are weak. The relatively strong CIII$\\lambda$1175 line is lost in the Ly$\\alpha$ forest at high-redshift, and the CIII$\\lambda$1892, FeIII$\\lambda$1925, and FeIII$\\lambda$1960 lines are often `clobbered' by the OH night-sky lines in galaxies at z $>$ 2.6.  Finally - based on local starbursts - it seems likely that the gas kinematics that are measured in the high-redshift galaxies using the interstellar absorption-lines are telling us a great deal about the hydrodynamical consequences of high-mass star-formation on the interstellar medium, and thus it will be difficult to use them to straightforwardly glean information about the gravitational potential or mass of the galaxy. Even the widths of the optical nebular emission-lines in local starbursts are not always reliable tracers of the galaxy potential well (cf. Lehnert \\& Heckman 1996a; but see Melnick, Terlevich, \\& Moles 1988). This means that it may be tricky to determine masses for the high-redshift galaxies without measuring real rotation-curves via spatially-resolved spectroscopy.  On the brighter side, there is now rather direct observational evidence that the kinematics of the interstellar absorption-lines in the high-redshift galaxies reflect an outflowing metal-enriched gas, thereby allowing us to directly study the `pollution' of the inter-galactic medium with metals in the early universe. The signature of these outflows is interstellar absorption lines that are systematically blueshifted by several hundred km s$^{-1}$ relative to the Ly$\\alpha$ emission-line (Franx et al 1997; Heckman 1997; Pettini 1997). In local starbursts, this is due to outflowing gas that both produces the blue-shifted absorption-lines and absorbs-away the blue side of the Ly$\\alpha$ emission-line (Lequeux et al 1995; Gonzalez-Delgado et al 1998a,b; Kunth et al 1998). The case for an outflow in the high-redshift galaxies could be strengthened by using the weak stellar photospheric lines (Fig. 1) to unambiguously determine the galaxy systemic velocity - cf. Heckman \\& Leitherer (1997) and Sahu \\& Blades (1997).  If these outflows escape the galactic potential well, we should see their cumulative effect in the form of a metal-enriched inter-galactic medium (IGM). The most dramatic such evidence is the existence of an IGM in rich clusters of galaxies whose metal content probably exceeds that of the cluster galaxies themselves. Recent ASCA X-ray spectra of hot gas in clusters of galaxies show approximately solar abundance ratios for the $\\alpha$-process elements like O, Ne, and Si relative to Fe (Mushotzky et al 1996; Ishimaru \\& Arimoto 1997). This implicates `core-collapse' supernovae (the end product of high-mass stars) and - by inference - starburst-driven superwinds as the source of at least 90\\% of the metals in the cluster IGM (Renzini 1997; Gibson, Loewenstein, \\& Mushotzky 1997). More generally, the presence of metals in the `Ly$\\alpha$ forest' at high-redshift (the cloudy component of the early IGM) is certainly suggestive of the dispersal of chemically- enriched material by early superwinds (cf. Cowie et al 1995; Madau \\& Shull 1996). The Ly$\\alpha$ forest clouds appear to have a high ratio of Si/C, suggestive of core-collapse supernovae as the source (Cowie et al 1995; Giroux \\& Shull 1997).  In conclusion, vacuum-ultraviolet observations give us a rich array of diagnostic probes of the stellar and interstellar components of starbursts in the local universe. The analysis described in this paper strongly suggests that these starbursts obey well-defined relationships between such key physical parameters as extinction, metallicity, starburst luminosity and gas dynamics, and host galaxy mass. These objects can therefore serve as local laboratories in which we can study the processes that were important in the galaxy- building events that can now be directly observed at high-redshift. The results of this synergistic study of the most actively-star- forming local and distant galaxies suggests that the latter have rather high UV extinctions (average factor 5 to 10), may have moderately high metallicities (typically 0.1 to 1 times solar), are probably building galaxies having surface mass-densities comparable to present-day ellipticals, and are suffering substantial losses of metal-enriched outflowing gas that can explain the observed `pollution' of the inter-galactic and intra-cluster medium.  {\\bf Acknowledgments}  We would like to thank Anne Kinney and Ralph Bohlin for helping us access and understand the IUE spectra that form the basis of this paper. We also thank Gerhardt Meurer and Daniela Calzetti for many enlightening discussions about the properties of local starbursts, and ditto Mark Dickinson regarding the high-redshift galaxies. This research was primarily supported by NASA LTSA grant NAGW-3138. DG acknowledges the support of the NASA LTSA grant NAG5-6416. This research has made use of the NASA/IPAC extragalactic database (NED), which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the National Aeronautics and Space Administration."
    },
    "9803/astro-ph9803070_arXiv.txt": {
        "abstract": "A phase transition in the nature of matter in the core of a neutron star, such as quark deconfinement or Bose condensation, can cause the spontaneous spin-up of a solitary millisecond pulsar. The spin-up epoch for our model lasts for $2\\times 10^7$ years or 1/50 of the spin-down time (Glendenning, Pei and Weber 1997). ",
        "introduction": "Neutron stars have a high enough interior density as to make phase transitions in the nature of nuclear matter a distinct possibility. According to the QCD property of asymptotic freedom, the most plausible is the quark deconfinement transition. According to lattice QCD simulations, this phase transition is expected to occur in very hot ($T\\sim 200$ MeV) or cold but dense matter.  Since neutron stars are born with  almost the highest density that they will have in their lifetime, being very little deformed by centrifugal forces, they will possess cores of the high density phase essentially from birth if the critical density falls in the range of neutron stars. However the global properties such as mass or size  of a slowly rotating neutron star are little effected by whether or not it has a more compressible phase in the core. In principle, cooling rates should depend on interior composition, but cooling calculations are beset by many uncertainties and competing assumptions about composition can yield similar  cooling rates depending on other assumptions about superconductivity and the cooling processes. Moreover,  for those stars for which a rate has been measured, not a single mass is known. It is unlikely that these measurements will yield conclusive evidence in the present state of uncertainty.  Nevertheless, it may be possible to observe the phase transition in millisecond pulsars by the easiest of measurements---the sign of $\\dot{\\Omega}$. The sign should be negative corresponding to loss of angular momentum by radiation. However, as we will show, a phase transition, either first or second order, that occurs near or at the limiting mass star, can cause spin-up during a substantial era compared to the spin-down time of millisecond pulsars. We sketch the conventional evolutionary history of millisecond pulsars with the addition of the supposition that the critical density for quark deconfinement falls in the density range spanned by neutron stars.  As already remarked, with this supposition, the star has a quark core from birth but its properties are so little effected that this fact cannot be discerned in members of  the canonical pulsar population. However, by some mechanism, usually assumed to be accretion from a companion during the radio silent epoch following its $10^7$ year spin down as a canonical pulsar, the neutron star may be spun up. As it spins up, it becomes increasingly centrifugally deformed and its interior density falls. Consequently,  the radius at which the critical phase transition  density occurs moves toward the center of the star---quarks recombine to form hadrons. When accretion ceases, and if the neutron star has been spun up to a state in which the combination of reduced field strength and increased frequency turn the dipole radiation on again, the pulsar recommences spin-down as a visible millisecond pulsar.  During spin-down the central density increases with decreasing centrifugal force. First at the center of the star, and then in an expanding region, the highly compressible quark matter will replace the less compressible nuclear matter. The quark core, weighed down by the overlaying layers of nuclear matter is compressed to high density, and the increased central concentration of mass acts on the overlaying nuclear matter, compressing it further. The resulting decrease in the moment of inertia  causes the star to spin up to conserve angular momentum not carried off by radiation. The phenomenon  is analogous to that of ``backbending'' predicted for rotating nuclei  (Mottelson and Valatin 1960) and discovered in the 1970's (Johnson et.al. 1972, Stephens and Simon 1972) (see Fig.\\ \\ref{nucleus}) In nuclei, it was established that the change in phase is from a particle spin-aligned state at high  nuclear angular momentum to a superfluid state at low angular momentum.  The phenomenon   is also analogous to an ice skater who commences a spin with arms outstretched. Initially spin decreases because of friction and air resistance, but a period of spin-up is achieved by pulling the arms in. Friction then reestablishes spin-down. In all three examples, spin up is a consequence of a decrease in moment of inertia beyond what would occur in response to decreasing angular velocity. ",
        "conclusions": ""
    },
    "9803/astro-ph9803246_arXiv.txt": {
        "abstract": "The observed Be and B relationships with metallicity clearly support the idea that both elements have a primary origin and that are produced by the same class of objects. Spallation by particles accelerated during gravitational supernova events (SNII, SNIb/c) seems to be a likely origin. We show, in the context of a model of chemical evolution, that it is possible to solve the Li, Be and B abundance puzzle with the yields recently proposed by Ramaty et al. (1997) provided that SNII are unable to significantly accelerate helium nuclei and that different mechanisms are allowed to act simultaneously. ",
        "introduction": "The origin of the light elements is still an open question in Astrophysics. It is widely accepted that standard Big Bang nucleosynthesis can only produce $^7$Li. The primordial abundance produced of this isotope is one order of magnitude below Solar System values and roughly coincides with that observed by Spite \\& Spite (1982) in the hot halo dwarfs ({\\it the lithium plateau}). For this reason, it is generally accepted that such an abundance is representative of the primordial nucleosynthesis and that, since the LiBeB isotopes do not have a cosmological origin, they must be created by galactic activity.  It is well known that the interaction of cosmic rays (CR) with the interstellar medium (ISM) can play an important role in the production of the LiBeB isotopes (Reeves, Fowler \\& Hoyle 1970). In fact, the bulk of the $^6$Li, $^9$Be and $^{10}$B Solar System abundances can be naturally accounted with the standard model of galactic cosmic ray (GCR) nucleosynthesis. In this model, high energy protons and alpha particles collide with heavier nuclei (CNONe) present in the ISM to produce the light element isotopes (Meneguzzi, Audouze \\& Reeves 1971). This model fails, however, to explain the present $^7$Li abundance and the Solar System $^7$Li/$^6$Li and $^{11}$B/$^{10}$B ratios. Furthermore, during the last decade, many observational studies have shown a linear relationship between Be and B abundances in metal-poor stars and the metallicity ([Fe/H]) (Rebolo et al. 1988; Gilmore et al. 1992; Boesgaard 1995; Molaro et al. 1997; Duncan et al. 1997). This {\\it primary} behaviour cannot be easily explained by the standard GCR model: e.g Prantzos et al. (1993) used an ad-hoc hypothesis concerning the time evolution of the CR's escape-length, Abia et al. (1995) introduced artificial time dependences of the CR flux with the metallicity and the star formation rate and Casuso \\& Beckman (1997) considered differential astration [see also Tayler (1995) and Yoshii, Kajino \\& Ryan (1997)]. The reason for using these hypotheses is that in the standard GCR model the LiBeB production rates are proportional to the global metallicity of the ISM and to the CR flux. Since the latter is assumed to be proportional to the supernova rate it is, in consequence, also proportional to the production rate of metals in the galaxy. Such a dependence would predict a slope of about two in the B and Be relationships with [Fe/H] rather than a slope of one as the observations show. Furthermore, Duncan et al. (1997) and Garc\\'\\i a-L\\'opez et al. (1998) have recently shown a nearly constant B/Be ratio of $\\sim 20$ in dwarf stars for a wide metallicity range, although the uncertainties in this ratio are important\\footnote{This B/Be value is obtained taking into account N-LTE effects in the derivation of B and Be abundances}. This value is still compatible with the idea of a spallative origin of Be and B (e.g Fields, Olive \\& Schramm 1994). Since this ratio is similar to that observed in the Solar System (Anders \\& Grevesse 1989) it cannot have experienced large variations during the galactic evolution, which strongly supports the idea of a similar origin for Be and B.  A straightforward interpretation of these results is that the net production rate of Be and B does not depend on the metallic abundance in the ISM, i.e. the light-element production is not dominated by protons and alpha particles colliding with CNO nuclei but by these nuclei colliding with ambient protons and alpha-particles, probably in regions of massive star formation heavily enriched in these nuclei. The $\\gamma$-ray observations from the Orion nebulae (Bloemen et al. 1994) provide additional support to this point of view since they are consistent with line emission from $^{12}$C$^*$ and $^{16}$O$^*$ produced by a large flux of low-energy ($<100$ MeV/nucleon) nuclei enriched in C and O. This might represent the first evidence of the existence of a considerable low-energy component in the spectrum of CRs (at least locally), as was suggested by different authors (Meneguzzi, Audouze  \\& Reeves 1971; Canal, Isern \\& Sanahuja 1980). Since the ejecta of gravitational supernovae (type II, Ib/c) naturally match the above conditions, these objects have been proposed as preferential sites for LiBeB spallation production (Gilmore et al. 1992). The ejecta in these explosions are indeed heavily enriched in CO nuclei and, under some conditions (type Ib/c supernovae), the concentration of these nuclei might even exceed that of H and/or He. Because the yield of CO nuclei in supernovae is almost independent of the metallicity of the progenitor star, the LiBeB produced by spallation during supernova outcomes would have a primary character as is observed for Be and B.  In fact, within the framework of a galatic evolutionary model, Vangioni-Flam et al. (1996) studied the production of Be and B, assuming a low-energy spectrum  of the form $\\rm{q(E)\\sim E^{-n}}$, with $n=9$, and constant for $\\rm{E\\leq 30~MeV/n}$ in the CRs associated with Orion-like regions. They were able to explain the observed behaviour of Be and B vs. [Fe/H], and they obtained upper and lower limits for the contribution of this mechanism to the galactic evolution of Be and B abundances. Their results show that this mechanism might contribute up to $70\\%$ of the observed Be and B abundances and that standard GCR nucleosynthesis is not the main source of $^9$Be and B in the galaxy. Recently, Ramaty et al. (1997; hereafter RKLR) studied the influence of different spectra and chemical compositions in the CRs produced by supernova on the LiBeB production. They showed from energetic arguments that due to the amount of Be necessary to account for its linear behaviour with [Fe/H], the CNO-rich, He-poor and H-poor CR source compositions are favoured. The observed Be/Fe ratio requires the investment of about $3\\times 10^{49}$ to $2\\times 10^{50}$ erg per gravitational supernova in these metallic CR, depending on whether or not H and He are accelerated with metals. Similar arguments led them to conclude that these CRs should have a hard-energy spectrum extending up to at least 50 MeV/n. From the constancy of the observed Be/Fe ratio and metallicity, they also derived the necessary Be yield per supernova that is needed. Assuming a $^{56}$Fe yield per SNII (Woosley \\& Weaver 1995) of $\\sim 0.11$ M$_\\odot$, they obtained a Be yield of $2.8\\times 10^{-8}$ M$_\\odot$ (within a factor of two of uncertainty from the observed Be/Fe ratio), almost irrespective of the progenitor star metallicity.  In this paper we have assumed that the Be yield from RKLR is representative of the Be produced per gravitational supernova. Since this yield automatically sets the corresponding Li and B yields for a given CR spectrum and chemical composition, we use this to study the impact of such objects on the galactic evolution of the light element abundances. We discuss the results in the framework of different scenarios for the progenitors of type II and Ib/c supernovae and possible mechanisms for the CR acceleration in such objects. ",
        "conclusions": "1) Type Ib/c supernovae seem to be key objects for the production of Be and B by spallation. However, due to their low rates in the galaxy, an additional contribution to the Be and B production by spallation due to type II supernovae is needed. Using a reasonable shock spectrum for the CR source, it is possible to explain the observed linear relation of Be and B abundances with metallicity although a neutrino contribution to B is necessary to account for the observed B/Be$\\sim 20$. In this situation, the standard GCR model would play a minor role in the early evolution of the Be and B abundances. Additional sources of $^7$Li are still necessary to account for the evolution of its abundance. Objects with longer lifetime than supernova like AGB stars (and/or novae) seem to be necessary. A multi-source nature of this element is thus obvious.  2) This scenario has a serious problem with the acceleration of CO-rich and HHe-poor matter in SNII ejecta to sufficient energies for spallation reactions. It is not clear how a huge Li production (due to $\\alpha+\\alpha$ reactions) in SNII relative to that of Be and B can be avoided if the acceletared particles are CO-poor. If that is the case a production ratio Li/Be$\\geq 100$ is predicted for a wide variety of CR spectra. This high Li production would be incompatible with the lithium plateau observed at low metallicity as shown in Figure 5 (short dash- dotted line).  3) The CR source spectrum shape associated with supernova explosions is another crucial point. It must extend to high kinetic energies (E$\\geq 50$ MeV/n) in order for this scenario to be energetically plausible and to avoid the overproduction of Li with respect to Be and B. Studies of LiBeB production using other CR spectra, not strict power laws, deserve special attention.  4) The existence of an additional source of $^{11}$B seems obligatory in order to explain the Solar System's $^{11}$B/$^{10}$B$=4.05$ value. Production of $^{11}$B by neutrino spallation seems to be the best candidate. In this case, a $^{11}$B/$^{10}$B ratio higher than 4 at low metallicities is predicted, which might be tested by using very high resolution spectroscopy on very large telescopes (Rebull et al. 1997).  This work has been partially supported by CICYT grants PB96-1428 and ESP95-00091 and by a CIRIT grant."
    },
    "9803/astro-ph9803160_arXiv.txt": {
        "abstract": "s{We present a method of subtracting the foreground contamination for the measurement of CMB polarization. We calculate the resultant errors on CMB polarization and temperature-polarization cross correlation power spectra for the high frequency instrument (HFI) aboard Planck Surveyor, and estimate the corresponding errors on cosmological parameters} ",
        "introduction": "The upcoming satellite CMB experiment Planck surveyor offers an unprecedented opportunity to measure CMB polarization. A major hurdle in extracting the primary CMB signal from data, apart from noise, is galactic and extragalactic foregrounds. However, as the foregrounds differ from CMB in both  frequency dependence and spatial distribution, one can hope to reduce their level in a multi-frequency CMB experiment. A multi-frequency Wiener filtering method to implement this scheme  was developed \\cite{bouchet,teg} and applied to cleaning the  simulated CMB temperature map   by   Bouchet {\\em et al.} \\cite{bouchet}. They showed  that the residual contamination after cleaning the map is much smaller than the CMB primary signal, and therefore the foregrounds may not be a major obstacle in the extraction of  CMB temperature angular power spectrum.  However, as the CMB  polarization signal is expected to be one to two orders of magnitudes below the temperature signal, it is likely to be comparable to both  the experimental  noise and the level of foregrounds. Sethi {\\em et al.}  \\cite{set}  modelled and estimated the level of dust polarized emission at high galactic latitudes. Dust polarization is expected to be the dominant contaminant for  measuring CMB polarization using Planck HFI.   In this paper, we extend the multi-frequency Wiener filtering method to include the polarization and temperature-polarization cross-correlation. The aim of this exercise is to quantify errors in estimating various power spectra and consequently the errors in cosmological parameters. Our results are relevant for Planck HFI. ",
        "conclusions": ""
    },
    "9803/astro-ph9803026_arXiv.txt": {
        "abstract": " ",
        "introduction": "When Baade \\& Zwicky proposed in 1934 that collapsed stars composed of neutrons could be formed in supernova explosions, they had little notion of how such creatures would manifest themselves observationally.  The surprise discovery of pulsars, and in particular, of the Crab and Vela pulsars in their respective supernova remnants, heralded the first visual image of the isolated neutron star: that of a compact, highly magnetized (surface field $\\sim 10^{12}$~G) star, spinning down slowly due to magnetic braking, emitting a collimated beacon, and exciting its surroundings via the injection of ultra-relativistic particles.  However, even the simplest follow-up questions, such as whether all neutron stars form in supernovae, what fraction of supernovae produce neutron stars, and in particular, whether all young pulsars are born with properties like those of the Crab and Vela pulsars remain to this day naggingly unanswered.  Here I summarize observations of neutron stars plausibly associated with supernova remnants.  The last similar review was published by Helfand \\& Becker (1984).  Their interesting synthesis of the observational data is beyond the scope of this paper. More recent reviews including only radio pulsar/supernova remnant associations can be found elsewhere (Kaspi 1996; Kaspi 1998). ",
        "conclusions": "The most striking conclusion following examination of the associations between neutron stars and supernova remnants is that the paradigm that every young neutron star is like the Crab pulsar requires reconsideration.  Neutron stars, it seems, come in many flavors, perhaps some yet to be discovered.  The continued analysis of valuable archival high-energy data, the upcoming launch of the {\\it Advanced X-ray Astrophysics Facility}, as well as the ongoing Parkes multi-beam and Nan\\c{c}ay Galactic plane survey should ensure significant progress, and even perhaps more surprises in this field in the near future."
    },
    "9803/astro-ph9803210_arXiv.txt": {
        "abstract": "We try to explain the unusually high far-infrared emission seen by IRAS in the double-lobed radio-loud quasars 3C\\,47, 3C\\,207 and 3C\\,334. High resolution cm--mm observations were carried out to determine their radio core spectra, which are subsequently extrapolated to the far-infrared in order to determine the strength of the synchrotron far-infrared emission.  The extrapolated flux densities being considerably lower than the observed values, a significant nonthermal far-infrared component is unlikely in the case of 3C\\,47 and 3C\\,334. However, this component could be responsible for the far-infrared brightness of 3C\\,207.  Our analysis demonstrates that nonthermal emission cannot readily account for the difference between quasars and radio galaxies in the amount of their far-infrared luminosity.  On the other hand, a significant role for this mechanism is likely; full sampling of the mm-submm spectral energy distributions is needed to address the issue quantitatively. ",
        "introduction": "\\label{intro}  Thermal emission from cool circumnuclear dust is widely accepted as an important mechanism for producing far-infrared radiation in radio-quiet and radio-loud AGN (Sanders et al.  1986, Chini et al.  1989, Antonucci et al.  1990).  This emission -- provided isotropic -- gives the opportunity to test the orientation unification model of radio-loud quasars (QSRs) and powerful radio galaxies (RGs) (Barthel 1989, Urry \\& Padovani 1995), since matched samples should yield similar far-infrared output for both classes.  However, from IRAS observations it appears that QSRs are stronger 60$\\mu$m emitters than Fanaroff \\& Riley class~II RGs (Heckman et al 1992, 1994, Hes, Barthel \\& Hoekstra 1995).  This would imply that the simple unification model does not hold or that the assumption of isotropic 60$\\mu$m radiation is incorrect.  In support of the latter, Pier \\& Krolik (1992) drew attention to moderate anisotropy effects in the thermal far-infrared radiation from optically thick tori surrounding AGN.  Furthermore, Hoekstra, Barthel \\& Hes (1997) demonstrated that beamed nonthermal far-infrared radiation is not insignificant in radio sources having prominent nuclei, and that the 60$\\mu$m differences might be attributed to a stronger, beamed, nonthermal component in the QSR class.  Within the unified model framework, QSRs would -- at decreasing angle to the line of sight -- naturally lead into the radio-loud blazar class, objects in which the dominance of the beamed component is without doubt.  Indeed, both blazars and core-dominated QSRs display a smooth, single component spectral energy distribution, extending from radio to X-rays, indicating that all continuum radiation is of synchrotron origin (Landau et al. 1986).  This causes blazars and core-dominated QSRs to be the most luminous IRAS AGN (Impey \\& Neugebauer 1988).  IRAS detected three $z\\sim0.5$, lobe-dominated 3CR QSRs at 60$\\mu$m, which corresponds to $\\approx 40\\mu$m emitted wavelength.  Comparable radio galaxies were not detected -- contrary to the expectation within the simple unification model.  The QSRs, 3C\\,47, 3C\\,207, and 3C\\,334, have relatively bright radio cores and large double-lobed structures. Two of these, 3C\\,47 and 3C\\,334, had been observed to display superluminal motion -- a clear sign of beamed emission (e.g., Zensus \\& Pearson 1987) -- thus raising the suspicion that their far-infrared brightness could be (partly) due to a beamed non-thermal component.  The goal of the present research is to investigate to what extent (beamed) nonthermal radiation can account for the infrared emission in lobe-dominated QSRs.  We have determined the cm--mm core spectra of 3C\\,47, 3C\\,207, and 3C\\,334 in order to assess their nonthermal 60$\\mu$m radiation by means of extrapolation.  Given the possibility of flaring submm components (observed in blazars -- e.g., Brown et al. 1989) this research is a first attempt to isolate nonthermal and thermal FIR radiation.  Observations over a wider frequency range are forthcoming. ",
        "conclusions": "\\label{conclusions}  We cannot readily explain the extraordinary high FIR flux density observed in the 3C\\,47, 3C\\,207, and 3C\\,334.  While relativistic beaming of a single nonthermal component cannot account for the total FIR radiation, the radiation of relativistically beamed multiple core components is likely to contribute significantly in the far-infrared. Thermal mechanisms could play a role, in combination with optical thickness and aspect effects, but we need more data in the submm and infrared to address these in detail.  We are currently engaged in measurements of matched pairs of radio galaxies and quasars, comparing their thermal FIR output by combining cm, mm, submm and FIR data from VLA, JCMT, and ISO observations.  In the meantime, we have little doubt that the importance of nonthermal FIR and submm radiation in double-lobed radio sources has been underestimated.  \\bigskip \\noindent  {\\bf Acknowledgments} \\noindent We acknowledge initial involvement in this work of R.~Hes and H.~Hoekstra, OVRO support by N.~Scoville and A.~Sargent, and the use of archival VLA data from R.~Ekers, R.~Perley, and J.~Wardle.  The NRAO VLA is operated by Associated Universities, Inc.  under contract with the National Science Foundation. OVRO research is supported in part by NSF Grant AST~93-14079."
    },
    "9803/astro-ph9803338_arXiv.txt": {
        "abstract": "Preliminary results of a search for distant clusters of galaxies using the recently released I-band data obtained by the ESO Imaging Survey are presented. In this first installment of the survey, data covering about 3 square degrees in I-band are being used.  The matched filter algorithm is applied to two sets of frames that cover the whole patch contiguously and these independent realizations are used to assess the performance of the algorithm and to establish, from the data itself, a robust detection threshold.  A preliminary catalog of distant clusters is presented, containing 39 cluster candidates with estimated redshifts $0.3 \\leq z \\leq 1.3$ over an area of 2.5 square degrees. ",
        "introduction": "\\label{sec:intro}   One of the primary goals for undertaking the ESO Imaging Survey (EIS; Renzini \\& da Costa 1997) has been the preparation of a sample of optically-selected clusters of galaxies over an extended redshift baseline for follow-up observations with the VLT.  High-redshift clusters are, of course, a primary target for 8-m class telescopes. A large and well defined sample of clusters can be used for many different studies, ranging from the evolution of the galaxy population, to the search for arcs and lensed high redshift galaxies, to the evolution of the abundance of galaxy clusters, a powerful discriminant of cosmological models. In addition, individual clusters may be used for weak lensing studies and as natural candidates for follow-up observations at X-ray and mm wavelengths, which would provide complementary information about the mass of the systems. For some of these applications it suffices to find a large number of clusters, while for others it is vital to obtain a full understanding of the selection effects, to generate suitable statistical samples.  Until recently, only a handful of clusters were known at redshifts $z \\gsim 0.5$; visual searches for high redshift clusters were conducted by Gunn \\etal (1986) and Couch \\etal (1991), but their samples are severely incomplete beyond $z \\sim 0.5$; at higher redshifts targeted observations in fields containing known radio-galaxies and QSOs have produced a handful of cluster identifications (\\eg Dickinson 1995; Francis \\etal 1996; Pascarelle \\etal 1996; Deltorn \\etal 1997).  The first objective search for distant clusters was conducted by Postman \\etal (1996; hereafter P96) using the 4-Shooter camera at the 5-m telescope of the Palomar Observatory. In their survey 10 out of the 79 cluster candidates have estimated redshift $\\gsim 0.7$.  Further evidence for the existence of clusters at high redshift has been obtained from X-ray (\\eg Gioia \\& Luppino 1994; Henry \\etal 1997; Rosati \\etal 1998), optical (\\eg Connolly \\etal 1996; Zaritsky et al. 1997) and infrared (Stanford \\etal 1997) searches. However, the existing samples are small, and their selection effects largely unknown.  Recently, observations of the first patch of the EIS, covering about 3 square degrees, have been completed, and the data made available to the community (Nonino \\etal 1998; hereafter Paper I). Although the data are still in a preliminary form, much can already be learned regarding the characteristics of the sample of candidate clusters that can be detected using the EIS data.  In this paper preliminary catalogs of objects detected on single 150 sec. I band frames are used (see Sects. \\ref{sec:obs_and_data} and \\ref{sec:gal_cat}) mainly to assess the capability of the EIS to detect clusters of galaxies at $z \\gsim 0.5$. A discussion of a full cluster sample based on the galaxy catalog extracted from the coadded EIS images, is postponed to a future paper (Scodeggio \\etal 1998). The reason for using here the single-frame catalogs is that they provide two independent datasets for the same area of the sky. The comparison between the cluster detections obtained using the two catalogs separately, can be used to quantify the reliability of the cluster detection procedure.  When the handling of catalogs extracted from the coadded images is fully implemented in the EIS data reduction pipeline, the cluster search will be carried out using those catalogs, instead of the single-frame ones, to benefit from the deeper limiting magnitude of the coadded images.  In the meanwhile a better quantification of the detection limits for distant clusters of galaxies within the EIS data could be obtained by comparing the results presented here with those obtained using independent cluster search methods.  In Sects. ~\\ref{sec:obs_and_data} and ~\\ref{sec:gal_cat} the observations, data reduction and the object catalogs, that are used for the cluster search, are briefly discussed. The cluster finding procedure, based on the matched-filter algorithm proposed by P96, is described in Sect. \\ref{sec:cluster_finding}.  In Sect. ~\\ref{sec:results} the preliminary cluster catalog is presented, and the properties of the detected candidates are discussed.  In Sect. ~\\ref{sec:future} conclusions of this work are summarized, and its possible extensions to the search for clusters using the coadded EIS images discussed. ",
        "conclusions": ""
    },
    "9803/astro-ph9803048_arXiv.txt": {
        "abstract": "Preliminary trigonometric parallaxes and BVI photometry are presented for two dwarf carbon stars, LP765$-$18 (= LHS1075) and LP328$-$57 (= CLS96).  The data are combined with the literature values for a third dwarf carbon star, G77$-$61 (= LHS1555).  All three stars have very similar luminosities (9.6$\\,<\\,$M$_{\\rm V}<\\,$10.0) and very similar broadband colors across the entire visual-to-near IR (BVIJHK) wavelength range. Their visual (BVI) colors differ from all known red dwarfs, subdwarfs, and white dwarfs. In the M$_{\\rm V}$ versus V$-$I color--magnitude diagram they are approximately 2 magnitudes subluminous compared with normal disk dwarfs with solar-like metallicities, occupying a region also populated by O-rich subdwarfs with $-1.5<$[m/H]$<-1.0$.  The kinematics indicate that they are members of the Galactic spheroid population. The subluminosity of all three stars is due to an as-yet-unknown combination of (undoubtedly low) metallicity, possibly enhanced helium abundance, and unusual line-blanketing in the bandpasses considered. The properties of the stars are compared with models for the production of dwarf carbon stars. ",
        "introduction": "The discovery of the first dwarf carbon star (\\cite{dah77}) sparked considerable research on the evolution of stars to and through the dwarf carbon star (dC) stage.  It is now accepted that dC stars are formed through the accretion of carbon-rich material from a close companion star that was evolving on the asymptotic giant branch at the time of mass transfer.   They are now recognized as members of a larger family of stars that have undergone mass-transfer binary evolution, a family that includes the halo CH giant stars, the disk Ba giants, the Ba dwarf or CH subgiant stars, and the extrinsic S stars (\\cite{gre97}).  De Kool \\& Green (1995) have modeled dC star formation by following the evolution of a large variety of binary systems constructed from the observed distributions of component masses, orbital separations, and metallicities for unevolved binaries. Although these simulations involve a large number of poorly constrained parameters, the results indicate that the formation of a dC star is strongly favored by low initial metallicity, and that virtually no dCs are produced in systems with metallicities above half solar.  Furthermore, the mass distribution is found to peak strongly in the 0.4 to 0.9 solar mass range.  Despite the studies to date, our understanding of dC stars remains highly speculative, based as much on plausibility arguments as on objective facts.  Distances of these stars are among the most essential missing data,  and must be known in order to determine each star's luminosity, space velocity, mass, and evolutionary status. Presently, only one dC star, G77$-$61, has a reliable trigonometric distance determination and it indicates a high space velocity and -- by implication -- low metallicity.  This implied low metallicity is consistent with the detailed atmospheric analysis of G77$-$61 by Gass et al. (1988) who derived the extremely low metallicity of [Fe/H]=$-5.6$, a value which certainly implies very early epoch halo formation.  A larger sample of dC stars with properties well constrained by fundamental observational data is clearly needed to understand these stars as a class.  Unfortunately, no known dC stars are bright enough to have had their parallaxes measured with high accuracy by the Hipparcos satellite.  In this paper we report preliminary trigonometric parallaxes for two additional dC stars and discuss some of their properties. ",
        "conclusions": "Results interpreted from the parallax data are shown in Table 2.  In calculating the formal uncertainties in M$_{\\rm V}$ and M$_{\\rm K}$ we have adopted the parallax uncertainties given in Table 1, along with $\\pm\\,$0.02 mag for the uncertainties in the V magnitudes and $\\pm\\,$0.03 mag for the uncertainties in the K magnitudes. The tabulated radial velocities (V$_{\\rm r}$) are from Bothun et al. (1991) for LP765$-$18 and LP328$-$57 and from Dearborn et al. (1986) for G77$-$61. The space velocity components (U,V,W) in Table 1 have been corrected for a solar motion of (10, 15, 7) km s$^{-1}$ to the local standard of rest; U is in the direction of the galactic center.  Inspection of Tables 1 and 2 reveals that these three dC stars are remarkably similar in terms of (1) overall kinematic properties, (2) infrared (JHK) colors, (3) optical (BVI) colors, and (4) luminosities (M$_{\\rm V}$ and M$_{\\rm K}$).  Regarding the kinematics specifically, the large overall space velocities and the large negative V galactic velocity components are indicative of membership in the Galactic spheroid population.  Regarding the colors, Green et al. (1992) have previously discussed the location of dC stars in the J$-$H versus H$-$K diagram and their apparent separation from the giant and subgiant carbon stars, presumably due (in part) to the higher surface gravities for the dC stars. In Figure 1 we show the location of these dC stars in the B$-$V versus V$-$I diagram.  Because all three were originally identified as stars with high proper motion, we include for comparison a selection of other field stars commonly identified in proper motion surveys: later-type degenerates (open circles), disk dwarfs (filled circles), and metal-poor subdwarfs (crosses).  The dC stars occupy a unique region in this diagram.  The lone exception to a clear-cut separation for the dC stars is LP701$-$29, the only known late-type degenerate with strong CaI absorption which blankets the B bandpass (\\cite{dah78}).  Figure 1 suggests that even broadband photometric surveys of faint, high proper motion stars might succeed in isolating additional dC candidates.  Their location to the left of the most extreme metal-poor field subdwarfs known at present suggests metallicities below [m/H]$\\sim-$2.  However, quantitative conclusions about metallicity are not possible because the field subdwarfs have O-rich rather than C-rich atmospheres; with the exception of the analysis of G77$-$61 by Gass et al. (1988), no complete studies based on model atmospheres have been presented for metal-poor, C-rich stars with dwarf-like gravities.  Figure 2 shows the location of these dC stars in the M$_{\\rm V}$ versus V$-$I color magnitude diagram.  The solid line represents the observed mean disk main sequence as defined by USNO parallax stars (cf. \\cite{mon92}).  The dashed lines show the metal-poor main sequences modeled by Baraffe et al. (1997) for metallicities (scaled from solar) of [m/H]= $-1.0, -1.5$ and $-2.0.$ These authors demonstrated that their models successfully reproduce the main sequence of several globular clusters over this range of metallicities.  However, these models are for O-rich atmospheres -- not the C/O$>$1 compositions which are appropriate for the dC stars. Furthermore, the position of dC stars in color-magnitude diagrams cannot necessarily be interpreted simply in terms of metallicity because of the possibility of enhanced helium abundance in their atmospheres, produced by the mass transfer event(s) that increased their carbon abundances.  The complicated atmospheric situation and the difficulty in uniquely determining both log$\\,g$ and the He abundance has been discussed by Gass (1988) and Gass et al. (1988).  Therefore, until more detailed spectroscopic investigations are undertaken to derive the helium abundance independently, we are left with the implication from the kinematic information that low metallicity and nearly normal helium abundances are most likely.  Figure 3 shows the location of the dC stars in the schematic M$_{\\rm K}$ versus I$-$K diagram. Here the solid line represents the ``young disk'' stars defined by Leggett (1992) while the dashed lines are again the metal-poor main sequences models from Baraffe et al. (1997).  In their analysis of G77$-$61, Gass et al. (1988) pointed out that the J and K fluxes should not depend strongly on the carbon abundance whereas the fluxes in the B,V,I and H bandpasses will depend strongly on composition. The three dC stars are also significantly subluminous with respect to the disk main sequence in this diagram.  However, it is also clear that going to a redder color index still does not allow a quantitative interpretation in terms of metallicity because the location of G77$-$61 above the [m/H]=$-1.0$ curve is clearly at odds with the value of [Fe/H]=$-5.6$ derived by Gass et al. (1988).  Figure 4 shows the location of G77$-$61 in the M$_{\\rm V}$ versus T$_{\\rm eff}$ plane using the value of T$_{\\rm eff}\\,=\\,$4200$^{\\circ}$K derived by Gass et al. (1988).  Once again the dashed lines are from Baraffe et al. (1997).  Here we find at least qualitative agreement in that this extremely metal-poor object lies below the [m/H]=$-2.0$ curve.  This emphasizes the need for higher resolution spectrophotometric observations and C-rich model atmosphere analyses for LP765$-$18 and LP328$-$57 in order to establish (at least) T$_{\\rm eff}$ values for them.  The models of Baraffe et al. (1997) indicate that the derived masses are primarily sensitive to the absolute magnitudes (or luminosities), and only weakly dependent on the colors.  Acknowledging the dangers of interpreting the dC stars with O-rich models, the masses so estimated are 0.39 M\\solar\\ for LP765$-$18, 0.37 M\\solar\\ for LP328$-$57, and 0.30 M\\solar\\ for G77$-$61.  Models for dC star formation predict that a range of masses from 0.2 to 1.0~M{\\solar} should exist (\\cite{dek95}).  However, the frequency distribution of these masses depends critically on the adopted Initial Mass Ratio Distribution (IMRD) for the unevolved binary which is quite uncertain.  For a flat IMRD (i.e., dN$\\,\\propto\\,$dq, where q is the mass ratio), the distribution for spheroid dC stars peaks rather sharply around 0.7 M\\solar\\ at a space density of dN/dlog$\\,$M = 4$\\,{\\rm x}\\,10^{-7}$ pc$^{-3}$, falling to roughly dN/dlog$\\,$M = 2$\\,{\\rm x}\\,10^{-7}$ pc$^{-3}$ for stars in the mass range inferred above.  On the other hand, for uncorrelated component masses, the models predict a broader peak spanning 0.2 to 0.8 M\\solar\\ with a space density of dN/dlog$\\,$M = 4$\\,{\\rm x}\\,10^{-7}$ pc$^{-3}$ for spheroid stars.  As noted by de Kool \\& Green (1995), the observed absolute magnitudes seem to support an IMRD with uncorrelated component masses.  However, the apparent similarity of these three dC stars is undoubtedly influenced to some extent by selection effects. Hotter, more massive dC stars are predicted by the models to be common, but hotter stars will have weaker carbon bands and may have spectra and broad-band colors that are not as obviously different from stars with oxygen-rich atmospheres as are these three dC stars.  Another factor that makes the model predictions uncertain is the amount of dilution of the carbon-rich material being transferred from the asymptotic-giant-branch star onto its main-sequence companion:  a more massive main sequence star (0.5--1.0~M{\\solar}, for example), has a less massive convective envelope than the stars studied here (reducing the dilution), but the carbon-rich material may be mixed into the radiative zone as well (\\cite{pro89}), thus increasing the dilution and making hotter dC stars more difficult to identify.  Furthermore, these three dC stars have been chosen for parallax observations based partly on their unusually-high proper motions.  This kinematic selection partially accounts for their all being members of the Galactic spheroid. Disk dC stars may be common --- the ratio of disk/spheroid dC stars is a quantity that will help constrain the models of dC-star formation.  Parallax data for a sample of dC stars with smaller proper motions could help to address this issue.  Several such dC stars are probably close enough to get significant parallax measurements -- however, none are presently being observed in the Naval Observatory parallax program pending time becoming available in the program at their respective right ascensions.  We conclude that the similar properties of these three dC stars and their membership in the Galactic spheroid may result in part from the way these stars were selected.  Accurate distances for a larger sample of dC stars are needed, and identification of dC stars without kinematic bias is essential in order to fully understand how and where dC stars are produced."
    },
    "9803/astro-ph9803081_arXiv.txt": {
        "abstract": "Oscillations observed in the light curve of Nova V1974 Cygni 1992 since summer 1994 have been interpreted as permanent superhumps. From simple calculations based on the Tidal-Disk Instability model of Osaki, and assuming that the accretion disc is the dominant optical source in the binary system, we predict that the nova will evolve to become an SU UMa system as its brightness declines from its present luminosity by another 2-3 magnitudes. Linear extrapolation of its current rate of fading (in magnitude units) puts the time of this phase transition within the next 2-4 years. Alternatively, the brightness decline will stop before the nova reaches that level, and the system will continue to show permanent superhumps in its light curve. It will then be similar to two other old novae, V603 Aql and CP Pup, that still display the permanent superhumps phenomenon 79 and 55 years, respectively, after their eruptions. We suggest that non-magnetic novae with short orbital periods could be progenitors of permanent superhump systems. ",
        "introduction": "\\subsection{The (permanent) superhump phenomenon}  Regular superhumps are quasi-periodic oscillations that appear in the light curve (LC) of the SU UMa subclass of dwarf novae, superimposed on the superoutburst LC of these systems. The superhump period is a few percent longer than the orbital period of the system in which it is observed (laDous 1993). Permanent superhumps (PSH), on the other hand, appear permanently or most of the time in systems where they prevail, and do not demand a superoutburst for their emergence. Superhumps are detected in the LC of cataclysmic variables (CVs) with short orbital periods ($P_{orbital}$=17-200 min. - Ritter \\& Kolb 1998). Stolz \\& Schoembs (1984) found a linear relation between the relative excess of the superhump period over the orbital one, and the orbital period.  The Tidal-Disc Instability model (for a review see Osaki 1996) is the commonly accepted explanation of the phenomenon. The superhump periodicity is the beat of the orbital period of the binary system with the period of the precession of the accretion disc around the white dwarf (WD) in the binary co-rotating frame. The difference between regular superhumps and PSH is understood as a result of the much larger mass transfer rate in the system showing PSH than in SU UMa systems. In fact, Osaki's schematic diagram (Fig. 1) states that the values of only two parameters, the orbital period and the mass transfer rate, determine the basic differences among the four major classes of CVs, namely, U Gem, SU UMa, PSH and Nova-like variables.  \\begin{figure}  \\centerline{\\epsfxsize=3.0in\\epsfbox{fig1nneps.eps}}  \\caption{A schematic theoretical diagram based on Osaki (1996), showing the location on the (Orbital Period, $\\dot{M}$) plane of four major groups of CVs: UG=U Geminorum, SU UMa, PSH=permanent superhumps and NL=nova like variables. The two dashed vertical lines represent the two ends of the well known period gap in the period distribution of CVs. Systems on the right hand side, the UG and NL groups, have accretion discs that are tidally stable. Systems on the left, PSH and SU UMa, are tidally unstable. Systems above the dotted tilted line, PSH and NL, have thermally stable discs. Systems below that line, UG and SU UMa, have thermally unstable discs. The upper dot represents the present location of V1974 Cyg in this plane. The arrow indicates the expected decrease in the $\\dot{M}$ value in the future. The lower point is the location of value of $\\dot{M}$, corresponding to the pre-nova magnitude according to Skiff (1997). If the fading of the nova, in magnitude units, continues linearly, the arrow will intersect the tilted line around the year 2000.}  \\end{figure}  \\subsection{Superhumps in Classical Nova Systems}  Photometric observations of V1974 Cyg (N. Cyg 1992) revealed two periodic oscillations in the LC of the star (Retter, Ofek \\& Leibowitz 1995). While there is an overall agreement that the shorter is the orbital period of the binary system, two interpretations for the nature of the second period, which is about 5\\% larger than the orbital one, have been suggested.  Semeniuk et al. (1995) and Olech et al. (1996), argued that it is the spin period of a rotating WD. Retter, Leibowitz \\& Ofek (1997) and Skillman et al. (1997) proposed, on the other hand, that the longer period is that of PSH oscillations. We think that the increase of this period during 1995 and 1996, as well as other features of the optical LC of the nova discussed in Retter et al. (1997), argue against the magnetic interpretation for the longer period, and for the PSH one, which we adopt in this letter.  Two other old novae, V603 Aql 1918 (Patterson \\& Richman 1991; Patterson et al. 1993; Thomas 1993 and Patterson et al. 1997) and CP Pup 1942 (White \\& Honeycutt 1992; White, Honeycutt \\& Horne 1993 and Thomas 1993), also show PSH in their LCs.  While the presence of PSH in the LC of V603 Aql is rather well established, there is still some controversy in the case of CP Pup.  White et al. (1993) and Balman, Orio \\& Ogelman (1995) proposed a WD spin interpretation for the second periodicity in the LC of this nova, which was initially thought to be 11\\% longer than the orbital period. However, an extensive photometric study by Thomas (1993) revealed that the period excess is only about 2\\%, and that the two periods obey the well-known relation of Stolz \\& Schoembs (1984) between orbital and superhumps periods in SU UMa systems. In addition, the spin periods in all but one intermediate polars are shorter than their orbital periods (Patterson 1994). Even in the one exceptional case (RX J19402-1025), the period excess is very minute - $(P_{spin} - P_{orbital})/ P_{orbital} \\approx$ 0.3\\% (Patterson et al. 1995), much smaller than in a typical superhump system.  We believe that the observed photometric features of CP Pup favour the superhump interpretation for this system and we adopt it in this work.  There are a few more reports on short periods in other novae.  RW UMi 1956, with a possible orbital period of $\\sim$2~hr (Szkody et al. 1989), is a natural permanent superhump candidate, but intensive continuous photometry of the nova during 25 nights in 1995 and 1997, carried out at the Wise Obs., suggests that the variation is irregular.  Two periods of $\\sim$3.3~hr, which differ from each other by less than 2\\%, were found in V1500 Cyg 1975. These periods, however, don't obey the Stolz \\& Schoembs relation. The variation in the polarization, found in this system, is also a very strong evidence for the magnetic nature of the WD. (Stockman, Schmidt \\& Lamb 1988).  There are also some evidence that the 85-m period found in GQ Mus 1983 is caused by the rotation of the magnetic pole of the primary (Diaz \\& Steiner 1994), thus cannot be identified with a superhump variation. ",
        "conclusions": "\\subsection{Consistency}  The value of the parameters of the three nova systems that we used in the previous section were derived by various researchers, in general independently of the superhumps phenomenon. Our calculations show that according to Osaki's theory, the very presence of superhumps in the LCs of these systems is indeed expected from these values.  Further support for the validity of the calculations presented in this work, particularly for the bolometric correction that we used in Section 2.1 (equation 5), comes from the agreement of the mass transfer rate that we obtain for V603 Aql with estimates by other authors using other methods. We derive in Section 2.3 for this parameter the value $8\\pm4 \\times 10^{17}$ gr/sec. It is in good agreement with  $4.8 \\times 10^{17}$ gr/sec (Krautter et al. 1981), $7.6 \\times 10^{17}$ gr/sec  (Wade 1988), and $2.6 \\times 10^{17}$ gr/sec (Duerbeck 1992).  \\subsection{The principal light source}  Our calculations in Section 2 were made with the assumption that the visual light of all three novae emanates from the accretion discs in these systems. This assumption is supported by the optical spectrum of the three stars. Shai Kaspi kindly took for us two spectrograms of V1974 Cyg with the FOSC camera at the Wise Obs., one in 1995 July and one in 1996 August. The two spectra are dominated by nebular emission lines and by a continuum that seems to be free of stellar absorption features.  Similarly, spectra of V603 Aql (Williams 1983) and of CP Pup (O'Donoghue et al. 1989, White et al. 1993), which were taken 69 years after outburst in the V603 Aql case, and 43, 45 and 46 years after outburst in the CP Pup case, do not show any obvious stellar absorption features. Thus, any contribution of the secondary to the light of these systems is indeed negligible.  Leibowitz (1993) has drawn attention to the appearance of a kink in the visual LC of many classical novae a few tens of days after maximum light. He suggested that the abrupt change in the slope of the decaying LC signifies a transition of the main light source in the system from the envelope of the contracting nova to the accretion disc in the system.  In the two old novae, V603 Aql and CP Pup, the photometry capable of detecting superhumps was performed many years after the outburst, long after the appearance of the kink in their LCs. In V1974 Cyg, the superhumps were also detected only after the kink in the LC of this nova, as shown in Fig.~2. This figure is a comprehensive visual light curve of the nova, from outburst in 1992 February to 1997 May. The data were taken from various amateur groups. The arrow in the figure indicates the time when superhumps were first observed. Since superhumps are a pure disc phenomenon and consistently with Leibowitz's hypothesis we may conclude that presently in V1974 Cyg, the accretion disc is indeed the major source of the optical continuum.  \\subsection{The future of the permanent superhumps in the two old novae}  We showed in Section 2 that in the three novae discussed in this letter, the present visual magnitude is brighter than the critical value for transition from the PSH to the SU UMa state. The future of the superhumps in their LCs is correlated with the future run of their visual LCs.  The question of what is the long term behaviour of the LC of classical novae, many years after outburst, is a wide open one. It is difficult to answer it observationally because of the scarcity of these objects on one hand, and the early age, of less than a century of modern observations in novae, on the other. This time duration is probably still shorter than the characteristic time of the returning of classical novae to their real quiescence state.  In two pioneering works, Vogt (1990) and Duerbeck (1992) attempted to systematically investigate the long term photometric behaviour of classical novae. Duerbeck's sample of 21 old novae includes also the two old novae discussed in this letter. There are large systematic difficulties and uncertainties in the determination of the decline rate of old novae, as underlined by Duerbeck himself. He, nevertheless, suggested for V603 Aql an average decline rate of 10.7 mmag/year, and for CP Pup 36.9 mmag/year (Duerbeck 1992 Table 1).  Taking these numbers at face value and using our results of Section 2.3, we obtain that V603 Aql will cross the critical line in Fig. 1 not before some 75 years in the future from today. CP Pup will undergo this phase transition not before some 55 years from today. In view of the large uncertainties in the value of the critical visual magnitude of these two novae, as well as in their average decline rate, these time intervals are rough lower limits at best.  \\subsection{The future of the permanent superhumps in Nova Cygni 1992}  V1974 Cyg is still in a phase of relatively steep decline from its recent outburst. In its visual LC shown in Fig. 2, three sections of nearly linear decline, at three different slopes, are clearly discernible. The decline rate in the last section, lasting now for more than two  years, is about 0.7 mag/year. If the nova keeps this decline rate for another few years, it will reach the critical visual magnitude, indicated by the dot in the figure, sometime around the year 2000.  The vertical error bars around the dot denote the formal uncertainty in the critical visual magnitude value (see Section 2.2). The horizontal bars represent the uncertainty in the time of crossing the line of phase transition, due to the uncertainty in the slope of the present average linear decline.  \\begin{figure}  \\centerline{\\epsfxsize=3.0in\\epsfbox{fig2neps.eps}}  \\caption{Five years light curve of Nova~Cyg~92. The data were taken from visual estimates of three groups of amateur astronomers - AFOEV (main source), VSNET and AAVSO. The solid line is a linear extrapolation based on the mean decline rate of the nova in the last two years. See text for further details.}  \\end{figure}  \\subsection{Sensitivity of the results to the adopted parameters of N. Cygni 1992}  Table 1 presents values of system parameters that we used in our calculations.  When more than one value is suggested in the literature, the table gives a range of possible values, where the range limits are two particular values that have been suggested for the corresponding parameter. The results of our calculations in Section 2.2 and 2.3 are accordingly given also as interval of possible critical values. For a few parameters of V1974 Cyg, however, there are in the literature more extreme estimates that put the parameter value outside the indicated interval. Here we show that even when using these more extreme values in our calculations, the results do not change very significantly.  Balman et al. (1997) suggested an upper limit of 1.37 $M_{\\odot}$ on the mass of the WD in V1974 Cyg. Paresce et al. (1995) proposed the low value of 0.75 $M_{\\odot}$ for this parameter. The use of these two extreme values in our calculation widens the range of the predicted critical visual magnitude of V1974 Cyg by no more than 0.4 magnitude on each side. Similarly, the upper limit on the distance to this system of 3.2 kpc (Paresce et al. 1995) increases the critical range of magnitudes by about 1.1 mag. It seems, however, that the expansion velocities of the nebulae used by these authors are rather large. All these possible modifications, while changing the calculated time that is required before phase transition occurs, do not alter the qualitative scenario, presented in this work. Nova Cyg '92 is indeed expected to decline further in its optical brightness, as it is presently still nearly four magnitudes brighter than the system pre-nova magnitude $m_{V}=19.5$ (upper horizontal line in Fig. 2 - Annuk et al. 1993). It may have even larger room to decline, if the progenitor had the magnitude $m_{B}=21$ (Fig. 2 - lower horizontal line - Skiff 1997).  Naturally, the possibility that V1974 Cyg changes its rate of decline, or even reverses it into brightening, before reaching the critical visual magnitude, cannot be ruled out. In this case PSH will continue to prevail in the LC of this system, much like they do in the LCs of V603 Aql and CP Pup.  \\subsection{A proposed evolution scenario}  Based on the example of the three classical novae V603 Aql, CP Pup and V1974 Cyg, and on the calculations presented in this letter, we may speculate that short period novae with non-magnetic WDs may be the progenitors of PSH systems. This quasi-stable stage in nova life might take a few decades or centuries, before the system returns to its quiescent state, and then becomes a regular SU UMa star. Naturally, such a scenario must be checked quantitatively by a proper statistical analysis of the population of the involved classes of stars."
    },
    "9803/hep-ph9803394_arXiv.txt": {
        "abstract": "We study the conditions under which an overdamped regime can be attained in the dynamic evolution of a quantum field configuration. Using a real-time formulation of finite temperature field theory, we compute the effective evolution equation of a scalar field configuration, quadratically interacting with a given set of other scalar fields. We then show that, in the overdamped regime, the dissipative kernel in the field equation of motion is closely related to the shear viscosity coefficient, as computed in scalar field theory at finite temperature. The effective dynamics is equivalent to a time-dependent Ginzburg-Landau description of the approach to equilibrium in phenomenological theories of phase transitions. Applications of our results, including a recently proposed inflationary scenario called ``warm inflation'', are discussed.  \\vspace{0.34cm} \\noindent PACS number(s): 98.80 Cq, 05.70.Ln, 11.10.Wx ",
        "introduction": "Kinetic equations describe the time evolution of a certain chosen set of physical variables. The choice of physical variables in principle is arbitrary, but often in practice is governed by the measurement of interest. Typical examples are the order parameter of a complex system or the coordinate of a Brownian particle in a heat reservoir. The kinetic approach is usually implemented through a proper separation of the microscopic equations of motion of the chosen physical variables into regular and random parts. An averaging over the random part then generates the effective partition function for the regular part. This averaging is often referred to as a coarse-graining.  One typical application of the kinetic approach is when the physical variables of interest possess energy in relative excess or deficiency to the rest of a large system. Kinetic theory then describes the approach to equilibrium of the chosen physical variables, as for example in the kinetics of phase transitions or in Brownian motion. In the former case, the system is able to release energy to the environment due to some change in its internal state. Provided the environment is disproportionately large relative to the system, the process is irreversible. For a continuous transition, the focus of the present work, this process of equilibration can be described by the monotonic change of an appropriate order parameter, which is the chosen physical variable. Many systems are known to relax in this manner. Phenomenologically, they are successfully described by the time-dependent Ginzburg-Landau theory \\cite{gold}.  Here we are interested in examining under which circumstances physical variables whose microscopic dynamics is second order in time, as for example the Higgs order parameter of spontaneous symmetry breaking, may have a dynamics which is effectively first order in time as in Ginzburg-Landau phenomenological models.  Qualitatively it is not difficult to argue the plausibility of this standard picture for the Higgs symmetry breaking scenario. A single variable, the Higgs order parameter, is modeled to control the release of energy to all the modes that couple to it. By basic notions of equipartition, one anticipates that some portion of the order parameter's energy will flow irreversibly to any given mode. Provided the Higgs order parameter couples to a sufficient number of modes, the motion of the order parameter will be overdamped.  In particle physics models, Higgs symmetry breaking is accompanied by mass generation. Thus the natural couplings for the Higgs field $\\phi$ to bosonic fields $\\chi_i$ is $\\phi^2 \\chi_i^2$, gauge fields $A^{\\mu}_i$ is $\\phi^2 A^{i\\mu} A_{i \\mu}$ and fermionic fields $\\psi_i$ is $\\phi {\\bar \\psi}_i \\psi_i$. For a microscopic realization of time dependent Ginzburg-Landau theory for the Higgs scalar order parameter in a particle physics setting, these are the most obvious types of couplings to investigate. In this paper we will examine the case of purely bosonic couplings in the ``symmetry restored'' regime. That is, we will study the relaxation of an order parameter which is initially away from the only minimum of the free energy density describing the system. Much of the formalism required for this has already been done in \\cite{GR} but we will extend that calculation to the overdamped regime. In an upcoming paper, we plan to study the symmetry broken case.  To our knowledge, this paper is the first study of overdamping in quantum field theory with realistic couplings between system and environment, as inspired by particle physics. Overdamping has been studied in quantum mechanical reaction rate theory for a particle escaping from a metastable state (for a review please see \\cite{hangii}). This is sometimes referred to as the Kramer's problem, with the overdamped limit also called the Smoluchowski limit. Quantum mechanical models describing this problem are commonly of the system-heat bath type. Microscopic quantum mechanical models have been constructed along these lines, in which the particle (system) is coupled to a set of otherwise free harmonic oscillators (heat bath). Such microscopic system-heat bath models are often referred to as Caldeira-Leggett (CL) models. In many cases they have been exactly solved \\cite{fordkac}. The overdamped limit has been derived in these models for the case where the coupling is linear with respect to the oscillator variables but arbitrary with respect to the particle variable \\cite{hangii,cl2}.  A Caldeira-Leggett type model has also been formulated for the case where the system is a self interacting scalar quantum field coupled linearly to a set of otherwise free fields and the overdamped limit has been obtained \\cite{ab54}. This model does provide a microscopic quantum mechanical realization of time-dependent Ginzburg-Landau dynamics in scalar quantum field theory. However, since the couplings between system and environment variables are linear, it should be considered as a first step toward more realistic treatments. More importantly, the calculational method used in \\cite{ab54} cannot be extended to the case when the system variable couples quadratically to other fields.  Although the analysis of overdamping in this paper has general applicability, it was motivated by the warm inflation scenario of the early universe \\cite{ab54,wi}. In \\cite{wi} it was realized that the standard Higgs symmetry breaking scenario, when put into a cosmological setting, provides suitable conditions for the universe to enter a de Sitter expansion phase and then smoothly exit into a radiation dominated phase. The overdamped motion of the order parameter in this scenario may sustain the vacuum energy sufficiently long for de Sitter expansion to solve the horizon and flatness problems. Simultaneously, the relaxational kinetics of the order parameter can maintain the temperature of the universe and permit a smooth exit from the de Sitter phase into the radiation dominated phase. Finally, the thermal fluctuations of the order parameter provide the initial seeds of density perturbations, which in addition could be scale free under specified conditions \\cite{wi,bf2}. An elementary analysis of this scenario, based on Friedmann cosmology for general realizations of order parameter kinematics, indicated that if the universe's temperature does not fall too much during de Sitter expansion, then the cosmological expansion factor from the de Sitter phase should be of order the lower bound set by observation \\cite{ab55}. Although this is not a tight constraint of this scenario, it is a natural one. An analysis of COBE data motivated by this expectation did indicate a slight preference for a small super-Hubble suppression scale, which could be interpreted as arising from a de Sitter expansion with duration near its lower bound \\cite{bfh}. Furthermore, the overdamped limit required by warm inflation, when expressed in different terms, was noted \\cite{ab54} to be an adiabatic limit, for which known methods from dissipative quantum field theory \\cite{GR,hs1,morikawa} are presumed valid. These facts provide further motivation to seek a microscopic model of the scenario, which is the goal of the present work.  The calculational methods used here, based on Schwinger's close-time path formalism, were developed in \\cite{GR}. There are several other works in the literature that apply this formalism to a variety of different situations. [See, for example, the works of Refs. \\cite{hu,boya1,morikawa,greiner,yoko}.] The new feature of the present paper is to shift focus to a kinematic regime dominated by strong dissipation, in order to establish under which conditions this regime leads to overdamped motion. This approach will allow us to have a unique understanding of the microphysical origin of such dynamical behavior, which is in general invoked phenomenologically in applications ranging from condensed matter physics to inflationary cosmology.   The paper is organized as follows. In Sec. II our model of interacting bosons is presented and the effective action is computed perturbatively for a homogeneous time dependent background field configuration $\\phi(t)$. In Sec. III the effective Langevin-like equation of motion is obtained for $\\phi$ in the symmetry-restored phase. In Sec. IV the overdamped limit of this equation of motion is derived and regions of validity are given. In Sec. V the results of the previous sections, which are for Minkowski space, are extrapolated into a cosmological setting and a preliminary examination is made of the warm inflation scenario. In Sec. VI concluding remarks are given. Two Appendixes are included to clarify a few technical details, like the evaluation of the imaginary part of the self-energies and to stress the importance of taking fully-dressed field propagators to properly describe dissipation in the adiabatic approximation for the field configuration. ",
        "conclusions": "\\label{sec6}  In this paper a microscopic quantum field theory model has been presented, describing overdamped motion of a scalar field. Commonly, such behavior is treated phenomenologically by Ginzburg-Landau order parameter kinetics. Our model provides a first principle explanation of how kinetics equivalent to the Ginzburg-Landau type, which is first order in time, arise for inherently second order dynamical systems. The microscopic treatment of this problem, in principle, should be well controlled, due to its fundamental reliance on the adiabatic limit, and our model exemplifies this expectation.  The calculational method for treating dissipation in this paper has one distinct difference from several other related works. In our calculation, we consider the effect of particle lifetimes in the effective equation of motion. To our knowledge, this effect has been discussed in only a few works in the past \\cite{GR,hs1,morikawa}.  Secs. II-IV presented a general, flat-space treatment, which offers a microscopic justification to the often used limit of diffusive Ginzburg-Landau scalar field dynamics. We have shown how it is possible to obtain an effective evolution for the scalar field which is first-order in time, due to its own thermal dissipation effects, interpreted microscopically as its decay into many quanta. In a sense, the field acts as its own brakes, the slowing of its dynamics being attributed to the highly viscous medium where it propagates, a densely populated sea of its own decay products.  The application that we considered in Sec. V was in expanding spacetime, for the cosmological warm inflation scenario. Although we did not formally derive the extension of our flat-space model of Secs. II-IV to an expanding background, we did present heuristic arguments that validate this extension for the special needs of warm inflation. The results of the simple analysis in Sec. V are strongly dependent on initial conditions and may be difficult to implement for models of observational interest. Nevertheless, these results will provide useful guidance both for modifications of this model and for our next study of the symmetry broken regime.  The direct significance of the present study to inflationary cosmology would be to the initial state problem \\cite{muw} in scenarios during symmetry breaking. The initial conditions required for warm inflation in the symmetry broken case are similar to new inflation. The requirement is a thermalized inflaton field, which at the onset of the warm inflation regime is homogeneous with expectation value $\\langle \\phi \\rangle_{\\beta}=0$. Although we have made no detailed application of our results to this problem, some general features are evident from the analysis in Sec. V. In particular, both the suppression of large fluctuations and thermalization are mutually consistent with strong dissipative dynamics. Many of the difficulties that have been discussed \\cite{muw,cl} in association with the initial state problem, are eliminated in the strong dissipative regime. In addition, the damping of fluctuations should simplify the formal problem of coupling this model to classical gravity. Thus, the strong dissipative regime appears to have the correct features both to carry the universe into an inflation-like phase and then to smoothly exit into a hot big-bang regime."
    },
    "9803/astro-ph9803204_arXiv.txt": {
        "abstract": "We present a new well defined sample of BL\\,Lac objects selected from the ROSAT All-Sky Survey (RASS). The sample consists of 39 objects with 35 forming a flux limited sample down to $\\rm f_{X}(0.5 - 2.0\\,keV) = 8\\cdot 10^{-13}\\,ergs\\,cm^{-2}\\,s^{-1}$, redshifts are known for 33 objects (and 31 of the complete sample). X-ray spectral properties were determined for each object individually with the RASS data. The luminosity function of RASS selected BL\\,Lac objects is compatible with results provided by objects selected with the {\\em Einstein} observatory, but the RASS selected sample contains objects with luminosities at least tenfold higher. Our analysis confirms the negative evolution for X-ray selected BL\\,Lac objects found in a sample by the {\\em Einstein} observatory, the parameterization provides similar results. A subdivision of the sample into halves according to the X-ray to optical flux ratio yielded unexpected results. The extremely X-ray dominated objects have higher redshifts and X-ray luminosities and only this subgroup shows clear signs of strong negative evolution. The evolutionary behaviour of objects with an intermediate spectral energy distribution between X-ray and radio dominated is compatible with no evolution at all. Consequences for unified schemes of X-ray and radio selected BL\\,Lac objects are discussed. We suggest that the intermediate BL\\,Lac objects are the basic BL\\,Lac population. The distinction between the two subgroups can be explained if extreme X-ray dominated BL\\,Lac objects are observed in a state of enhanced X-ray activity. ",
        "introduction": "The most common view about BL\\,Lac Objects is that we are looking into a highly relativistic jet (Blandford \\& Rees, \\cite{blandford}). The high variability, the polarization, and the spectral energy distribution can in principle all be explained with this model. But there are still unsolved problems for this model. Examples are the nature of the mechanism(s) that generates and collimates the jet. The evolution of physical parameters as the energy and momentum density along the jet is not yet clear, too. An important question is also whether the jets are composed of highly relativistic hadronic or leptonic plasma (Kollgaard, \\cite{kollgaard}). Furthermore there exists the competing theory that at least some BL\\,Lac objects originate from microlensed QSO (Ostriker \\& Vietri, \\cite{ostriker}). Models of BL\\,Lac objects must explain why X-ray and radio selected BL\\,Lac objects differ in some observational characteristics (degree of variability and polarization (Jannuzi et al., \\cite{jannuzi})) so that it is not necessarily true that both object classes are of the same astrophysical origin. A sensitive test for the models is the cosmological evolutionary behaviour.  To evaluate the luminosity function and evolution behaviour of an object class complete samples with known distances of each individual object are necessary.  For BL\\,Lac objects two problems exist which render the determination of their luminosity function difficult. First, on account of the inconspicuous optical spectral properties of BL\\,Lacs (the absence of strong emission and absorption lines being one of the defining criteria) more effort is required in their selection than is the case with most other object classes. Although there are defining criteria suitable for optical selection (variability, polarisation) no optically selected sample of BL\\,Lac objects with more than 10 objects exists (Kollgaard, \\cite{kollgaard}). Second, their redshifts are difficult to determine for the same reason.  As a consequence, flux limited samples of BL\\,Lac objects with nearly complete redshift information are rare. Stickel et al. (\\cite{stickel}) selected a sample of 35 objects in the radio wavelength region with radio fluxes above 1\\,Jy. They found a flat redshift distribution between $0.1 < \\rm{z} < 1.2$ and their analysis of the sample revealed positive evolution; radio selected BL\\,Lac objects are more numerous in cosmological distances than in our local neighbourhood. At  X-ray wavelengths a complete well defined sample was built up with the EMSS (Stocke et al., \\cite{emss}). Morris et al. (\\cite{morris}) presented a sample of 22 objects down to $f_{X}(\\rm 0.3 - 3.5\\,keV) = 5\\cdot 10^{-13}\\,ergs\\,cm^{-2}\\,s^{-1}$, which was later expanded to 30 objects and lower fluxes by Wolter et al. (\\cite{wolter}). The EMSS objects exhibit a redshift distribution with strong concentration to low redshifts $\\rm z < 0.3$ and the sample shows clear signs for negative evolution. The EMSS sample was constructed with serendipitously found BL\\,Lac objects from pointed observations of the {\\em Einstein} satellite. Because most these pointed observations were of short exposure, objects with low fluxes are less frequent and are rare in the sample.  We contribute to this discussion with a new sample consisting of 39 X-ray selected BL\\,Lac objects with redshifts available for more than $80\\%$ of them. The selection is based on the ROSAT All-Sky Survey (RASS, Voges et al., \\cite{rass}). With a flux limit of $f_{X}(\\rm 0.5 - 2.0\\,keV) = 8\\cdot 10^{-13}\\,ergs\\,cm^{-2}\\,s^{-1}$ and redshifts up to $\\rm z \\sim 0.8$ the sample is suitable to test the evolutionary behaviour of the sample and to determine the luminosity function of RASS selected BL\\,Lac objects. In an earlier paper (Nass et al., \\cite{nass}) we presented a larger sample to discuss selection processes of BL\\,Lac objects within the RASS, but we could not deal with the evolutionary behaviour because the flux limit was not deep enough and many redshifts were unknown. The new sample consists of objects with a wide range of spectral energy distributions, from extremely X-ray dominated objects to objects with intermediate spectral energy distribution. Objects with a maximum of $\\nu f_{\\nu}$ in the radio or mm wavelength region, which are frequently found among the 1\\,Jy sample, are not contained in our X-ray flux limited sample. We were able to subdivide the sample into two subgroups according to \\alphaox\\footnote{we define \\alphaox\\ as the power law index between 1\\,keV and 4400\\,\\AA\\ with $\\rm f_{\\nu} \\propto \\nu^{-\\alphaox}$}. Several properties are different in these subgroups.  The following section describes the observational characteristics of the sample in the radio-, optical- and X-ray region. Details of the redshift determination are discussed, bcause it is a delicate process for BL\\,Lac objects. Section III analyses the properties of the sample, in particular the luminosity function and the evolutionary behaviour. In Sect. IV the new results  are discussed in the context of unified schemes for BL\\,Lac objects. Several new objects need comments to their redshift determination which are presented in an appendix.  We use cosmological parameters $\\rm H_{0} = 50\\,km\\,s^{-1}\\,Mpc^{-1}$ and $q_{0} = 0$ in this paper. ",
        "conclusions": "We have shown that the unified models do not predict the distinction between the two subgroups in our sample and additional free parameters would be necessary to adjust the model with the observations. Currently the observational constraints are of low statistical significance. Two populations of BL\\,Lac objects with different physical origin are also consistent with the results. The dividing point, $\\alphaox = 0.91$, is the median of the sample and has no special physical significance.  Our study has confirmed the negative evolution of X-ray selected BL\\,Lac objects and we find that this behaviour could be restricted to the extremely X-ray dominated BL\\,Lac objects. The cosmological evolutionary behaviour of the intermediate objects is consistent with no evolution. Since these two subgroups show a different redshift distribution the interesting evolution time of X-ray dominated BL\\,Lac objects is shifted outside to $\\rm 1.0 < z < 2.0$ (the $\\rm z_{max}$ from the $V/V_{max}$ analysis). There is still no answer for this odd behaviour, which up to now has not been found in other AGN classes. Possibly the jets in these objects, which are especially powerful in the X-rays, need a minimal time for development. This speculation has gained some plausibility since the analysis of our sample has shown that the relevant times lie farther away. Nevertheless X-ray dominated BL\\,Lac objects must have been considerably rarer between $\\rm 2 < z < 3$, the time when QSO activity was at its highest point. The distinct properties in the two subgroups are already revealed in the bright part of the $\\log N(> S) - \\log S$ distribution. It is therefore implausible that selection effects, as described by Browne \\& March\\~{a} (\\cite{browne}), are responsible for this distinction. Nevertheless, BL\\,Lac objects are ``susceptible'' to selection effects and therefore, in future, galaxy clusters should be carefully analyzed in the selection process.  The observations suggest the intermediate BL\\,Lac objects as the basic population. They have the lowest luminosity in X-rays and radio wavelengths, and they have the highest space density. Both the extremely radio and X-ray dominated BL\\,Lac objects have higher luminosities. This led us to speculate about a beaming scenario. The spectral energy distribution of extremely X-ray dominated BL\\,Lac objects can be interpreted with a high energy cutoff of the synchrotron spectrum. The high X-ray luminosity of objects with $\\alphaox < 0.91$ can be explained with large bulk Lorentz factor of relativistic electrons in the jets of these objects. This result is compatible with conclusions from multiwavelength observations of outbursts in Mrk\\,421 and Mrk\\,501. Therefore we suggest that extreme X-ray dominated objects are observed in a state of enhanced activity which would explain the anticorrelation between X-ray luminosity and \\alphaox.   Our investigation has shown the importance of the redshift parameter for the study of samples of BL\\,Lac objects. Advances in observing techniques have made easier the determination of redshifts in optical spectra which are devoid of strong features. In order to have stronger observational constraints it is necessary to enlarge the sample and to determine redshifts for objects for which this has not yet been possible. The uncertainty of the $V/V_{max}$ statistics depends on $1/\\sqrt{12N}$. In the last years several projects began to select BL\\,Lac objects with combined X-ray - radio methods (e.g. Wolter et al., \\cite{pilot}). They will certainly yield valuable results about intermediate objects and low flux BL\\,Lac objects below the adopted flux limit in this paper. However, the extremely X-ray dominated BL\\,Lac objects are rare and their negative evolution leads to a very flat $\\log N - \\log S$. Their distinct population properties are an important tool to understand the BL\\,Lac phenomenon, and they can be selected most preferably with the full sky coverage of the  RASS."
    },
    "9803/astro-ph9803032_arXiv.txt": {
        "abstract": "We have undertaken toward 30 Mira or semi-regular variables and one OH/IR object highly sensitive observations of the $v=1, J=2 \\rightarrow 1$ and $3 \\rightarrow 2$ transitions of SiO simultaneously with observations of the $J = 1 \\rightarrow 0$ transition of CO during three observing sessions in the period 1995 to 1996. As in our previous observations of 1994, we observe that for several stars the SiO profiles exhibit unusually broad wings which sometimes exceed the terminal velocity of the expanding circumstellar envelope traced by the thermal CO emission. We have discovered a clear dependence of the SiO wing emission on the optical phase. These wings are probably due to peculiar gas motions and varying physical conditions in relation with the stellar pulsation. However, we cannot exclude other mechanisms contributing to the observed wings. In particular, SiO turbulent motions for the semi-regular variables or the asymmetric mass loss mechanism may play a role. We conclude that the SiO wing emission is due to masing processes and that this emission very likely arises from the inner part of the circumstellar envelope. ",
        "introduction": "The SiO molecule exhibits widespread maser emission from O-rich circumstellar envelopes (CEs) around Long Period Variables (LPVs). The variety of rotational transitions emitted from several vibrational levels and the strength of the emission  make these lines very useful for a study of the innermost layers of CEs. Recently, during sensitive SiO observations of several O-rich late-type stars, we discovered unexpectedly broad wing SiO emission (Cernicharo et al. 1997, hereafter referred as CABG). We found that the $v=1, J=2 \\rightarrow 1$ SiO emission reaches and sometimes exceeds the maximum velocity traced by the quasi-thermal emission of the CO molecule. Several mechanisms may be invoked to explain such wings: turbulent motions, rotation of dense SiO clumps, high velocity shocks produced during the pulsation of the star, or high velocity bipolar ejection of gas from the star. These mechanisms were discussed by CABG, but no firm conclusions were reached, although it was recognized that the pulsation of the star and asymmetric mass loss could play an important role. Obviously interferometric observations and a monitoring of the SiO wing emission are needed to shed light on the location and the physical origin of the high velocity SiO emission.  The main purpose of this paper is to present and analyze new data gathered at different epochs on the SiO velocity wings in order to study the physical mechanisms responsible for this phenomenon, and, at the same time, to obtain deeper insight into the complex kinematics of the circumstellar shells around late-type stars. The latter question is clearly important since it is related to other crucial problems such as the expansion of the CE and the mass loss \\begin{figure*} [ht] \\begin{center} \\epsfxsize=15.cm \\epsfbox{figure1.epsf} \\end{center} \\caption []{Examples of SiO $v=1, J=2 \\rightarrow 1$ spectra for {\\em Mira variables} taken with $0.136$ {\\kms} spectral resolution in June 1995 (continuous line), April 1996 (dashed line) and October 1996 (dotted line). For each star the top panel shows the full main beam brightness temperature scale and the bottom panel corresponds to an enlargement of the same data, with the CO line emission width ($\\Delta V(CO)$ above the $2\\sigma$ level) represented by a horizontal segment. For R Leo, no data are available for April 1996. To transform the intensities into Jy one must multiply by $4.4$ Jy/K.} \\label{asio_spectra} \\end{figure*} phenomenon, or the processes leading to the formation of dust. We have considered a rather small but homogeneous sample of stars as it includes Miras and a few semi-regulars for which all stellar characteristics (mass loss, temperature, spatial distribution \\ldots) are uniformly represented. We have monitored the SiO maser line profiles in order to investigate the relation of broad (weak) emission in wings with the stellar light phase. In particular, we wish to test whether the shocks driven by the stellar pulsation could play a dominant role in the occurrence of SiO high velocity features. In Sects. 2 and 3 we give details of our observations and present our main results. In Sect. 4 we discuss the different hypotheses which could explain the SiO line wings and their location. \\newpage ",
        "conclusions": "We have observed a rather large sample of late-type stars including Miras and semi-regulars with the IRAM 30-m radiotelescope at four epochs covering the period 1994 to 1996 in order to investigate the correlation between the SiO linewing activity and the stellar light phase. The SiO $v=1, J=2 \\rightarrow 1$ and $J=3 \\rightarrow 2$ lines were observed simultaneously with the CO $J=2 \\rightarrow 1$ quasi-thermal emission line. Several high velocity wings have been detected in the red and blue edges of the SiO profile.  The SiO wing emission could result from complex mechanisms combining stellar pulsation of the fundamental mode, and perhaps of other modes, with asymmetric mass loss (as for R Leo) and structure changes. We deduce from our observations that the time evolution of the SiO line wings is related to the stellar pulsation. However, it is difficult to specify how the pulsation induces local physical variations responsible for variations of the SiO wing emission. Pulsation, through shocks, may produce high velocity emissions which then vary according to the optical phase. For the semi-regular variables there is some indication on the importance of turbulent motions in the formation of the high velocity emission as well.  Finally, we have shown that the SiO wing emission results from masing processes and that this emission very likely arises from the innermost gas layers of the circumstellar envelope."
    },
    "9803/astro-ph9803342_arXiv.txt": {
        "abstract": "We present ultra-high resolution (0.32 km s$^{-1}$) spectra obtained with the 3.9m Anglo-Australian Telescope (AAT) and Ultra-High-Resolution Facility (UHRF), of interstellar Na~{\\sc i} D$_1$, Na~{\\sc i} D$_2$, Ca~{\\sc ii} K, K~{\\sc i} and CH absorption toward two high galactic latitude stars HD~141569 and HD~157841. We have compared our data with 21-cm observations obtained from the Leiden/Dwingeloo H~{\\sc i} survey. We derive the velocity structure, column densities of the clouds represented by the various components and identify the clouds with ISM structures seen in the region at other wavelengths. We further derive abundances, linear depletions and H$_2$ fractional abundances for these clouds, wherever possible.  Both stars are located in regions of IRAS 100$\\mu$m emission associated with high galactic latitude molecular clouds (HLCs) : HD~141569 lies, in projection, close to MBM~37 and the Lynds dark cloud L~134N while HD~157841 is in the vicinity of the MBM~151.  Toward HD~141569, we detect two components in our UHRF spectra : a weak, broad $b$ = 4.5 \\kms\\ component at -- 15 \\kms\\ , seen only in Ca~{\\sc ii} K absorption and another component at 0 \\kms\\ , seen in Na~{\\sc i} D$_1$, Na~{\\sc i} D$_2$, Ca~{\\sc ii} K, K~{\\sc i} and CH absorption. The cloud represented by the -- 15 \\kms\\ component, is warm and may be located in a region close to the star. The cloud represented by the 0 \\kms\\ component has a Ca linear depletion $\\delta$(Ca) = 1.4 $\\times$ 10$^{-4}$ and shows evidence for the presence of dust, consistent with strong 100$\\mu$m emission seen in this region. The H$_2$ fractional abundance  $f$(H$_2$), derived for this cloud is 0.4, which is typically what is observed toward HLCs. We conclude that this 0 \\kms\\ cloud is associated with MBM~37 and L~134N based on the presence of dust and molecular gas (CH) and good velocity agreement with CO emission from these two clouds. This places HD~141569 beyond MBM~37 and L~134N, which are estimated to be at $\\approx$ 110 pc.  In the case of the HD~157841 sightline, a total of 6 components are seen on our UHRF spectra in  Na~{\\sc i} D$_1$, Na~{\\sc i} D$_2$, Ca~{\\sc ii} K, K~{\\sc i} and CH absorption. 2 of these 6 components are seen only in a single species. The cloud represented by the components at 1.85 \\kms\\ has a Ca linear depletion $\\delta$(Ca) = 2.8 $\\times$ 10$^{-4}$, indicating the presence of dust. The $f$(H$_2$) derived for this cloud is 0.45 and there is good velocity agreement with CO emission from MBM~151. To the best of our knowledge, this 1.85 \\kms\\ component towards HD~157841 is the first one found to have relative line widths that are consistent with pure thermal broadening only. We associate the 1.85 \\kms\\ cloud seen in our UHRF spectra with MBM~151 and conclude that HD~157841 must lie beyond $\\sim$ 200 pc, the estimated distance to MBM~151. ",
        "introduction": "In this paper, we report the results of an ultra-high-resolution (0.32 \\kms) absorption line study of two stars, HD~141569 and HD~157841. Both stars were chosen as background sources for detailed spectroscopic studies of high galactic latitude molecular clouds (HLCs), as well as gas in the nearby halo.   HLCs show structure, with dense clumps as small as 0.03 pc, which appear to be virially unstable (Pound et al. 1990). Some of these clouds appear to be embedded in the local hot interstellar medium as suggested by the probable detection of X-ray shadowing from two MBM clouds (Burrows and Mendenhall, 1991). If this is indeed the case, it is not clear what physical process or combination of processes is causing such enhanced molecular abundances. This has led some authors (Blitz, 1990) to conclude that HLCs may represent a completely distinct population of clouds. More recent work by Gir et al. (1994) indicates that HLCs are almost always located along filamentary or looplike HI structures and appear to have condensed from atomic gas in situ rather than having been entrained in the HI.   These clouds are an important factor both in the composition and the energy balance of the nearby interstellar medium. They have molecular masses about 50 M$_{\\odot}$ and the molecular gas contribution of these high latitude clouds in the solar neighborhood up to 100 pc, is about 5000  M$_{\\odot}$ (Magnani 1993). The relation and exchange of energy between the molecular gas in these clouds and the ambient highly ionized, moderately ionized and atomic gas is unclear although it has been suggested by Blitz (1990) that compression of these clouds could lead to a phase transition of atomic hydrogen into molecular hydrogen.   Absorption line studies toward stars whose lines of sight intercept regions close to HLCs, probe the diffuse component of the ISM associated with these high latitude clouds. Typical values of the kinetic temperatures for the diffuse component of the ISM range between 45 to 130 K (Savage et al. 1977). To resolve absorption lines arising from clouds with cloud kinetic temperatures of \\Tk\\ = 50 K, a  velocity resolution of 0.30 \\kms\\  corresponding to a resolving power R = 10$^6$ is necessary.  The Ultra High Resolution Facility (UHRF)  at the Anglo Australian Telescope (AAT) is currently the world's highest-resolution optical astronomical spectrograph (Diego et al., 1995, Barlow et al., 1995) and is ideal for studies of such cool, interstellar clouds.  We have begun a study of the spatial distribution and properties of atomic and molecular interstellar species towards HLCs at ultra high resolution (velocity resolution = 0.32 \\kms\\ ), and have compiled a list of background stars  with reliable photometric and spectroscopic data and lying in the line of sight to HLCs.   We have already observed a small fraction of this sample, utilizing the UHRF (Blades et al. 1997), and in the following sections, we present UHRF observations toward HD~141569 and HD~157841.  \\subsection{Overview of the HD~141569 sightline}  HD~141569 has been variously classified as an isolated Herbig Ae/Be star (Th$\\acute{e}$ et al. 1994) and a B9.5V star by Penprase (1992) and an AOVe star by Dunkin et al. (1997). It is an emission-line star associated with an IRAS source, IRAS~15473-0346, which is not exactly coincident with it. The heliocentric radial velocity of the star is --6.4 \\kms\\ (Frisch, 1987). In the direction of HD~141569, which is located in galactic coordinates at ($l$, $b$) = (4.$^\\circ$2, +36.$^\\circ$9), this translates into a LSR radial velocity of --20.1 \\kms\\ . The star exhibits excess far-IR emission (Oudmaijer et al. 1992) believed to arise due to emission from dust grains present in a circumstellar disk, and Sylvester et al. (1996) include it in their list of Vega-like systems.  A high galactic latitude dark cloud complex at ($l$, $b$) $\\sim$ (4$^\\circ$, 36$^\\circ$) lies close to the star in projection. This cloud complex includes the Lynds dark clouds L~134, L~183 (which is also often referred to as L~134N; henceforth we will refer to this cloud as L~134N) and L~1780. The dense molecular CO core MBM~36 is located within L~134, while the CO core in L~134N is referred to as MBM~37 (Magnani et al. 1985). Of these two dark clouds, L~134N, centered at ($l$, $b$) = (6$^\\circ$.0, 36$^\\circ$.8) and 2$^\\circ$ north of L~134, lies closest to HD~141569. HD~141569 lies well within the 100$\\mu$m emission contours of the L~134N dark cloud (Laureijs et al. 1991).  Penprase (1992) estimated a distance of 190$\\pm$110 pc for HD~141569 based on spectroscopic and UBV and Str\\\"{o}mgren photometric data. There are various distance estimates to the L134, L134N and L1780 complex: 100$\\pm$50 pc, based on optical extinction and surface brightness observations by Mattila (1979), 110$\\pm$10 pc based on photometric data by Franco (1989) and 160 pc  based on reddenings of a large number of stars in this direction by Snell (1981). The parallax of HD~141569 from the Hipparcos Catalogue is 10.10$\\pm$0.83~mas, corresponding to a distance of 99$\\pm$8~pc. This implies that the distance to the star is very similar to that of the L~134, L~134N and L~1780 complex. Penprase (1992) estimated an E(B--V) of 0.12 $\\pm$ 0.09 for HD~141569. The Hipparcos Catalogue lists Tycho magnitudes which transform to give (B--V) = 0.078$\\pm$0.007, which for an A0V star corresponds to an E(B--V) of 0.10$\\pm$0.03 (the uncertainty of 0.03 is due to the uncertainty in the (B--V)$_{o}$ calibration for dwarf stars). The star falls just outside the outermost CO contours which trace MBM~37, the CO core in L~134N (Caillault et al. 1995). Based on available information, although HD~141569 appears to lie approximately at the distance of L134N and clearly does not lie within the dark cloud itself, it is not certain whether it lies in front of, or behind this dark cloud. This point is discussed further in $\\S$ 3.1.3. and our UHRF data indicate that HD~141569 lies beyond L~134N and MBM~37.  \\subsection{Overview of the HD~157841 sightline}  HD~157841, a B9 star (SIMBAD database) at ($l$, $b$) = (16$^\\circ$.8, 15$^\\circ$.7), lies in projection near the HLC, MBM~151, which is located at ($l$, $b$) = (21$^\\circ$.5, 20$^\\circ$.9). We have searched the SIMBAD database and related literature and have been unable to find any radial velocity measurement for this star.\\\\  Based on CO velocity agreement, Penprase (1993) suggested that MBM~57 located at ($l$, $b$) = (5$^\\circ$.1, 30$^\\circ$.8) is related to MBM~151, which is $\\sim$ 19$^\\circ$ away. Photometric data of stars in the region around MBM~57 (Franco, 1989), show an increase in extinction beyond 100 pc. Penprase (1992) also observed this increase in the reddening and estimated a distance of 150 to 210 pc for the near edge of MBM~57. Assuming that MBM~57 and MBM~151 are related, the distance to MBM~151 is $\\sim$ 200 pc. The CO (J= 1$\\rightarrow$0) LSR velocity of MBM~151 is --0.8 \\kms\\ (Magnani et al., 1985). The region near HD~157841 has not been mapped in CO and it is not known if there is any CO emission present at the position of the star or in its vicinity. The Hipparcos Catalogue lists a parallax for HD~157841 of 5.70$\\pm$1.16~mas corresponding to a distance of 175$\\pm$$\\stackrel{45}{_{30}}$~pc. It also lists Tycho magnitudes which transform to yield (B--V) = 0.213$\\pm$0.008. For a B9V star, this corresponds to an E(B--V) of 0.28$\\pm$0.03. In $\\S$ 3.2.5., we identify the various components seen in the UHRF spectra toward HD~157841. ",
        "conclusions": "We have made an ultra-high resolution study of the Na~{\\sc i}, Ca~{\\sc ii} K, K~{\\sc i} and CH interstellar absorption lines toward two stars, HD~141569 and HD~157841. These absorption spectra have been compared to 21-cm data obtained from the Leiden/Dwingeloo H~{\\sc i} survey. Both stars probe the gas in regions close to high galactic latitude molecular clouds : the HD~141569 sightline intercepts a region close to MBM~37 and L~134N while the HD~157841 sightline probes a region close to MBM~151.   The results of our investigation are as follows.  \\noindent 1) Toward HD~141569, two components are seen in our UHRF spectra : one at --15 \\kms\\ and another at 0 \\kms\\ . The --15 \\kms\\ component, seen only in Ca~{\\sc ii} K, is weak, broad and shows low linear depletion, $\\delta$(Ca) $>$ 0.12. In the absence of turbulence, the kinetic temperature derived is \\Tk\\ $\\le$ 47,900~K and the cloud represented by this component may be located in a region close to HD~141569, whose stellar radial LSR velocity is --20.1 \\kms\\ .   The $\\sim$ 0 \\kms\\ component is seen in Na~{\\sc i}, Ca~{\\sc ii} K, K~{\\sc i} and CH absorption. The cloud represented by these components has a linear Ca depletion of $\\delta$(Ca) = 1.4 $\\times$ 10$^{-4}$, implying the presence of dust, consistent with strong 100 $\\mu$m emission from this region, and H$_2$ fractional abundance $f$(H$_2$) = 0.4. We conclude that this  0 \\kms\\ cloud is associated with MBM~37 and L~134N because of the presence of dust and molecular (CH) gas at velocities close to CO velocities derived for the two HLCs. The Hipparcos distance of 99$\\pm$8~pc for HD~141569 is comparable to the distance of 100 -- 110~pc that has been estimated for the MBM~37 and L~134N clouds. However, our absorption line data imply that these clouds must lie a little closer than HD~141569.   \\noindent 2) A total of 6 components are seen in our UHRF spectra in Na~{\\sc i}, Ca~{\\sc ii} K, K~{\\sc i} and CH absorption toward HD~157841. The cloud represented by the 1.85 \\kms\\ absorption components has a Ca linear depletion $\\delta$(Ca) = 2.8 $\\times$ 10$^{-4}$ and the derived H$_2$ fractional abundance is $f$(H$_2$) = 0.45. The relative $b$-values determined for the 1.85 \\kms\\ Na~{\\sc i}, Ca~{\\sc ii} and K~{\\sc i} absorption lines are found to be consistent with thermal broadening (with no significant turbulent component), corresponding to a kinetic temperature of \\Tk\\ = 5700$\\pm$500~K. To our knowledge, this is the first velocity component found whose line widths can unambiguously be attributed to pure thermal broadening. We conclude that the 1.85 \\kms\\ cloud is associated with MBM~151 because of the presence of dust and molecular (CH) gas at velocities close to CO velocities derived for MBM~151. HD~141569 must therefore lie beyond MBM~151, estimated to be at $\\sim$ 200 pc and its Hipparcos distance of 175$\\pm$$\\stackrel{45}{_{30}}$~pc is consistent with this conclusion."
    },
    "9803/astro-ph9803174_arXiv.txt": {
        "abstract": " ",
        "introduction": "The cluster Abell~370 is a distant ($<z>=0.374$~\\cite{M88}), rich, massive cluster, with two central giant galaxies dominating its optical image. The X-ray luminosity of the cluster is\\footnote{We assume a Hubble constant $H_0=75$~km~s$^{-1}$~Mpc$^{-1}$ and a deceleration parameter $q_0=1/2$ throughout this paper, giving a luminosity-distance to the cluster of 1613~Mpc.} $3.7 \\times 10^{44}$ erg~s$^{-1}$ \\cite{Henry}, which is as expected from the high value of the velocity dispersion of this cluster, 1340~km~s$^{-1}$\\cite{M88}.  A370 was shown by \\cite{Bautz}, \\cite{BO84} and \\cite{M88} to contain an anomalous fraction of blue objects, as compared to nearby clusters (50~\\% of the cluster members in the central region show evidence for star-formation activity, as compared to a mere 3~\\% in Coma). Based on optical and near-IR photometry, it has been shown (\\cite{Aragon}, \\cite{Stanford}, \\cite{McLean}) that most galaxies in the cluster are reasonably well fitted by passive evolution models, while the existence of galaxies in a post-starburst phase is controversial (\\cite{Aragon} vs. \\cite{Stanford}). The spectral energy distribution of the cluster galaxies must be known over a wider wavelength range before their nature can be firmly established, and in this context ISO observations are very important.  {\\psfig{file=a370_opt_c.ps,width=15.cm,angle=0}} {\\bf Fig.1.} I-band optical image of A370 (from \\cite{Kneib}); the original image has been transformed to match the ISOCAM image field-of-view. The coordinates are arcsec. The strange feature in the upper part of the figure is straylight from a very bright star off the image. \\medskip  A370 contains the first spectroscopically confirmed gravitational arc, A0 at $\\mbox{z=0.724}$~\\cite{Soucail}. Other arcs have since been detected in the cluster (\\cite{Kneib} and references therein). Based on {\\em Hubble Space Telescope} images, \\cite{Smail} found evidence for a faint spiral structure and an apparent bulge. The spiral nature of the A0 arc was also indicated by its optical and near-IR colours and it was detected by its CO line in emission by \\cite{Casoli}.  In this paper we report on the mid-IR observations of A370 done with ISOCAM onboard ESA's {\\em Infrared Space Observatory} (ISO) satellite. A370 and another three clusters were selected for observation in the context of an ISO guaranteed-time project aimed at generating mid-IR images of known giant arcs and providing their mid-IR fluxes.  In \\S~2 we give a description of our observations and data-reductions; in \\S~3 we present the ISOCAM image, and compare it with an I-band image in order to derive a colour map. We give our conclusions in \\S~4. ",
        "conclusions": "In this paper we have presented the results of observations of A370 with ISOCAM onboard ISO. We have detected the giant arc at 7 and 10 $\\mu$m; this is the first detection ever of a gravitational arc in a galaxy cluster at these wavelengths. Moreover, we have also detected many other galaxies, in the cluster field.  The comparison of an I-band image and the ISOCAM 7 $\\mu$m image in the LW2 filter has allowed us to build a colour-map which seems to indicate a starbursting phase for the gravitationally lensed galaxy. Most cluster galaxies however appear to have a ratio of 7 $\\mu$m to 0.9 $\\mu$m flux density as predicted from spectral energy distribution models of normal ellipticals \\cite{Mazzei1}.  While our data reduction methodology is still being developed and refined, and while a full exposition of quantitative results remains to be achieved, the present results already demonstrate the capability of ISO in detecting faint mid-IR sources at large distances. With our ongoing program we hope to detect other gravitational arcs in clusters, and so to constrain the evolutionary status of these high redshift sources. We will widen the wavelength coverage of the spectral energy distributions for cluster galaxies, thus providing a better leverage for constraining models of galaxy formation and evolution. An additional exciting perspective of our future observations is the serendipitous discovery of optically undetected \"IR-only\" arcs or arclets."
    },
    "9803/astro-ph9803319_arXiv.txt": {
        "abstract": "The dispersion in the peak luminosities of high redshift type Ia supernovae will change with redshift due to gravitational lensing.  This lensing is investigated with an emphasis on the prospects of measuring it and separating it from other possible sources of redshift dependent dispersion.  Measuring the lensing induced dispersion would directly constrain the power spectrum of density fluctuations on smaller length scales than are easily probed in any other way.  The skew of the magnification distribution is related to the bispectrum of density fluctuations.  Using cold dark matter models it is found that the amount and quality of data needed is attainable in a few years.  A parameterization of the signal as a power law of the angular size distance to the supernovae is motivated by these models.  This information can be used in detecting lensing, detecting other systematic changes in supernovae and calculating the uncertainties in cosmological parameter estimates. ",
        "introduction": "There is presently a large effort underway to predict and detect the weak gravitational lensing caused by Large Scale Structure (LSS) or the ``cosmic shear'' (see \\cite{valdes83,mould94,vill96,kais92,kais96}).  Such a measurement would constitute a direct probe of the mass density fluctuations on large scales irrespective of how light and baryons are distributed.  This would make it possible to measure one of cosmology's least well understood processes, how light traces mass and whether this is a function of scale.  Gravitational lensing causes images to be both magnified and demagnified as well as stretched asymmetrically (shear).  Most of the methods proposed for detecting LSS lensing are based on measuring the shear in high redshift galaxy images.  Although the distortions in the ellipticities of individual galaxies are expected to be small they are distorted in coherent ways.  Indications of lensing are sought in the alignment of the galaxy images with the assumption that they are not intrinsically aligned.   This technique has already been used with great success on galaxy cluster lensing and is presently being applied to random fields in an attempt to detect the lensing effects of large scale structure.  In this paper a method of detecting lensing directly through its magnification rather than shear is proposed.  The study of high redshift supernovae (SNe) is another area of cosmology and astrophysics that is seeing a lot of activity.  Type Ia SNe, the brightest type of supernova (SN), are believed to be caused by the thermonuclear explosion of an accreting white dwarf. It has been found empirically that the peak magnitude of type Ia's have a dispersion of only 0.2 - 0.3 mag in B band.  It has further been found that the peak magnitude is related to the width of the SN's light-curve which can then be used to reduce the dispersion to about 0.17 mag if one color is used \\cite{hamuy96} and 0.12 mag if multiple colors are used \\cite{RPK96}.  In addition to the light-curve width, the SN's color and spectral features are related to the peak luminosity \\cite{BNF97,nugent95}.  It may be possible to reduce this dispersion in the future by incorporating additional observables into the correction procedure.  The uniformity in type Ia SNe combined with their high luminosity makes them an excellent tool for doing cosmology.  Using them to measure the redshift-luminosity distance relation has recently resulted in tightened constraints on the cosmological parameters $\\Omega$ and $\\Omega_\\Lambda$ (\\cite{Perl97,Perl98,garn98,Riess98}).  There are now systematic searches for Ia SNe at high redshift which can reliably discover on the order of ten SNe in a night's observing and do spectroscopic followup (\\cite{Perl97}).  In addition there are several ongoing searches for low redshift SNe.  To date there have been more then 100 type Ia SNe discovered with redshifts between $z=0.4$ and $0.97$ and many additional ones at lower redshift.  In this paper I concentrate on the lensing produced by dark matter composed of microscopic particles.  Matter in macroscopic compact objects can cause microlensing.  However, the known populations of stars will not cause a significant number of microlensing events.  The microlensing of SNe has been discussed by \\cite{SW87} and \\cite{LSW88}. A SN could also be lensed by one dominant galaxy cluster or individual galaxy.  There is the possibility of getting multiple observable images and high magnifications (strong lensing) in this case.  It has been suggested that observing SNe behind galaxy clusters would be a way of lifting the mass sheet degeneracy that exists in shear measurements of the gravitational lensing \\cite{kolatt97}.  The likelihood of a SN at $z=1$ being strongly lensed by a galaxy or cluster is small unless they are specifically sought out.  In general a SN will be lensed by many galaxies and larger structures, each have a weak contribution to the total magnification.  This will increase the dispersion of high redshift SNe and decrease the precision of cosmological parameter determinations.  This decrease in precision has been investigated by \\cite{Frieman97} using analytic methods and by \\cite{wamb97} using N-body simulations.  \\cite{HW98} calculates the lensing of point sources using numerical simulations which assume that all matter is in unclustered galaxy halos.  \\cite{Kant98} (and \\cite{Kant95}) does analytic calculations of this effect under the assumption that all matter is in compact objects and none of them are close enough to the line of sight to a SN to cause significant lensing. In this paper the problem will be turned around and we will ask how well the lensing itself can be determined from the SN data.  There are several important differences between the lensing of SNe and the lensing of galaxies. For SNe the signal to noise ratio in each measurement can in fact be greater.  Lensing is estimated to contribute about $5$ to $10\\%$ to the observed root-mean-squared ellipticity of a $z=1$ galaxy.  For SNe the variance in the lensing contribution is comparable to the intrinsic variance in the peak magnitude at this redshift. In addition the unlensed ellipticity distribution of galaxies at high redshift is not known and can not be easily extrapolated from zero redshift galaxies.  For this reason lensing must be inferred by correlations between galaxies, either between lensed galaxies or between lensed galaxies and foreground galaxies.  The result is that the lensing of galaxies measures the shear averaged over a finite area on the sky.  This average shear drops rapidly with increasing area and the signal to noise is reduced to something more like $1\\%$ per galaxy on the $1\\mbox{ deg}^2$ scale.  This is made up for by the large number of evaluable galaxies ($\\sim 10^5\\mbox{ deg}^{-1}$) to the extent that fluctuations in the projected mass density at $1-100$~arcmin scales are expected to be detectable in the near future.  In contrast, the dispersion in the absolute magnitude of type Ia SNe is presumably independent of redshift.  The dispersion can then be measured in a low redshift population where lensing is not important.  The lensing of galaxies does have the advantage of sources that are generally at higher redshifts where the lensing is stronger.  On the other hand, the redshift distribution of faint galaxies is not strongly constrained which adds systematic uncertainty.  The redshifts of SNe are individually measured.  It is clear that the lensing of SNe and the lensing of galaxies probe different scales of density fluctuations for several reasons.  Because the lensing of galaxies can only be detected through correlations in their shear, or positions and shear, lensing structures that are smaller than roughly the separation between galaxies are not detectable.  Since each SN is an independent measure of the lensing at a point, not an average over a region on the sky, they will be sensitive to smaller scale structures.  Also the lensing of SNe is a direct measurement of the magnification which, in the weak lensing limit, is directly related to the mass density along the line of sight.  The shear is dependent on the mass outside of the ``beam'' and thus a shear map is in a way a smoothed version of a surface density map.  In the thin lens approximation the shear and the magnification are related to each other through a differential operator \\cite{KS93} which makes the shear insensitive to a uniform offset in the surface density - the ``mass sheet degeneracy''. Although the lensing of galaxies has limitations on small scales it does have the potential of measuring lensing over a large range of scales, arcminutes to degrees. With the possible exception of SNe viewed through galaxy clusters, it will be difficult to find enough SNe that are close enough together to measure correlations in their lensing.  In the next section I describe first how the lensing of SNe is related to the density fluctuations and then I describe how the lensing signal could be identified in the data. In section~\\ref{models} a specific model of structure formation is used to estimate the level of signal and to motivate some parameterizations.  The last section contains conclusions and remarks about possible complications. ",
        "conclusions": "It has been shown that measuring the gravitational lensing of type Ia SNe is feasible if the noise can be reasonably constrained.  It would be best to solve for the best fit cosmological parameters (ie. $\\Omega$, $\\Omega_\\Lambda$), lensing strength (ie. $\\eta_o$) and intrinsic noise ($\\sigma_M$) simultaneously using SNe at all redshifts.  The photometric uncertainties should be comparatively well determined for each SN.  One could then marginalize over the intrinsic variance, $\\sigma_M$.  The greatest worry is of course that the type Ia SNe properties or their galactic environments are changing with redshift.  Observations of spectral features and colors (\\cite{Perl97}) suggest that this is not the case, but the possibility remains.  It is possible that a systematic change in the metallicity of progenitors could change both the average peak luminosity and its dispersion.  Another worry is that extinction corrections change with redshift.  This could systematically reduce $\\overline{m}(z)$, increase its dispersion and make its distribution non-Gaussian.  Extinction should be accompanied by reddening, which can be detected, but the extinction law is not certain. These changes would affect cosmological parameter parameter estimates as well as lensing estimates.  The methods discussed here can be directly applied to detecting any redshift dependent change in the dispersion.  In testing for possible redshift evolution lensing must be incorporated.  There are difficulties in searching for SNe at higher redshifts.  The region of the spectrum that is used to do the light curve correction passes out of the visible at $z\\gtrsim 1$.  To go to significantly higher redshift may require switching to the IR. There is also difficulties with the K-correction and the subtraction of atmospheric lines. But with this in mind it seems that gravitational lensing of SNe could be detectable in the next few years when hundreds of high redshift SNe are observed and systematic effects are better understood.  CCD cameras with fields of view approaching a square degree and very small pixel sizes are being built now.  They will be used for weak lensing measurements using galaxy shear. SNe searches could be incorporated into these surveys with the benefits of improved cosmological parameter estimation and complimentary weak lensing measurements.  Combined with the limits on the cosmological parameters the lensing of SNe can constrain the power spectrum of the true mass density, unbiased by the light distribution, on the scale of galaxy halos.  At present the mass distribution on these scales is not well known with the exception of within galaxy clusters which are certainly atypical regions.  In this paper it has been assumed that the the large majority of matter in the universe is in the form of WIMPS or some other small particle.  If the dark matter is in compact objects like MACHOs the distribution of magnifications will be different.  In this way the lensing of SNe also provides information on the composition of dark matter. In addition to the variance of the magnification distribution the skewness would provide an important constraint on the nature of structure in the nonlinear regime.  A future paper will treat this subject in more detail and relate the magnification distribution to the nature of dark matter and the structure of galaxy halos."
    },
    "9803/astro-ph9803296_arXiv.txt": {
        "abstract": "We present new observations of the isolated young stars \\hd\\ and \\cd.  Pointed \\rosat\\ observations show that their X-ray properties, including X-ray luminosity and variability, are consistent with those of \\pms\\ (PMS) stars.  These observations do not reveal any additional PMS candidates in 40\\arcmin\\ fields centered on \\hd\\ and \\cd. Hipparcos observations of TW Hya (Wichmann \\etal\\ 1998\\mc{w98}) and \\hd\\ (Soderblom \\etal\\ 1998\\mc{s98}) show that both stars are roughly 50 pc away and are PMS with ages of $\\sim 10^7$ yr.  We searched the Hipparcos catalog (complete down to $\\sim$ 2--3 $L_\\odot$ at this distance) for other PMS stars in the same area. In a 10-pc radius volume of space centered on the previously known PMS stars, we find one additional candidate PMS star (CD $-$36$\\arcdeg$7429) with a low space velocity, X-ray emission comparable to that of \\hd, and Li absorption.  There are eight other stars in this area that have dwarf spectral types and lie above the main sequence, but based on their weak X-ray emission, high space velocities, and lack of Li in low-resolution spectra (i.e.\\ EW(Li) $< 0.1$ \\AA), these are probably mis-classified subgiants or giants. The current positions and proper motions of TW Hya, \\hd, and CD $-$36$\\arcdeg$7429 are inconsistent with them having formed as a group. ",
        "introduction": "Recent observations have revealed a small population of stars that bear many of the hallmarks of low-mass pre--main-sequence stars but lie far from any obvious region of recent star formation (as revealed by substantial dark clouds of molecular gas and dust).  Gregorio-Hetem \\etal\\ (1992\\mc{gh92}) identified 33 candidate T Tauri stars based on spectroscopy of stars in the {\\it IRAS\\/} Point Source Catalog and a few additional emission-line stars.  One of the more interesting findings of their work was the identification of four stars (HD 98800, \\cd, \\cdb, and \\hen) within 10\\arcdeg\\ of TW Hya, earlier identified as a possible isolated T Tauri star by Rucinski \\& Krautter (1983\\mc{rk83}).  The proximity of these five systems to each other suggested a possible loose cluster of young stars, possibly formed by a small molecular cloud that has since dissipated (e.g., Feigelson 1996\\mc{f96}).  In the absence of association with any known cloud complex, the distances to these stars were very uncertain, and thus their pre--main-sequence status was in question. To investigate whether two of these stars, \\hd\\ and \\cd, are in fact young and whether they are part of a larger group of young stars, we observed them with the \\rosat\\ X-ray satellite.  One distinguishing feature of low-mass pre--main-sequence stars is strong X-ray emission, presumed to arise from solar-like chromospheric activity (see, e.g., \\rne\\ 1997\\mc{n97} for a recent review).  Thus, X-ray observations can be used to search for young stars that may not have other hallmarks of youth such as infrared excesses or strong H$\\alpha$ emission. Observations of Taurus-Auriga with the {\\it Einstein} satellite led to the discovery of the naked T Tauri stars, stars that have no circumstellar material but that are nonetheless coeval with classical T Tauri stars (Walter 1986\\mc{w86}).  Much observational attention has been focused on the stars in the vicinity of TW Hya recently, and several other investigations have proceeded in parallel with ours.  Kastner \\etal\\ (1997\\mc{k97}) reported \\rosat\\ observations of some of the same stars we report on here.  They found that the strength of X-ray emission from these stars is consistent with them being young stars and argued that the five young stars in this area make up a physical association.  Hoff \\etal\\ (1996\\mc{hph96}, 1998\\mc{hhp98}) investigated TW Hya and \\cdb\\ and their surroundings with low spatial resolution using \\rosat\\ PSPC pointed observations.  Soderblom \\etal\\ (1998\\mc{s98}) and Wichmann \\etal\\ (1998\\mc{w98}) report Hipparcos distances to HD 98800 and TW Hya, respectively; both systems have ages of $\\sim$ 1--2 $\\times\\ 10^7$ yr and distances of $\\sim 50$ pc.  The work we report here is complementary to this other recent work. We analyze the X-ray data of our target stars in more detail than Kastner \\etal\\ (1997\\mc{k97}) and explore the status of other X-ray sources in the surrounding fields. We also combine X-ray and Hipparcos data to search for other young stars in the vicinity in order to address the larger question of the origin of these isolated young stars.  In \\sec\\ \\ref{sec:observations} we report the details of our observations. In \\sec\\ \\ref{sec:xproperties} we show that the X-ray properties of the two target stars are consistent with their being \\pms\\ (PMS) stars.  We then investigate (\\sec\\ \\ref{sec:othermembers}) whether there is any evidence that the young stars near TW Hya formed as a group, using our X-ray observations as well as data from the Hipparcos satellite.  Finally, in \\sec\\ \\ref{sec:discussion} we discuss the implications of our findings for our understanding of isolated young stars. ",
        "conclusions": "\\label{sec:summary}  We have presented pointed \\rosat\\ observations of the isolated young stars \\hd\\ and \\cd.  Their X-ray fluxes and X-ray variability are consistent with them being \\pms\\ stars.  No other \\pms\\ stars were found among X-ray sources within 40\\arcmin\\ of these stars.  Hipparcos observations of the area reveal the presence of nine additional stars that lie above the main sequence in a volume roughly 10 pc in diameter centered on the five previously known young stars. The X-ray properties and kinematics of these stars indicate that one of them (CD $-$36$\\arcdeg$7429) is quite likely a PMS star with an age of $10^7$ yr, while the others are more likely post--main-sequence stars. Observations of Li abundances in these stars confirm this conclusion.  Comparison with other fields at the same Galactic latitude shows that other fields selected in the same way show similar numbers of stars above the main sequence.  However, the field around TW Hya has significantly more stars with low space velocities (less than 5 km s\\per\\ relative to the LSR) than the other fields.  Thus, there is some indication of an excess of young stars in this area. Nonetheless, the proper motions of \\hd, TW Hya, and CD $-$36$\\arcdeg$7429, if extrapolated back in time, do not indicate a common place of origin.  \\noindent \\centerline{\\bf Note added in proof:}  Our analysis of whether HD 98800, CD $-$36$\\arcdeg$7429, and TW Hya could have formed as a group neglected an important source of uncertainty. The uncertainty in distance to these stars, while it does not affect the observed proper motions, does introduce uncertainty into the solar motion that must be subtracted in order to determine the stars' intrinsic proper motions with respect to the LSR.  The 1$\\sigma$ effect of this uncertainty is shown by the dashed lines in Figure 3. Including this uncertainty indicates that TW Hya and CD $-$36$\\arcdeg$7429 could plausibly have formed together; only a $1.5\\sigma$ error on their distances is required. Regarding HD 98800, bringing its current distance $2\\sigma$ closer, to 36.8 pc, would allow its projected point of origin to overlap with those of the other two stars.  However, this distance makes HD 98800 a main-sequence star, in conflict with its observed properties and, more importantly, of a very different age than the other two stars.  We conclude that HD 98800 is unlikely to share a common origin with TW Hya or CD $-$36$\\arcdeg$7429."
    },
    "9803/astro-ph9803225_arXiv.txt": {
        "abstract": "Results obtained from 9 X-ray observations of 3C 273 performed by {\\it ASCA} are presented (for a total exposure time of about 160 000 s). The analysis and interpretation of the results is complicated by the fact that 4 of these observations were used for on-board calibration of the CCDs spectral response. In particular, we had to pay special attention to the low energy band and 5--6 keV energy range where systematic effects could distort a correct interpretation of the data.  The present standard analysis shows that, in agreement with official recommendations, a conservative systematic error (at low energies) of $\\sim$ 2--3 $\\times$ 10$^{20}$ cm$^{-2}$ must be assumed when analyzing {\\it ASCA} SIS data. A soft-excess, with variable flux and/or shape, has been clearly detected as well as flux and spectral variability that confirm previous findings with other observatories. An anti-correlation is found between the spectral index and the flux in the 2-10 keV energy range. With the old response matrices, an iron emission line feature with EW $\\sim$ 50--100 eV was initially detected at $\\sim$ 5.6-5.7 keV ($\\sim$ 6.5-6.6 keV in the quasar frame) in 6 observations and, in two occasions, the line was resolved ($\\sigma \\sim$ 0.2-0.6 keV). Comparison with the Crab spectrum indicates however that this feature was mostly due to remaining calibration uncertainties between 5--6 keV. Indeed, fitting the data with the latest publicly available calibration matrices, we find that the line remains unambiguously significant in (only) the two observations with lowest fluxes where it is weak (EW $\\sim$ 20-30 eV), narrow and consistent with being produced by Fe K$_{\\alpha}$ emission from neutral matter.  Overall, the observations are qualitatively consistent with a variable, non-thermal X-ray continuum emission, i.e., a power law with $\\Gamma$ $\\sim$ 1.6 (possibly produced in the innermost regions of the radio-optical jet), plus underlying ``Seyfert-like'' features, i.e., a soft-excess and Fe K$_{\\alpha}$ line emission. The data are consistent with some contribution (up to a few 10\\% level in the {\\it ASCA} energy band) from a ``Seyfert-like'' direct continuum emission, i.e. a power law with $\\Gamma$ $\\sim$ 1.9 plus a reflection component, as well. When the continuum (jet) emission is in a low state, the spectral features produced by the Seyfert-like spectrum (soft-excess, iron line and possibly a steep power law plus a reflection continuum) are more easily seen. ",
        "introduction": "The remarkable discovery by EGRET on-board {\\it CGRO} that blazars (i.e. BL Lacertae objects and flat-spectrum radio quasars) are strong $\\gamma$-ray emitters has drawn in recent years the attention of the astronomical community to this class of objects. Observations indicate that the overall energy distribution of blazars shows the signature of two different types of emission mechanisms: beamed non-thermal jet radiation producing the overall broad band (from radio to $\\gamma$-rays) continuum emission common to all blazars, and quasi-isotropic thermal radiation by an accretion-disk producing a ``UV Bump'' observed in a large number of quasars and Seyfert galaxies but absent in BL Lac objects (e.g. Sambruna, Maraschi \\& Urry, 1996, Elvis et al. 1994). The non-thermal continuum consists of IR-optical and $\\gamma$-ray peaks. The first peak is interpreted in terms of synchrotron emission and the second peak in terms of inverse Compton emission (see Urry \\& Padovani 1995 for a review on the subject). From object to object, the observed different spectral characteristics may be due to the relative importance of one emission mechanism to the other which, in turn, is likely to be related to the amount of beaming in one object or the other (e.g. Dondi \\& Ghisellini 1995).  One of the most well-studied and characteristic example of blazars is the bright quasar 3C 273 (z$\\simeq$0.158). It is a good example where both non-thermal and thermal emission mechanisms might be observed because its broad band energy distribution exhibits two large peaks, one peaking in the IR-optical and one peaking in the $\\gamma$-rays with a `UV bump'' superimposed on it (Courvoisier et al. 1987, Lichti et al. 1995, von Montigny et al. 1997). The study of its X-ray properties may provide important clues to understand the origin of both emission mechanisms because a) soft-X-ray excess emission has been observed and interpreted as the high-energy tail of the UV bump (Turner et al. 1985, Courvoisier et al. 1987, Walter et al. 1994, Leach, Mc Hardy \\& Papadakis 1995) and b) the 2-10 keV spectrum which is most likely associated to the $\\gamma$-ray emission is known to be variable in time and shape (Turner et al. 1985), thus allowing for tests of different X- and $\\gamma$-ray emission models. \\pn To clarify the mechanism responsible for the X-ray emission, high quality data are first necessary to disentangle the contributions from the different spectral components, namely the jet and Seyfert components.  In this paper, we report on observations with {\\it ASCA} during the first year of the mission. The source spectral properties are shown with particular attention to calibration uncertainties, most relevant in this source because it was used for on-board calibration of the CCDs. We show evidence of complex spectral features (soft-excess, Fe K emission line, flux and spectral correlated variability) and discuss their possible interpretation. Throughout the analysis we use $H_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$ and $q_{0}$ = 0. ",
        "conclusions": "To date, {\\it ASCA} has observed 3C 273 10 times. Results from the first 9 observations, all performed during the first year of the mission have been presented here. These confirm and expand the evidence that the X-ray emission of 3C 273 is complex, with different spectral components contributing to its X-ray emission. Because 4 of the 9 observations were used for the on-board calibration of the CCDs, great care had to be taken when interpreting the observational results for this source. As a rule, {\\it absolute} values, in particular those obtained from the SIS, require a detailed estimate and knowledge of the instrumental systematic errors to be trusted. {\\it Relative} measurements (like flux and/or spectral variability) obtained from comparing different observations are, however, more reliable since they should not be affected by calibration uncertainties.  With this caveat in mind, it is found that: \\pn \\begin{enumerate} \\item A conservative systematic error at low energies that corresponds to an extra-absorption column of $\\sim$ 2--3 $\\times$ 10$^{20}$ cm$^{-2}$ is found for the SIS response, consistent with the ASCA Team's official prescriptions. \\item 2--10 keV flux variations by up to $\\sim$ 60\\% on a time-scale of $\\sim$ 200 days and day-to-day variations as large as $\\sim$ 20\\% were observed. \\item Extra soft X-ray emission is required by the data during the first observation, when the source was in its lowest flux level. \\item Flux and spectral slope variations are clearly detected as well and, for the first time, there is a statistically significant evidence that the index and flux are anti-correlated. \\item Iron line emission is detected in (only) the two observations with the lowest flux levels. The line is in both cases weak (EW$\\sim$20-30 eV), but statistically significant at more than 99\\% confidence level, narrow and consistent with Fe K$_{\\alpha}$ emission from neutral matter. \\end{enumerate}  We then speculate that all the above observable properties of the X-ray spectrum of 3C 273 can be interpreted in terms of the sum of two emission mechanisms. These are a non-thermal emission from the innermost regions of the jet which dominates the 2-10 keV region and whose signatures are the spectral variability and a flat ($\\Gamma \\sim 1.6$) power law continuum (that extrapolates well into higher energies), plus a diluted Seyfert-like spectrum whose signatures are the soft-excess and iron line emission. The newly discovered index-flux anti-correlation may be interpreted either by intrinsic variations of the jet power law index or by some contribution of the Seyfert-like continuum spectrum (say, a power law with $\\Gamma \\sim 1.9$) as the jet component varies. The overall scenario predicts that when the (dominant) jet component is in a low flux state, the spectral features produced by the Seyfert-like spectrum should be more easily detected (owing for variability of the Seyfert-like spectrum itself)."
    },
    "9803/astro-ph9803155_arXiv.txt": {
        "abstract": "We report on the mid-infrared imaging at 5, 7, 10 \\& 15 $\\mu$m of the galaxy cluster Abell 2218 obtained with the ISOCAM instrument onboard ESA's Infrared Space Observatory (ISO), as part of an on-going program to image gravitational arcs and arclets in distant clusters. Several cluster galaxies as well as field galaxies are detected. We discuss their mid-IR flux properties. ",
        "introduction": "Abell 2218 is a very rich galaxy cluster (richness class 4, according to Abell et al. 1989), characterized by a large velocity dispersion of the galaxy population ($\\sigma_v=1370$~km/s, Le Borgne et al. 1992), and a high X-ray temperature and luminosity ($T_x=7.2$~keV, see e.g. Markevitch 1997, $L_x = 2.9 \\times 10^{44}$ erg/s\\footnote{In this paper we adopt $H_0=75$~km/s/Mpc and $q_0=1/2$} in the energy band 0.5--4.4 keV, see e.g. Kneib et al. 1996).  These properties, coupled to a relatively small distance ($<z>=0.175$, Le Borgne et al. 1992), made this cluster an attractive target for studies of the Sunyaev-Zeldovich effect (Birkinshaw \\& Hughes 1994), and one of the closest clusters where gravitational arcs are detected. The amazing concentration of gravitational arcs and arclets has stimulated a huge observational effort first to get sub-arcsec imaging (from the ground, see, e.g. Kneib et al. 1995, and from space with HST, Kneib et al. 1996), and then to get spectra of the arcs (Ebbels et al. 1997).  The optical and near-IR observations have allowed a very detailed modelling of the mass distribution within the cluster (Kneib et al. 1996), confirmed later to a great level of accuracy, (Ebbels et al. 1997). It was found that the cluster mass distribution is bi-modal, with the main concentration centred on the cD galaxy. The X-ray surface brightness, as obtained via ROSAT observations, does not trace the gravitational potential as derived from the lensing analysis, a possible indication that the X-ray emitting gas is far from hydrostatic equilibrium (Kneib et al. 1996). This would also explain the discrepancy in the X-ray and lensing mass estimates (Markevitch 1997).  In this paper we report on the mid-IR observations of Abell 2218 done with ISOCAM on-board ESA's {\\em Infrared Space Observatory} (ISO) satellite. The high sensitivity of ISOCAM allows us to determine the photometric properties of lensed and cluster galaxies in the mid-IR band, thus widening the wavelength coverage of their spectra. Mid-IR observations are critical in understanding the intrinsic nature of these distant galaxies, since a starburst is known to emit a large fraction of its energy in the mid-IR, due to dust grain re-processing of the emitted radiation. ",
        "conclusions": "The LW1, LW2 and LW3 maps are presented in figures 1 \\& 2a, as overlays on top of the HST image (courtesy J.-P. Kneib) of Abell 2218.  \\begin{figure} \\plottwo{altierib1.eps}{altierib2.eps} \\caption{Abell 2218 4.5$\\mu$m (a) \\& 7$\\mu$m (b) contour maps} \\label{fig-1} \\end{figure}    In the cluster core, no arclet is detected at a significant level at any wavelength in the mid-IR. This non-detection of arcs or arclets contrasts with our recent mid-IR imaging of Abell 370 (Metcalfe et al. 1997), where the A0 giant arc is clearly detected as the main feature and emitter in the cluster core at 15$\\mu$m. But this giant arc is already a prominent feature in the optical, (first giant arc discovered), whereas arcs and arclets are fainter in Abell 2218. \\\\ The cD galaxy which is apparently centred on the cluster potential is clearly detected in the LW1 \\& LW2 filter bands. At 4.5$\\mu$m (rest-frame wavelength 3.8 $\\mu$m) the cD emission is extended, covering part of the optical halo; at 7$\\mu$m the emission is much more confined to the centre, but with an extension along the optical major axis, probably contaminated by the neighbouring merging dwarf galaxies. At 10$\\mu$m and 15$\\mu$m it lies just above the detection limit of ISOCAM and no statement can be made on its extension. The mid-IR  spectral energy distribution up to 15$\\mu$m seems to follow a simple Rayleigh-Jeans tail of the cold stellar component as found in optically-selected normal early-type galaxies (E, S0, S0a) in the Virgo cluster (Boselli et al. 1998). But due to the contamination by close-by galaxies it is very difficult to estimate the 4.5$\\mu$m flux of the cD. \\\\ The brightest cluster member galaxies are also detected at 4.5 $\\mu$m and 7 $\\mu$m. At the redshift of the cluster the 4.5 $\\mu$m emission is mostly coming again from cold stellar photospheres. The 7 $\\mu$m (LW2 band) (rest-frame 5.7 $\\mu$m) emission includes a small fraction of PAH emission but comes mainly again from the cold stellar photospheres. At 10 \\& 15 microns most of cluster members have vanished (ie. are below or very close to our detection limit), as expected for normal early-type galaxies (ellipticals) where there is only the Rayleigh-Jeans tail mid-IR emission from cold stellar photospheres, in dust \\& gas poor cluster core ellipticals.\\\\  Still one cluster member galaxy \\#373 (numbering system from Le Borgne et al. 1992 in the following) is detected at 15$\\mu$m, but also at 4.5 $\\mu$m and 7 $\\mu$m, it has a 'cart-wheel' like aspect, from the ring surrounding it and is probably a big face-on spiral galaxy. The 4.5$\\mu$m and 7$\\mu$m emission comes mostly from the nucleus whereas the 15$\\mu$m originates apparently from the 'ring' structure. It could be PAH and small hot grain emission enhanced by star formation in tidal shocks in the 'ring'.  \\begin{figure} \\plottwo{altierib3.eps}{altierib4.eps} \\caption{Abell 2218 15$\\mu$m contour maps (a) and spectral energy distributions of 4 objects} \\label{fig-2} \\end{figure}  Object \\#323 is detected at 15 $\\mu$m. This object was suspected to be an arclet candidate in the first deep multi-filter ground-based imaging (Pello et al. 1992), since it appeared as a fine unresolved extended structure, together with another 32 arclets in total. It is classified as one of the 235 arclet candidates in the first HST observation of Abell 2218 (Kneib et al. 1996). However, as noted by Kneib et al., it was also suspected not to be a strongly lensed object when comparison is made with shear orientation. This was again confirmed by Ebbels et al. (1997), object \\#323 being very elongated but at an angle of 45 degrees from the shear direction. Identification of several absorption features in its spectrum reveal it to be a cluster member at a redshift of z=0.179, a spiral seen edge-on.  To our surprise the brightest source by far from 7$\\mu$m to 15$\\mu$m is object \\#395, an apparently insignificant z=0.1032 (Le borgne 1992) foreground Sb galaxy, showing an ultraviolet excess and a strong H$\\alpha$ emission line, both indicating strong star formation of massive stars. Most of this mid-IR emission could then originate from the reprocessing by the small hot dust heated by a strong UV radiation field.  The second strongest emitter at long wavelength (from 10$\\mu$m to 15$\\mu$m), but much less at shorter wavelengths is object \\#317, a lensed galaxy at z=0.474 (Ebbels et al. 1997), a rather red object from the visible to the mid-IR. First indications show that this type of lensed object is common on our larger field imaging towards lensing clusters (Metcalfe et al. 1998 in preparation)  The spectral energy distributions from the visible (B,R,I) to the mid-IR of the 4 objects discussed above are shown in fig. 2b"
    },
    "9803/astro-ph9803188_arXiv.txt": {
        "abstract": "In a universe reionized in patches, the Doppler effect from Thomson scattering off free electrons generates secondary cosmic microwave background (CMB) anisotropies. For a simple model with small patches and late reionization, we analytically calculate the anisotropy power spectrum. Patchy reionization can, in principle, be the main source of anisotropies on arcminute scales. On larger angular scales, its contribution to the CMB power spectrum is a small fraction of the primary signal and is only barely detectable in the power spectrum with even an ideal, i.e.~cosmic variance limited, experiment and an extreme model of reionization. Consequently patchy reionization is unlikely to affect cosmological parameter estimation from the acoustic peaks in the CMB.  Its detection on small angles would help determine the ionization history of the universe in particular the typical size of the ionized region and the duration of the reionization process. ",
        "introduction": "It is widely believed that the cosmic microwave background (CMB) will become the premier laboratory for the study of the early universe and classical cosmology.  This belief relies on the high precision of the upcoming MAP\\footnote{{\\tt http://map.gsfc.nasa.gov}} and Planck Surveyor\\footnote{{\\tt http://astro.estec.eas.nl/SA-general/Projects/Planck}} satellite missions and the high accuracy of theoretical predictions of CMB anisotropies given a definite model for structure formation (\\cite{Hu95}\\ 1995).  To realize the potential of the CMB, aspects of structure formation affecting anisotropies at only the percent level in power must be taken into account.  A great uncertainty in models for structure formation is the extent and nature of reionization. Fortunately, this uncertainty is largely not reflected in the CMB anisotropies due to the low optical depth to Thomson scattering at the low redshifts in question. Reionization is known to be essentially complete by $z \\sim 5$ from the absence of the Gunn-Peterson effect in quasar absorption spectra (\\cite{Gun65}\\ 1965). Significant reionization before $z\\sim 50$ will be ruled out once the tentative detections of the CMB acoustic peaks at present are confirmed (\\cite{Sco95} 1995). It should be possible to deduce the reionization redshift $z_i$ through CMB polarization measurements (\\cite{Zal97}\\ 1997).  Nevertheless, the duration of time spent in a partially ionized state will remain uncertain. Moreover as emphasized by \\cite{Kai84} (1984) secondary anisotropies generated by the Doppler effect in linear perturbation theory are suppressed on small scales for geometric reasons (gravitational instability generates potential flows, leading to cancellations between positive and negative Doppler shifts).  Higher order effects which are not generally included in the theoretical modeling of CMB anisotropies are likely to be the main source of secondary anisotropies from reionization below the degree scale.  Such effects rely on modulating the Doppler effect with spatial variations in the optical depth.  Incarnations of this general mechanism include the Vishniac effect from linear density variations (\\cite{Vis87}\\ 1987), the kinetic Sunyaev-Zel'dovich effect from clusters (\\cite{SZ}\\ 1970), and the effect considered here: the spatial variation of the ionization fraction.  Reionization commences when the first baryonic objects form stars or quasars that convert a part of the nuclear or gravitational energy into UV photons.  Each such source then blows out an ionization sphere around it.  Before these regions overlap is a period when the universe is ionized in patches. The extent of this period and the time evolution of the size and number density of these patches depend on the nature of the ionizing engines in the first baryonic objects. Theories of reionization do not give robust constraints ({\\it cf.} \\cite{Teg94}\\ 1994; \\cite{Ree96}\\ 1996; \\cite{Agh96}\\ 1996; \\cite{Loe97}\\ 1997; \\cite{Hai97}\\ 1997; \\cite{Sil98}\\ 1998; \\cite{Hai98}\\ 1998).  We therefore take a phenomenological approach to studying the effects of patchy reionization on the CMB.  We introduce a simple but illustrative three parameter model for the reionization process based on the redshift of its onset $z_i$, the duration before completion $\\delta z$, and the typical comoving size of the patches $R$. It is then straightforward to calculate the CMB anisotropies generated by the patchiness of the ionization degree of the intergalactic medium.  We find that only the most extreme models of reionization can produce degree scale anisotropies that are observable in the power spectrum given the cosmic variance limitations. A large signal on degree scales requires early ionization, $z_i\\gtrsim 30$, long duration, $\\delta z\\sim z_i$, and ionization in very large patches, $R\\gtrsim 30$Mpc. Thus the patchiness of reionization is unlikely to affect cosmological parameter estimation from the acoustic peaks in the CMB (\\cite{Jun96}\\ 1996; \\cite{Zal97}\\ 1997; \\cite{Bon97}\\ 1997).  On the other hand, the patchy reionization signal on the sub-arcminute scale can, in principle, surpass both the primary and the secondary Vishniac signals.  These may be detectable by the Planck Surveyor and upcoming radio interferometry measurements (\\cite{Par97}\\ 1997) if point sources can be removed at the $\\Delta T/T \\sim 10^{-6}$ level.  An explicit expression for the CMB anisotropies power spectrum generated in a universe reionized in patches is given \\S \\ref{sec:explicit}. Simple order of magnitude estimate of the anisotropy from patches is given in \\S \\ref{sec:order}. In \\S \\ref{sec:power} we give a rigorous definition of our three-parameter reionization model, and calculate the patchy part of the power spectrum. We discuss illustrative examples in \\S \\ref{sec:discussion}. ",
        "conclusions": "The signal from patchy reionization in our model depends on four quantities: the rms peculiar velocity $\\left< v^2 \\right>^{1/2}$ today, the redshift of reionization $z_i$, its duration $\\delta z$ and the characteristic comoving size of the patches $R$. The structure formation model specifies the power spectrum of fluctuations which in turn tells us the rms peculiar velocity. Let us now consider the patchy reionized signal in the context of a specific model for structure formation.  For illustrative purposes, let us consider a cold dark matter model with $h=0.5$, $\\Omega_b  =0.1$, and a scale-invariant $n=1$ spectrum of initial fluctuations.  Normalizing the spectrum to the COBE detection via the fitting formulae of \\cite{Bun97} (1997) (their equations [17]-[20]) and employing the analytic fit to the transfer function of \\cite{Eis98} (1998) (their equations [15]-[24]) we find an rms velocity of $\\left<v^2\\right>^{1/2}=3.9\\times 10^{-3}$. With the present optical depth of $\\tau_0 = 0.122 \\Omega_b h = 0.0061$, we have a maximal anisotropy of \\begin{equation} ({l^2C_l\\over 2\\pi })_{\\rm max}=2.41\\times 10^{-15}{R\\over {\\rm Mpc}}\\delta z(1+z_i)^{3/2}, \\end{equation} at \\begin{equation} l_{\\rm max}={16958\\over R/{\\rm Mpc}} [ 1 - (1+z_i)^{-1/2}]. \\end{equation}  The power spectrum of the model in principle also tells us the remaining parameters of the ionization: its redshift $z_i$, duration $\\delta z$ and typical patch size $R$. Unfortunately, these quantities depend on details of the cooling and fragmenting of the first baryonic objects to form the ionizing engines.  We therefore consider $5 \\lesssim z_i \\lesssim 50$ which spans the range of estimates in the literature (\\cite{Teg94}\\ 1994; \\cite{Ree96}\\ 1996; \\cite{Agh96}\\ 1996; \\cite{Loe97}\\ 1997; \\cite{Hai97}\\ 1997; \\cite{Sil98}\\ 1998; \\cite{Hai98}\\ 1998).  Reionization, once it commences, is generally completed in a time short compared with the expansion time at that epoch $\\delta z /(1 +z_i) < 1$ by the coalescence of patches that are small compared with the horizon at the time $R/\\eta_i \\ll 1$ at the time.  Again the exact relations depend on the efficiency with which the first objects form and create ionizing radiation (see e.g. \\cite{Teg94}\\ 1994).  \\begin{figure}[htb] \\psfig{figure=patchf1.eps,width=3.3in} \\caption{CMB anisotropy power spectra in a CDM model with extreme patchiness. Shown here are the primary anisotropy and the patchy reionization anisotropy, eq.~(\\protect\\ref{eqn:Cl}) with $z_i=10$, $\\delta z=3$, $R=20$Mpc. These signals are compared with the cosmic variance of the primary anisotropy and the noise of the MAP satellite (in logarithmic bins).} \\label{fig:patch10} \\end{figure}   Let us consider an extreme example of $z_i=10$, $\\delta z=3$, $R=20{\\rm Mpc}$. Then the maximal power is $\\approx 5.3\\times 10^{-12}$ at $l\\approx 590$, the primary signal at these scales is $\\approx 3\\times 10^{-10}$ -- the contribution of the patchy reionization is small in comparison (see Fig.~\\ref{fig:patch10}).  However, in light of the high precision measurements expected from the MAP and Planck satellites such a signal is not necessarily negligible. The ultimate limit of detectability through power spectrum measurements is provided by so-called cosmic variance. This arises since we can only measure $2\\ell+1$ realizations of any given multipole such that power spectrum estimates will vary by \\begin{equation} \\delta C_\\ell = \\sqrt{2 \\over 2\\ell+1} C_\\ell^{(\\rm primary)}. \\end{equation} Detection of a broad feature such as that from patchy reionization is assisted in that we may reduce the cosmic variance by averaging over many $\\ell$'s. We show an example of this averaging in Fig.~\\ref{fig:patch10} (lower left boxes). In this model, the patchy reionization signal can be detected at the several $\\sigma$ level if cosmic variance were the main source of uncertainty. Of course a realistic experiment also has noise and systematic errors. We also show the noise error contributions expected from the MAP experiment in Fig.~\\ref{fig:patch10}.  An important additional source of uncertainty is provided by other unknown aspects of the model.  Indeed it is hoped that the CMB power spectrum can be used to measure fundamental cosmological parameters to high precision. Excess variance from patchy reionization can in principle cause problems for cosmological parameter estimation from the CMB if not included in the model. It would remain undetected and produce parameter misestimates if its signal can be accurately mimicked by variations in the other parameters. Fortunately, the angular signature we find here -- $\\ell^2$ white noise until some cut off due to the patch size -- does not resemble the signature of other cosmological parameters which alter the positions and amplitudes of the acoustic peaks (see \\cite{Bon97}\\ 1997; \\cite{Zal97}\\ 1997).  Coupled with the small amplitude of the effect on the 10 arcminute to degree scale for even this extreme model, it is unlikely that patchy reionization will significantly affect parameter estimation through the CMB.  We have called the ($z_i=10,\\delta z=3,R=20$) model extreme, because of the size of patches; the reionization redshift and duration would be considered reasonable by a number of theories. For example the early quasar model of Haiman \\& Loeb (1998) does predict $z_i\\sim 10$ and $\\delta z \\sim 3$. However, their ``medium quasar'' emits only $\\sim 10^{67}$ ionizing photons during its life time. These photons cannot ionize a bubble larger than $R\\sim 1$Mpc comoving.  Perhaps more interesting is the case where reionization takes place at a higher redshift with for example $z_i=30$, $\\delta z=5$, $R=3{\\rm Mpc}$.  The reduction in the patch size causes the signature to move to smaller angles where the primary signal is negligible due to dissipational effects at recombination. The increase in the optical depth at this higher redshift is counterbalanced by the reduction in the rms fluctuation due to the number of patches along the line of site such that the amplitude of the signal increases only moderately.  Here the maximal power is $\\approx 6.2\\times10^{-12}$ at $l\\approx 4650$ (see Fig.~\\ref{fig:patch30}).  Patchy reionization effects exceed the Vishniac signal at these scales ($\\approx 3\\times 10^{-12}$) which is believed to be the leading other source of secondary anisotropies (\\cite{Hu96}\\ 1996).  \\begin{figure}[htb] \\psfig{figure=patchf2.eps,width=3.3in} \\caption{CMB anisotropy power spectra in a CDM model with early reionization. Shown here are the primary anisotropy suppressed by rescattering and the patchy reionization anisotropy, eq.~(\\protect\\ref{eqn:Cl}) with $z_i=30$, $\\delta z=5$, $R=3$Mpc. These signals are compared with the cosmic variance of the primary anisotropy achievable by an ideal experiment in the absence of galactic and extragalactic foregrounds.} \\label{fig:patch30} \\end{figure}  Although the morphology and amplitude of the patchy reionization and Vishniac signals are similar, the Vishniac effect is fully specified by the ionization redshift and the spectrum of initial fluctuations and hence may be removed once these are determined from parameter estimation at larger angular scales.  Likewise, since the rms peculiar velocity $\\left< v \\right>$ and the ionization redshift $z_i$ will be specified by the large scale observations, the amplitude of the signal can be used to estimate the duration of reionization $\\delta z$ and its angular location the typical comoving size of the bubbles $R$.  In summary, the patchiness of reionization leaves a potentially observable imprint on the CMB power spectrum, but one that is unlikely to affect cosmological parameter estimation from the acoustic peaks in the CMB.   We show how the signature scales with the gross properties of reionization -- its redshift, duration, and typical patch size.   Observational detection of this signature would provide useful constraints on the presently highly uncertain reionization scenarios but will likely require experiments with angular resolution of an arcminute or better and foreground subtraction at better than the $\\delta T/T \\sim 10^{-6}$ level."
    },
    "9803/astro-ph9803231_arXiv.txt": {
        "abstract": "We present radio continuum observations of the spiral-shaped ionized feature (Sgr A West) within the inner pc of the Galactic center at three epochs spanning 1986 to 1995.  The VLA A-configuration was used at $\\lambda$2cm (resolution of 0\\dasec1$\\times$0\\dasec2).  We detect proper motions of a number of features in the Northern and Eastern Arms of Sgr A West including the ionized gas associated with IRS 13 with V(RA)= 113$\\pm$10, V(Dec)=150$\\pm$15 km s$^{-1}$, IRS 2 with V(RA)= 122$\\pm$11, V(Dec)=24$\\pm$34 km s$^{-1}$ and the Norther Arm V(RA)= 126$\\pm$30, V(Dec)=--207$\\pm$58 km s$^{-1}$.  We also report the detection of features having transverse velocities $>$1000 \\kms\\ including a head-tail radio structure, the ``Bullet'', $\\approx$4$''$ northwest of Sgr A$^*$ with V(RA)= 722$\\pm$156, V(Dec)=832$\\pm$203 km s$^{-1}$, exceeding the escape velocity at the Galactic center.  The proper motion measurements when combined with previous H92$\\alpha$ radio recombination line data suggest an unambiguous direction of the flow of ionized gas orbiting the Galactic center.  The measured velocity distribution suggests that the ionized gas in the Northern Arm is not bound to the Galactic center assuming a 2.5 million solar mass of dark matter residing at the Galactic center.  This implies that the stellar and ionized gas systems are not dynamically coupled, thus, supporting a picture in which the gas features in the Northern Arm and its extensions are the result of an energetic phenomenon that has externally driven a cloud of gas cloud into the Galactic center. ",
        "introduction": "The ionized gas known as Sagittarius A West (Sgr A West) appears as a three-Arm spiral-like structure (North, East, and West Arms) engulfing the inner pc of the Galaxy where Sgr A$^*$, the compact radio source at or near the dynamical center of the Galaxy lies (Ekers et al. 1983). These features are surrounded by neutral gas in the circumnuclear disk (CND) rotating with the velocity of about 100 \\kms\\ at the distance of 2 pc from the Galactic center (e.g. G\\\"usten et al. 1988).  The kinematics of ionized gas surrounding Sgr A$^*$ show systematic velocities along various components of Sgr A West including Western Arc with a radial velocity structure which varies regularly between --100 and +100 \\kms\\ in the North-South direction (e.g. Serabyn et al. 1988; Herbst et al. 1993; Roberts \\& Goss 1993). However, the velocity structure of the inner 10$''$ where there is a hole in the distribution of ionized gas, known as the ``mini-cavity'', becomes increasing more negative $\\approx-$350 \\kms\\ approaching Sgr A$^*$ (Yusef-Zadeh, Morris \\& Ekers 1989; Roberts, Yusef-Zadeh \\& Goss 1996, hereafter RYG).  Recent observations of stellar proper motions shows evidence of a 2.5$\\times10^6$ \\msol\\ object lying close to the position of Sgr A$^*$ (Eckart and Genzel 1997).  The stars orbit randomly around the Galactic center with increasing velocity dispersion around Sgr A$^*$, reflecting the gravitational potential of central mass.  The ionized gas, on the other hand, is part of a coherent flow with systematic motion that is decoupled from the stellar orbits. Understanding the kinematics of the system of ionized gas is complicated by its incomplete view of the 3-dimensional geometry with respect to Sgr A$^*$ as well as by the interaction of orbiting gas with non-gravitational forces, such as the winds from the cluster of hot mass-losing stars near the Galactic center. To examine the gas kinematics and the true geometry of the ionized flow, in this {\\it Letter} we present the results of proper motion measurements of ionized gas at $\\lambda$2cm. ",
        "conclusions": "By combining the transverse and radial velocities, we are able to unambiguously determine the direction of ionized flow at the Galactic center.  The predominant component of the motion in the plane of the sky is from east to west for most of the measured features with the exception of few places where the velocity of ionized gas is anomalously large.  It appears that the flow of ionized gas in the Northern Arm (Box 12) originates in the northeast with red-shifted velocities in the orbital plane. The ionized gas then follows an orbital trajectory to the southwest as it crosses the plane of the sky and passes by  Sgr A$^*$ before the ionized gas moves to the northwest.   One idea that has been suggested to explain the origin of the mini-cavity and its unusual characteristics (e.g. Lutz et al. 1993; Melia et al. 1996) is the the collision of fast-moving Blobs with the orbiting ionized gas. The Blobs are hypothesized to be formed as a result of high velocity winds of IRS 16 cluster escaping from but focussed by the gravitational potential of Sgr A$^*$.  This focusing mechanism allows the diffuse outflowing materials to collide with each other and form dense Blobs of ionized gas leaving the gravitational potential of Sgr A$^*$ (Wardle \\& Yusef-Zadeh 1992).  The anomalous high-velocity features seen in Boxes 1, 4, 8 and 9 are consistent with the outflow picture. In particular, the Bullet is clearly escaping from the gravitational potential of the Galactic center region even when the mass of the stellar cluster is included. In this model, however, it is not expected to see a tail produced behind the fast moving Blobs.  The existence of a bow shock or X-ray emitting gas associated with the head of the Bullet would favor a model in which these fast-moving features are ejected by mass-losing stellar sources.  Further observations of the Bullet are needed to understand its origin.  The comparison between the measured total velocities and the upper limits to the escape velocities at projected distance (r) from the center of each box to Sgr A$^*$ (Table 1) suggests that the ionized flow is on an unbound orbit around the Galactic center. This is consistent with the model of RYG that the ionized gas in the Northern Arm is on a hyperbolic orbit.  Like the compact and relatively isolated features with high transverse velocities discussed above (e.g. the Bullet and the Blobs), the rest of the orbiting ionized features are extended and follow a global velocity field which may also be unbound to the Galactic center. If we use the projected distance as cos$^{-1}$ (45$^\\circ$) of the actual distance, almost all the measured velocities are greater than the escape velocities listed in Table 1. The mass within the inner 10-20$''$ is assumed to be dominated by a compact source centered on Sgr A$^*$ having a mass of 2.5$\\times10^6$ \\msol\\ as measured recently from stellar proper motion measurements (Eckart \\& Genzel 1997).  The escape velocity estimates may not be applicable at large distances where the mass of the evolved stellar cluster becomes important. From the comparison of the three dimensional stellar and ionized gas motions, it appears that these two systems are not dynamically coupled in the inner 20$''$ of the Galactic center.  The effect of a strong gravitational potential due to the large concentration of dark matter near Sgr A$^*$ is manifested as high velocity gradients of over 600 \\kms\\ pc$^{-1}$. However, the existence of ionized gas in an unbound orbit is inconsistent with the notion that the ionized gas in the Northern Arm is a tidally stretched infalling feature (e.g. Serabyn et al. 1988).  Additionally, the present proper motion data do not support the interpretation that the Northern Arm is a segment of a one-armed spiral pattern induced as a result of an instability in the rotating disk (e.g. Lacy et al. 1991). A significant variation in the velocity distribution of ionized gas along Northern Arm is not consistent with a small variations expected from Keplerian motion, thus supporting that the Northern Arm is a material feature.  We believe that the high velocity of ionized gas on an unbound orbit supports a scenario in which an energetic phenomenon outside the inner few parsecs of the Galactic center accelerated a cloud to pass by the Galactic center, which then collides with the CND and results in the loss of angular momentum of the material in the CND (Serabyn et al. 1988).  There is evidence of disturbed neutral gas and shocked molecular gas based on OH 1720 MHz maser emission at the interface of the extension of the Northern and Eastern Arms of Sgr A West and a ``gap'' in the CND (Yusef-Zadeh et al. 1996; Jackson et al. 1993).  In this picture, the Northern and Eastern Arms delineate the edges of the intruding cloud photoionized by the UV radiation field at the Galactic center (Jackson et al. 1993).  The neutral gas in the ``gap'' of the CND is interpreted to be the site of collision with a cloudlet pushed into the Galactic center, possibly by the energetic explosion of Sgr A East.  Future modeling of the three-dimensional motion of ionized gas should constrain the inclination of the orbital plane of the ionized gas with respect to the orbital plane of the CND and examine the possibility that the CND is origin of the Northern and Eastern Arms."
    },
    "9803/astro-ph9803007_arXiv.txt": {
        "abstract": "Since their discovery in 1973, Gamma-Ray Bursts (GRBs) have remained for many years one of the most elusive mysteries in High Energy-Astrophysics. The main problem regarding the nature of GRBs has usually been the lack of knowledge of their distance scale. About 300 GRBs are detected annua\\-lly by BATSE in the full sky, but only a few of them can be loca\\-lized accurately to less than half a degree. For many years, follow-up observations by other satellites and ground-based telescopes were conducted, but no counterparts at other wavelengths were found. The breakthrough took place in 1997, thanks to the observation by {\\it BeppoSAX} and {\\it RossiXTE} of the fa\\-ding X-ray emission that follows the more energetic gamma-ray photons once the GRB event has ended. This emission (the afterglow) extends at longer wavelengths, and the good accuracy in the position determination by {\\it BeppoSAX} has led to the discovery of the first optical counterparts -for GRB 970228, GRB 970508, and GRB 971214-, greatly improving our understanding of these puzzling sources. Now it is widely accepted that most bursts originate at cosmological distances but the final solution of the GRB problem is still far away. \\vspace {5pt} \\\\    Key~words: multiwavelength observations; gamma-ray bursts. ",
        "introduction": "In 1967-73, the four {\\it VELA} spacecraft (named after the spanish verb {\\it velar}, to keep watch), that where originally designed for verifying whether the former Soviet Union abided by the Limited Nuclear Test Ban Treaty of 1963, observed 16 peculiarly strong events (Klebesadel, Olson and Strong 1973, Bonnell and Klebesadel 1996). On the basis of arrival time differences, it was determined that they were related neither to the Earth nor to the Sun, but they were of cosmic origin. Therefore they were named cosmic $\\gamma$-ray Bursts (GRBs hereafter).  GRBs appear as brief flashes of cosmic high energy photons, emitting the bulk of their energy above $\\approx$ 0.1 MeV (Fig. 1). They are detected by instruments somewhat similar to those used by the particle physicists at their laboratories. The difference is that GRB detectors have to be placed onboard balloons, rockets or satellites. In spite of the abundance of new observations of GRBs, their energy source and emission mechanism remain highly speculative.  The KONUS experiment on {\\it Veneras 11} and {\\it 12} gave the first indication that GRB sources were isotro\\-pically distributed in the sky (Mazets et al. 1981, Atteia et al. 1987).  Based on a much larger sample, this result was nicely confirmed by BATSE on board the {\\it CGRO} (Meegan et al. 1992). About 300 GRBs occur annually in the full sky, but only few of them are localized accurately.  The apparent isotropy was interpreted in terms of GRBs produced at cosmological distances, although the possibility of a small fraction of the sources lying nearby, within a galactic disc scale of few hundred pc, or in the halo of the Galaxy, could not be discarded. Another result was that the time profiles of the bursts are very diffe\\-rent, with some GRBs lasting a few ms and others lasting for several minutes.  In general, there was no evidence of periodicity in the time histories of GRBs.  However there was indication of a bimodal distribu\\-tion of burst durations, with $\\sim$25\\% of bursts ha\\-ving durations around 0.2 s and $\\sim$75\\% with durations around 30 s. A extensive review of the observational character\\-istics can be found in Fishman and Meegan (1995).  A deficiency of weak events was also noticed, and all these observational data led many researchers to believe that GRBs are indeed at cosmological distances.  In this case, the released energies could be as high as 10$^{53}$ erg and models could involve coalescence of neutron stars in double systems, neutron star-black hole systems, accretion induced collapse in white dwarfs or $^{\\prime\\prime}$failed$^{\\prime\\prime}$ Type I supernovae. It was also proposed the possibility of having a mixture of two popu\\-lations: a cosmological plus a galactic one, but the latter, probably formed by accreting neutron stars, will account for only a very small fraction of the total population. See Nemiroff (1984) for a review of the di\\-fferent theoretical models.  It is well known that an important clue for resol\\-ving the GRB puzzle is the detection of transient emission -at longer wavelengths- associated with the bursts. Here I review all the efforts in the search for GRB counterparts throughout the electromagnetic spectrum. I will first review the searches prior to 1997, and afterwards I will discuss the important discove\\-ries achieved last year. Previous reviews can be seen in Schaefer (1994), Hartmann (1995), Vrba (1996), Greiner (1996a) and Hurley (1998).  \\begin{figure}[htp] \\epsfxsize=70mm \\epsfysize=60mm \\centerline{\\epsfbox{fig1.eps}} \\caption{One of the GRBs detected by {\\it BATSE}, lasting for about 150-s. Further details in Connaughton et al. (1997).} \\end{figure} ",
        "conclusions": "The existence of an X-ray afterglow in {\\it all} bursts seems to be confirmed. The first optical/IR counterparts have been found in 1997. GRB 970508 was also detected in radio and mm, but prompt searches at other wavelengths failed to detect GRB 970111, GRB 970402 and GRB 970828.  {\\it BeppoSAX} and {\\it RossiXTE} have opened a new window in the GRB field and it is widely accepted now that most GRBs, if not all, lie at cosmological distances. It is expected that {\\it BeppoSAX}, {\\it RossiXTE} and {\\it CGRO} will facilitate the discoveries of other counterparts and, together with the new high-energy observatories ({\\it AXAF}, {\\it SPECTRUM X/$\\Gamma$}, {\\it XMM}, {\\it INTEGRAL}, {\\it HETE 2}) and other satellites of the future {\\it 4th Interplanetary Network}, will definitively solve the long-standing Gamma-Ray Burst mystery."
    },
    "9803/astro-ph9803282_arXiv.txt": {
        "abstract": "Microlensing is increasingly gaining recognition as a powerful method for the detection and characterization of extra-solar planetary systems.  Naively, one might expect that the probability of detecting the influence of more than one planet on any single microlensing light curve would be small.  Recently, however, Griest \\& Safizadeh (1998) have shown that, for a subset of events, those with minimum impact parameter $u_{min} \\lsim 0.1$ (high magnification events), the detection probability is nearly 100\\% for Jovian mass planets with  projected separations in the range 0.6--1.6 of the primary Einstein ring radius $R_E$, and remains substantial outside this zone.  In this Letter, we point out that this result implies that, regardless of orientation, {\\it all} Jovian mass planets with separations near 0.6--1.6$R_E$ dramatically affect the central region of the magnification pattern, and thus have a significant probability of being detected (or ruled out) in high magnification events. The probability, averaged over all orbital phases and inclination angles, of two planets having projected separations within $0.6$--$1.6R_E$ is substantial: 1-15\\% for two planets with the intrinsic orbital separations of Jupiter and Saturn orbiting around 0.3--1.0$M_\\odot$ parent stars.  We illustrate by example the complicated magnification patterns and light curves that can result when two planets are present, and discuss possible implications of our result on detection efficiencies and the ability to discriminate between multiple and single planets in high magnification events. ",
        "introduction": "A planetary microlensing event occurs whenever the presence of a planet creates a perturbation to the standard microlensing event light curve. These perturbations typically have magnitudes of $\\lsim 20\\%$ and durations of a few days or less. First suggested by Mao \\& Paczy\\'nski (1991) as a method to detect extra-solar planetary systems, the possibility was explored further by Gould \\& Loeb (1992), who found that roughly 15\\% of microlensing light curves should show evidence of planetary deviations if all primary lenses have Jupiter-mass planets with orbital separations comparable to that of Jupiter. Although these probabilities are relatively high, the use of microlensing to discover planets was largely ignored since in order to detect the primary events the microlensing survey teams must monitor millions of stars in very crowded fields, resulting in temporal sampling that is too low ($\\sim 1\\, {\\rm day}$) and photometric errors that are too high ($\\gsim 5\\%$) to detect most secondary planetary deviations (Alcock et al.\\ 1997a).  Recently, the situation has changed dramatically as the real-time reduction of the survey teams has enabled them to issue electronic ``alerts,'' notification of on-going events detected before the peak magnification (Udalski et al.\\ 1994, Pratt et al.\\ 1996), allowing other collaborations to perform special purpose observations of the alerted events.  These additional observations include denser photometric sampling by the PLANET and GMAN collaborations (Albrow et al.\\ 1996, 1997, 1998 and Pratt et al.\\ 1996, Alcock et al. 1997b) as well as spectroscopic follow-up of particular events (Lennon et al.\\ 1997). Over 60 events are currently alerted per year towards the Galactic Bulge. Since only a handful of these are on-going at any given time, monitoring teams can sample events very densely and with high photometric accuracy, enabling the detection of many second order effects, including --in principle-- planetary anomalies.  No clear planetary detections have yet been made in this way, but preliminary estimates of detection efficiencies show that PLANET, over the next two observing seasons, should be sensitive to planetary anomalies caused by Jovian planets orbiting a few AU from their parent star (Albrow et al.\\ 1998). Thus, if these kinds of planets are common, they should be detected soon. If not, microlensing will be able to place interesting upper limits on the frequency of such systems.  These observational developments have been accompanied by an explosion of theoretical work, including further studies of detection probabilities and observing strategies incorporating a variety of new effects (Bolatto \\& Falco 1994, Bennett \\& Rhie 1996, Peale 1997, Sackett 1997, DiStefano \\& Scalzo 1998a,b), demonstration of planetary microlensing light curves (Wambsganss 1997), explorations of the degeneracies in the fits of planetary events (Gaudi \\& Gould 1997, Gaudi 1998), and a study of the relation between binary and planetary lenses (Dominik 1998). It would thus seem that the theoretical understanding of the detection and characterization of planetary systems using microlensing should be well in hand.  Surprisingly, however, the field still has surprises to offer. Recently, Griest \\& Safizadeh (1998, hereafter GS) came to a rather startling",
        "conclusions": "In this {\\it Letter\\/}, we have demonstrated that: (1) the probability of two planets having projected separations that fall in the ``standard lensing zone'' ($0.6 < b < 1.6$) is quite high, $\\sim 1-15\\%$ for planets with true separations corresponding to Jupiter and Saturn orbiting stars of typical mass; (2) the influence of multiple planets in and somewhat beyond the standard lensing zone can be profound for high magnification events ($u_{min}<0.1$,) however (3) for some geometries, the magnification pattern and resulting light curves from multiple planets are qualitatively degenerate with those from single-planet lensing, and (4) for high magnification events, finite source effects are likely to suppress more substantially the amplitude of multiple planet deviations than single planet deviations.  Given these results, it would appear that the effects of multiple planets on the detection and characterization of planetary systems warrant future study. All previous theoretical studies have calculated microlensing planet detection sensitivities either by ignoring multiple planets or by treating each planet independently. For high impact parameter events (low magnification), this is probably a fair assumption, but as the magnification maps in Fig.~2 illustrate detection probabilities will need to be revised for small impact parameters (large magnification). The sense of revision will likely depend on finite source effects. It is also likely that for some geometries serious degeneracies exist between light curves arising from multiple and single planet high magnification events; these degeneracies are above and beyond those present in the single planet case discussed by Griest \\& Safizadeh (1998).  This possible degeneracy is especially pertinent in light of the fact that the conditional probability of having two planets in the lensing zone is substantial. Thus, the interpretation of any given high magnification event may be difficult: the degeneracies should be characterized and their severity determined in order to have a clear understanding of the kinds of systems whose parameters can be unambiguously determined from the deviations. Finally, the calculation of planet detection efficiencies for high magnification events should consider multiple planets in order to be able to reliably convert the observed frequency of planetary deviations into a true frequency of planetary systems."
    },
    "9803/astro-ph9803018_arXiv.txt": {
        "abstract": "In observations with the Rossi X-ray Timing Explorer we have discovered quasi-periodic oscillations (QPOs) near 1 kHz from 4U~1705-44, a low--mass X-ray binary with a neutron star classified as an atoll source. In six separate observations, we detect one QPO with a frequency ranging between 770 and 870 Hz and a 4\\% rms fraction in the full detector energy band.  There is evidence for a second QPO at 1073 Hz in one interval. The separation in frequency of the two QPOs is $298\\pm11$ Hz.  The QPOs are present only in observations where the mass accretion rate is inferred to be at an intermediate level, based on the atoll source phenomenology.  At the highest accretion rates, the QPOs are not detected with upper limits to the rms fraction of about 2\\%. At the lowest accretion rates the upper limits are about 4\\%.  The QPO frequency increases with inferred mass accretion rate. This is expected in models where the QPO frequency is generated by motion at an inner edge of the accretion disk. An increased mass accretion rate causes the disk edge to move in, increasing the orbital frequency.  Five Type--I X-ray bursts are observed with no detectable oscillations. ",
        "introduction": "Quasi-periodic oscillations (QPOs) with frequencies near 1 kilohertz were discovered from X-ray binaries with the Rossi X-Ray Timing Explorer (RXTE) almost immediately after its launch (for reviews and references see van der Klis 1997, Kaaret \\& Ford 1997a, Ford 1997). These QPOs are likely produced very near the accreting neutron star where the dynamical time scale is near $10^{-3}$ seconds. Observations of such fast oscillations provides a new opportunity to measure fundamental properties of neutron stars and the effects of strong field gravity in low mass X-ray binaries.  Kilohertz QPOs have been discovered in both major subclasses of LMXBs: the Z-sources and the atoll sources (see Hasinger \\& van der Klis 1989). Certain trends are already apparent.  For example, the kilohertz QPOs in Z-sources are much weaker than those in the lower luminosity atoll sources. In the sources with luminosities intermediate between the majority of Z and atoll sources, the persistently bright galactic bulge sources, QPOs are apparently absent or very weak (Wijnands, van der Klis, \\& van Paradijs 1997a). To understand the physical mechanisms at work, further observations of a large sample of sources are needed with the unique capabilities of RXTE, likely to be unequaled in the next decade.  Here we report the discovery of fast QPOs from the low mass X-ray binary 4U~1705-44 using RXTE.  4U~1705-44 has been classified as an atoll source (Hasinger \\& van der Klis 1989). Its different source states are supposed to reflect changes in mass accretion rate. We observe a link between the source state and the presence and frequencies of fast QPOs.  In Section~2 we summarize the observations and analysis techniques. Section~3 presents the detections of fast QPOs and X-ray burst properties.  In Section~4 we identify the atoll source states during our observations. Section~5 is a discussion of the relation of the source states to the QPO properties and physical implications. ",
        "conclusions": "The 760--870 Hz QPO in 4U~1705-44 is probably a single feature and is likely the lower frequency of two QPO peaks. The second QPO, which we identify at 1074 Hz, then is the higher frequency peak. More observations in the lower banana state are needed to confirm the presence of this second QPO. Such phenomenology is very similar to other atoll sources. A single strong QPO near 800 Hz is present at times in several sources, e.g. 4U~1608-52 (Berger et al. 1996; Mendez et al. 1998), 4U~1820-30 (Smale, Zhang \\& White 1997), 4U~1636-53 (Zhang et al. 1996; Wijnands et al. 1997b). In all of these sources a second QPO is sometimes detectable (see above references). The frequency separation of the two peaks is similar to that in 4U~1705-44.  Such double QPOs, and also the QPOs observed in X-ray bursts (e.g. Strohmayer, Zhang \\& Swank 1997), suggest a beat frequency mechanism (Alpar \\& Shaham 1985; Miller, Lamb \\& Psaltis 1998). If such an interpretations holds, then the frequency separation of the QPOs in 4U~1705-44 implies that the spin period of the neutron star in this system is $3.35\\pm0.12$ msec.  The appearance and frequency of the fast QPOs in 4U~1705-44 are correlated with the states of the source. The states in turn are thought to be related to the mass accretion rate, with the accretion rate increasing from the island to the banana.  At the lower end of the banana branch a QPO appears at about 780 Hz.  Further along the branch the frequency increases to 835 Hz and then 865 Hz.  This behavior supports models where the QPO modulation is generated at an inner disk edge which shrinks as the mass accretion rate increases (e.g. Miller, Lamb \\& Psaltis 1998).  We detect the second QPO at 1073 Hz only in the the lowest part of the banana where the frequency of the stronger QPO is at its lowest.  Farthest along the banana branch the QPO becomes undetectable with rms fraction upper limits of about 2\\%.  At the highest accretion rates, the reduced QPO amplitude is perhaps due to spreading of the accretion stream, alternatively a puffed--up inner disk may obscure the central source (Smale, Zhang \\& White 1997).  Similar behavior is observed in other X-ray binaries.  Atoll source state identifications have been made simultaneous with fast QPO detection in 4U~1636-53 (Zhang et al. 1996; Wijnands et al.  1997b), 4U~1735-44 (Wijnands et al. 1997c), 4U~1820-30 (Smale, Zhang \\& White 1997) and KS~1731-260 (Wijnands \\& van der Klis 1997d).  These sources were observed in the banana branch. Consistent with the present data, the QPO in all cases is detected only in the lower part of the banana and disappears as the source moves up along the banana branch. In 4U~0614+091 a similar effect appears as the QPO amplitude decreases as the count rate increases (Ford 1997).  The situation is less clear at low mass accretion rates, i.e. in the island states.  Our upper limits of about 4\\% rms fraction are not very constraining.  QPOs have been detected in island states in both 4U~0614+091 (Mendez et al. 1997) and 4U~1608-52 (Yu et al. 1997). At the lowest count rates in the island state of 4U~0614+091 the QPO is not detected (Mendez et al. 1997). This fits the general trend we observe here.  The QPO frequency shows no correlation with count rate or energy flux in the 2--10 keV band, though we note that the range of count rates is small (1205 to 1236 c/s). An effect similar to that seen in 4U~0614+091 may be at work, where there is no unique correlation between rate and frequency in different observations separated by several months (Ford et al. 1997a). The present data indicate that the QPO frequency is better correlated with the atoll state of the source. Changes in the source state are reflected as changes in the x-ray colors, which in turn are simply changes in the energy spectrum of the source. Therefore the state--frequency correlation also manifests as a spectrum--frequency correlation.  From other sources we know that the changes in QPO frequency are linked to changes in the energy spectrum.  The frequencies of the QPOs in 4U~0614+091 are well correlated with the flux of a blackbody component in the energy spectrum (Ford et al. 1997b), and in that source and also 4U~1608-52 the QPO frequencies are correlated with the spectral index of the power law component in the energy spectrum (Kaaret et al. 1998).  We thank Mariano Mendez for assistance in implementing the frequency shifting technique.  We thank Rudy Wijnands, Rob Fender and the referee for helpful discussions and comments. ECF acknowledges support by the Netherlands Foundation for Research in Astronomy with financial aid from the Netherlands Organization for Scientific Research (NWO) under contract numbers 782-376-011 and 781-76-017."
    },
    "9803/astro-ph9803312_arXiv.txt": {
        "abstract": "In this talk, I have discussed some issues of recent interest and activity in the field of neutrino astrophysics and cosmology. The topics are: (1)~The origin of high peculiar velocities of pulsars; (2)~Energization of the supernova shock wave; (3)~Ultra-high energy neutrino astronomy; (4)~Possible implications of the recent measurements of low deuterium abundance. ",
        "introduction": "It was known, since the birth of modern astrophysics in the early part of the 20th century, that neutrinos play an important role in various processes that occur within a stellar core and which are responsible for energy generation in a star. Gradually, the importance of neutrinos were understood in stars outside the main sequence. And, since the discovery of the microwave blackbody radiation, it was taken for granted that there is a similar cosmic background of neutrinos, although experimentally this background has not been detected so far. Various constraints from neutrino properties have been deduced from this belief, some of which are much better than the corresponding constraints from earth-based experiments. For example, one can cite the mass bound on stable neutrinos which are derived from the energy density of the universe as a whole. This sets an upper bound of order of a few tens of eV, whereas the direct measurement of the mass of $\\nu_\\tau$ sets upper bounds in the range of a few tens of MeVs. If the neutrinos are unstable, then also there exists quite severe bounds on their lifetimes.  Unfortunately, in this talk I cannot review all of these aspects. Rather, I will have to assume that the audience is familiar with these concepts. The reason is that, fortunately, there has been a lot of progress in the field of neutrino astrophysics in the last year and a half, and quite a few of them are remarkable. I have to concentrate on these recent developments. I cannot guarantee that I will cover even all of the interesting recent developments. Let me say I will cover what I know, with the restriction that I will leave out topics such as solar neutrinos and atmospheric neutrinos, which are covered by other speakers in this conference. ",
        "conclusions": ""
    },
    "9803/astro-ph9803124_arXiv.txt": {
        "abstract": "Echelle spectra have been obtained of the \\caII\\ H and K lines for a sample of metal-poor subdwarf stars as well as for a number of nearby Population~I dwarfs selected from among those included in the Mount Wilson HK survey. The main conclusion of this paper is that \\caII\\ H- and K-line emission does occur among subdwarfs. It is particularly notable among those subdwarfs with colours of $B-V \\geq 0.75$; all such stars observed exhibit chromospheric emission, although emission is observed among some subdwarfs bluer than this colour. The \\caII\\ K emission profile in most subdwarfs exhibits an asymmetry of $V/R >1$, similar to that seen in the integrated light of the solar disk. Two quantitative indicators of the contrast between the peaks in the emission profile and the neighbouring photospheric line profile are introduced. Measurements of these indicators show that the level of \\caII\\ emission among the subdwarfs is similar to that among low-activity Population~I dwarfs. ",
        "introduction": "Chromospheric activity among main-sequence stars is well known to be a decreasing function of age. The fluxes in the chromospheric emission features formed in the cores of the \\caII\\ H and K lines show a decrease with age among stars younger than 4.5 Gyr (Skumanich 1972).  Within most main-sequence stars chromospheric heating is thought to be governed by the activity of an interior magnetic dynamo, the strength of which depends on the stellar rotation rate. As main-sequence stars age they spin down, with a consequent decrease in dynamo activity and chromospheric heating (Hartmann \\& Noyes 1987).  The dynamo model has met with considerable success in describing the evolution of chromospheric activity among main-sequence stars. However, much of the data published to date on chromospheric activity among dwarf stars of spectral types F-G-K pertains to stars younger than $\\sim$ 5 Gyr. Such is true, for example, of the studies by Skumanich (1972) and Simon, Herbig, \\& Boesgaard (1985). By comparison, there is much less spectroscopy available pertaining to chromospheric activity among older dwarfs, particularly halo subdwarfs. Consequently, little is known about the evolution of chromospheric activity among the oldest main-sequence stars in the Galaxy. To study the low levels of activity expected among such stars, we have carried out a spectroscopic study of the \\caII\\ H and K lines for a sample of low-metallicity subdwarfs. Metal-poor subdwarfs which are members of the Galactic halo are amongst the oldest stars in the Galaxy. They are useful for studying chromospheric activity at great ages.  The observations of \\caII\\ H and K line emission among subdwarfs reported in this paper complement the HST GHRS observations of \\mgII\\ lines among metal-poor solar-like stars by Peterson \\& Schrijver (1997). ",
        "conclusions": "The main conclusion of this paper is that Population II subdwarfs do have chromospheres, i.e., their outer atmospheres are heated by some form of non-radiative process. In terms of the contrast of the peaks in the K-line emission profile relative to the neigbouring absorption profile, and the preponderance of an asymmetry of $V/R > 1$, the subdwarfs appear similar to low-activity dwarfs sampled from the Mount Wilson survey. These observations are consistent with data on the \\mgII\\ lines of metal-poor solar-type stars reported by Peterson \\& Schrijver (1997).  The origin of an asymmetry of $V/R > 1$ among both dwarfs and subdwarfs is not well understood. As Linsky (1980) has pointed out such an asymmetry can be produced by downward motions in the region of formation of the central K$_3$ absorption feature, or by upwards motions in the K$_2$ formation region if there are no systematic motions in the region of K$_3$ formation (see Ayres \\& Linsky 1975 for the definitions of this terminology). Presumably the reason for the $V/R >1$ asymmetry among the subdwarfs is the same as for the Population~I dwarfs, including the Sun. However, the origin of this asymmetry even for the best studied case of the Sun is still not completely clear. This subject has been reviewed by Linsky (1980), Cram (1983), and Rutten \\& Uitenbroek (1991). The concensus seems to be that the asymmetry is associated with upwards propagating gravity or acoustic waves (see, e.g., Carlsson \\& Stein 1992, 1997).  In addition to the uncertainty in the origin of the $V/R$ asymmetry, the \\caII\\ K-line emission observations do not permit a determination of whether the chromospheric activity of the subdwarfs is being modulated by a magnetic dynamo, as for the Sun, or whether it is produced by a process which is active even in the absence of such a dynamo. Statistical analyses of the chromospheric emission line fluxes of late-type Population~I stars have been used to argue that the emission can be divided into two components (Schrijver 1987a,b, 1995): a ``magnetic dynamo'' component which correlates with stellar rotation rate and decreases with age, and a rotation-independent ``basal'' component. The magnetic dynamo component is important for stars like the Sun, and is particularly strong in late-type dwarfs younger than 1 Gyr. The basal component is thought to reflect the contribution of acoustic heating to the maintenance of a chromosphere (Schrijver 1987b; Cuntz, Rammacher, \\& Ulmschneider 1994).  Whether the chromospheres of subdwarfs are ``basal chromospheres'' or are maintained by a magnetic dynamo is not clear from the \\caII\\ K line observations. The finding by Baliunas et~al. (1995) that the subdwarf HD~103095 exhibits an activity cycle of period 7.3 years in the \\caII\\ H and K emission lines, suggests an analogy with the solar cycle, implying that the subdwarfs have chromospheres whose long-term behaviour is dominanted by a magnetic dynamo.  The similarity in \\caII\\ K-line emission strength among the subdwarfs and the lowest-activity stars sampled from the Mount Wilson survey seems consistent with the spectroscopy of main-sequence stars in the old open clusters NGC 188 and M67 by Barry, Cromwell, \\& Hegge (1984) which suggests that the rate of decline of the fluxes in the Ca II H and K line cores flattens out after $\\sim$ 3 Gyr. These observations, and those noted in the previous paragraph, would be consistent with a scenario in which the magnetic dynamo ceases to decline significantly in activity once main-sequence stars reach an age of around 5 Gyr. There could be two reasons for this. Either the dynamo reaches a true minimum level, as in the turbulent magnetic field scenario of Durney, De Young, \\& Roxburgh (1993), or else the rate of spin down of dwarfs and subdwarfs older than 5 Gyr becomes very slow.  The alternative to these dynamo scenario is that chromospheres among stars older than 5 Gyr are largely maintained by an age-invariant basal heating mechanism, and that activity produced by an age-dependent magnetic dynamo has dropped to a low level by comparison. In sufficiently old main-sequence subdwarf stars the dynamo might be relatively inactive and acoustic heating may become the dominant mechanism for maintaining a ``basal chromosphere.'' This is the scenario proposed by Peterson \\& Schrijver (1997) on the basis of similarities observed between the \\mgII\\ emission line profiles of metal-poor solar-type stars and solar quiet regions. In this respect it is perhaps relevant to note that the most extreme $V/R >1$ \\caII\\ K line asymmetry in the integrated spectrum of the solar disk occurs around the times of solar minimum (White \\& Livingston 1981).  The type of observation that may differentiate between these scenarios is a search for soft x-ray emission among subdwarfs. If subdwarfs are found to have solar-like coronae, with soft x-ray properties consistent with temperatures of 1-2$\\times 10^6$ K, then the case for dynamo-maintained chromospheres and coronae would be strengthened.  We thank the referee for a number of useful comments and for drawing our attention to a number of relevant references in the literature.  \\onecolumn"
    },
    "9803/astro-ph9803254_arXiv.txt": {
        "abstract": "We have imaged a sample of 45 face-on spiral galaxies in the K-band, to determine the morphology of the old stellar population, which dominates the mass in the disk. The K-band images of the spiral galaxies have been used to calculate different characteristics of the underlying density perturbation such as arm strengths, profiles and cross-sections, and spiral pitch angles.  Contrary to expectations, no correlation was found between arm pitch angle and Hubble type, and combined with previous results this leads us to conclude that the morphology of the old stellar population bears little resemblance to the optical morphology used to classify galaxies.  The arm properties of our galaxies seem inconsistent with predictions from the simplest density wave theories, and some observations, such as variations in pitch angle within galaxies, seem hard to reconcile even with more complex modal theories. Bars have no detectable effect on arm strengths for the present sample.  We have also obtained B-band images of three of the galaxies. For these galaxies we have measured arm cross-sections and strengths, to investigate the effects of disk density perturbations on star formation in spiral disks. We find that B-band arms lead K-band arms and are narrower than K-band arms, apparently supporting predictions made by the large scale shock scenario, although the effects of dust on B-band images may contribute towards these results. ",
        "introduction": "Ever since the pioneering work of Lin \\& Shu (1964, 1966) many theoretical models have been proposed to explain the existence of spiral structure in disk galaxies. Many are based upon the Lin-Shu Hypothesis (e.g. Lin, Yuan \\& Shu 1969; Roberts 1969; Roberts, Roberts \\& Shu 1975) which describes spiral patterns as density waves, which cause compression of the gas component as it flows through the arms leading to subsequent star formation. Modal theories (Bertin et al. 1989a, 1989b; Bertin \\& Lin 1996) are more complex density wave models, which describe disks of galaxies as resonant cavities within which density waves of different modes can co-exist and interfere to produce a range of observed phenomena. Tidal models (Toomre \\& Toomre 1972; Kormendy \\& Norman 1979) set up spiral structure through transient density waves, which are caused by the tidal field of a nearby neighbour. Bars can also potentially drive spiral structure (Sanders \\& Huntley 1976). In this case the formation of arms is driven by the effect of the bar potential on the dissipative interstellar medium. Finally Stochastic Self-propagating Star Formation (SSPSF) can describe flocculent structure. Here, short irregular arms are produced due to shearing of a part of the disk that has recently formed stars (Gerola \\& Seiden 1978). SSPSF alone cannot describe Grand-Design structure, but some models (e.g. Sleath \\& Alexander 1995, 1996) use SSPSF with a weak imposed density wave to describe such global spiral modes.  Many of the recent advances in this area have arisen from a study of the atomic and molecular gas component in spiral galaxies, through studies of HI and CO line emission.  This has permitted the detailed mapping of both the distribution and velocity field of the gas in the spiral arms of nearby galaxies, e.g. M81 (Visser 1980) M51 (Tilanus \\& Allen 1991, 1993; Rand 1993; Nakai et al. 1994), M100 (Knapen 1993; Rand 1995), NGC~3627 (Reuter et al. 1996) and NGC~6946 (Regan \\& Vogel 1995).  These studies find streaming velocities of gas through the spiral arms of order a few 10s of km~s$^{-1}$, and also find offsets between the peaks of the gas density and the old stellar population in the arms and the star formation as revealed by H$\\alpha$ emission.  All of these findings are in general agreement with the predictions of density wave theories, with the gas being compressed and shocked as it flows into the spiral arm, and subsequently forming stars downstream from the density peak of the spiral arm. However, these studies are only possible in the strongest arms of the nearest spiral galaxies, and hence yield little information on the global importance of density waves in the general population of disk galaxies. Here we undertake the complementary approach of calculating the strength of density waves as mapped out by the old stellar population of a large sample of galaxies, using near-IR K-band ($2.2\\mu m$) imaging. In this waveband the old stellar population is observed (Rix \\& Rieke 1993), making this the best observational tracer of the stellar mass distribution. Also, K--band measurements are less affected by extinction than measurements in the optical.  In a previous paper (Seigar \\& James 1998 - hereafter Paper I) we described the properties of the bulges, disks and bars of a sample of 45 galaxies. In this paper we describe how spiral arm structures can be extracted for these galaxies and determine quantitative parameters for them. In section 2 we describe the observations; section 3 describes the data reduction and extraction of all of the spiral arm parameters used in the analysis; section 4 then draws on these parameters in discussing, in turn, each of the theoretical models of spiral structure; in section 5 we describe tests of models of star formation in spiral galaxy disks and in section 6 we summarise our findings. ",
        "conclusions": "We now summarise the implications of these observations for the main models of the formation and stability of spiral structure. No one theory fits all of the observations, but a combination can explain most or all of the features we have observed. This section summarises these conclusions.  Simple density wave theories (e.g. Lin \\& Shu 1964, 1966) are undermined in three ways. Firstly, we do not get the correlation between the fraction of light in the disk and pitch angle which is predicted by Lin \\& Shu (1964) and Roberts et al. (1975). Secondly, Lin \\& Shu (1964) predict that pitch angle should be the same for arms at the same radius within a given galaxy. We find that this is not the case. Finally, we have found that the FWHM of arm cross-sections are somewhat narrower than predicted by simple density wave theories.  More complicated density wave theories, e.g. Modal theory (Bertin et al. 1989a, 1989b; Bertin \\& Lin 1996), are supported by modulation effects that we have found in arm strength as a function of radius, and which appear to be fairly common.  Bar driven models of spiral structure (Sellwood \\& Sparke 1988; Sanders \\& Huntley 1976) are not supported by our observations, as we find no correlation between arm strength and bar strength, even in the inner part of the disk.  Tidal effects from near neighbours seem to have some effect. We find that 4 out of 7 galaxies with arm EA greater than $35^{\\circ}$ have near neighbours whereas only 12 out of 45 galaxies in whole sample have near neighbours. We also find that in the Fourier analysis of galaxies with near neighbours, the even low-order modes are dominant, especially the $m=2$ mode.  We have also discussed factors affecting star formation in the disks of spiral galaxies. We have found evidence both for and against the Large Scale Shock Scenario (Roberts 1969). Firstly, we find a correlation between arm contrast and the normalised star formation rate, in agreement with the predictions of the large scale shock scenario. Secondly, we find that B-band arms are more loosely wound and narrower than K-band arms, confirming the predictions made by density wave theories (Roberts 1969). However, this can also be interpreted as an extinction effect because dust lies on the trailing edges of arms.  Finally, and possibly most significantly, we find no correlation between the arm properties of this sample of galaxies and their classified Hubble type.  Combining this with the poor correlation between Hubble type and K-band B/D ratio (Paper I), it would appear impossible to allocate a Hubble type to these galaxies on the basis of these K-band observations, and the morphology of the old stellar population varies surprisingly little between Hubble types Sa and Sd. We speculate that the determining parameter for Hubble type is cold gas content, which controls the star formation rate and the distribution of the young stars, which dominate the optical appearance."
    },
    "9803/astro-ph9803062_arXiv.txt": {
        "abstract": "We report the initial results of a survey for intracluster planetary nebulae in the Virgo Cluster.  In two $16' \\times  16'$ fields, we identify 69 and 16 intracluster planetary nebula candidates, respectively.  In a third $16' \\times  16'$ field near the central elliptical galaxy M87, we detect 75 planetary nebula candidates, of which a substantial fraction are intracluster in nature.  By examining the number of the planetaries detected in each field and the shape of the planetary nebula luminosity function, we show that 1) the intracluster starlight of Virgo is distributed non-uniformly, and varies between subclumps A and B{}, 2) the Virgo Cluster core extends $\\sim 3$~Mpc in front of M87, and thus is elongated along the line-of-sight, and 3) a minimum of 22\\% of Virgo's stellar luminosity resides between the galaxies in our fields, and that the true number may be considerably larger.  We also use our planetary nebula data to argue that the intracluster stars in Virgo are likely derived from a population that is of moderate age and metallicity. ",
        "introduction": "The concept of intracluster starlight was first proposed by Zwicky (1951), when he claimed to detect excess light between the galaxies of the Coma cluster.  Follow-up photographic searches for intracluster luminosity in Coma and other rich clusters (Welch \\& Sastry 1971; Melnick, White, \\& Hoessel 1977; Thuan \\& Kormendy 1977) produced mixed results, and it was not until the advent of CCDs that more precise estimates of the amount of intracluster starlight were made (cf.~Guldheus 1989; Uson, Boughn, \\& Kuhn 1991; V\\'ichez-G\\'omez, Pell\\'o, \\& Sanahuja 1994; Bernstein \\etal 1995). All these studies suffer from a fundamental limitation: the extremely low surface brightness of the phenomenon.  Typically, the surface brightness of intracluster light is less than 1\\% that of the sky, and measurements of this luminosity must contend with the problems presented by scattered light from bright objects and the contribution of discrete sources below the detection limit.  Consequently, obtaining detailed information on the distribution, metallicity, and kinematics of intracluster stars through these types of measurements is extremely difficult, if not impossible.  An alternative method for probing intracluster starlight is through the direct detection and measurement of the stars themselves.  Recent observations have shown this to be possible.  In their radial velocity survey of 19 planetary nebulae (PN) in the halo of the Virgo Cluster galaxy NGC~4406 (M86), Arnaboldi \\etal (1996) found 3 objects with $v > 1300$~\\kms; these planetaries are undoubtably intracluster in origin. Similarly, Ferguson, Tanvir, \\& von Hippel (1998) detected Virgo's intracluster component via a statistical excess of red star counts in a Hubble Space Telescope (HST) Virgo field over that in the Hubble Deep Field. Finally, intracluster stars have been unambiguously identified from the ground via planetary nebula surveys in Fornax (Theuns \\& Warren 1997) and Virgo (M\\'endez \\etal 1997; Ciardullo \\etal 1998).  Motivated by these results, we have begun a large scale [O~III] $\\lambda 5007$ survey of intergalactic fields in Virgo, with the goal of mapping out the distribution and luminosity function of intracluster planetary nebulae (IPN){}.  Depending on the efficiency of tidal stripping, Virgo's intracluster component is predicted to contain anywhere from 10\\% to 70\\% of the cluster's total stellar mass (Richstone \\& Malumuth 1983; Miller 1983). A survey of several square degrees of Virgo's intergalactic space with a four meter class telescope should therefore detect several thousand PN, and shed light on both the physics of tidal-stripping and on the initial conditions of cluster formation.  Here, we present the first results of our survey. ",
        "conclusions": "We report the results of a search in three fields in the Virgo Cluster for intracluster planetary nebulae, and have detected a total of 95 intracluster candidates.  From analysis of the numbers of the planetaries, we find that the amount of intracluster light in Virgo is large (at least 22\\% of the cluster's total luminosity), distributed non-uniformly, and varies between subclump~A and B{}.  By using the planetary nebulae luminosity function, we derive an upper limit of $\\sim 12$~Mpc for the distance to the front edge of the Virgo Cluster and use this to show that the cluster must be elongated along our line of sight.  We also use the properties of planetary nebulae to suggest that the intracluster stars of Virgo have moderate age and metallicity.  The large fraction of intracluster stars found has potentially serious consequences for models of cluster formation and evolution, and for cosmological models.  Finally, we note that this survey included less than 0.2\\% of the traditional 6 degree core of the Virgo Cluster.  Many more intracluster planetary nebulae wait to be discovered.  We thank Allen Shafter for some additional off-band observations and Ed Carder at NOAO, for his measurements of our on-band filter so that we could begin our observations on time.  We would also like to thank the referee, R. Corradi, for several suggestions that improved the quality of this paper.  Figure~1 was extracted from the Digitized Sky Survey, which was produced at the Space Telescope Science Institute under U.S. Government grant NAGW-2166.  This work was supported in part by NASA grant NAGW-3159 and NSF grant AST95-29270.  \\pagebreak"
    },
    "9803/gr-qc9803031_arXiv.txt": {
        "abstract": "Binary systems comprising at least one neutron star contain strong gravitational field regions and thereby provide a testing ground for strong-field gravity. Two types of data can be used to test the law of gravity in compact binaries: binary pulsar observations, or forthcoming gravitational-wave observations of inspiralling binaries. We compare the probing power of these two types of observations within a generic two-parameter family of tensor-scalar gravitational theories. Our analysis generalizes previous work (by us) on binary-pulsar tests by using a sample of realistic equations of state for nuclear matter (instead of a polytrope), and goes beyond a previous study (by C.M.~Will) of gravitational-wave tests by considering more general tensor-scalar theories than the one-parameter Jordan-Fierz-Brans-Dicke one. Finite-size effects in tensor-scalar gravity are also discussed. ",
        "introduction": "The detection of gravitational waves by kilometric-size laser-interferometer systems such as LIGO in the US and VIRGO in Europe will initiate a new era in astronomy. One of the most promising sources of gravitational waves is the inspiralling compact binary, a binary system made of neutron stars or black holes whose orbit decays under gravitational radiation reaction. The observation of these systems will provide important astrophysical information, e.g. masses of neutron stars, and direct distance measurements up to hundreds of Mpc \\cite{reviews}. It is also said that detecting gravitational waves from inspiralling binaries should provide rich tests of the law of relativistic gravity in situations comprising strong-field regions (like near a neutron star or a black hole). However, present binary-pulsar experiments already provide us with deep and accurate tests of strong-field gravity \\cite{taylor94,twdw92,dt92}. It is therefore interesting to compare and contrast the probing power of present (and foreseeable) pulsar tests with that of future gravity-wave observations.  A convenient quantitative way of doing this comparison is to work within a multi-parameter family of physically motivated (and physically consistent) theories of gravity which differ from Einstein's theory in their radiative and strong-field predictions. The most natural framework of this type is the general class of tensor-scalar theories in which gravity is mediated both by a tensor field $g_{\\mu \\nu}^*$ (``Einstein metric'') and by a scalar field $\\varphi$. These theories contain one arbitrary ``coupling function'' $A(\\varphi)$ defining the conformal factor relating the pure spin-2 Einstein metric $g_{\\mu \\nu}^*$ to the ``physical metric'' $\\widetilde{g}_{\\mu \\nu} = A^2 (\\varphi) \\, g_{\\mu \\nu}^*$ measured by laboratory clocks and rods. The usual Jordan-Fierz-Brans-Dicke theory \\cite{Jordan,Fierz,BransDicke} is the one parameter class of theories defined by a coupling function $A(\\varphi) = \\exp (\\alpha_0 \\, \\varphi)$. [The coupling strength $\\alpha_0$ is related to the often used parameter $\\omega$ by $\\alpha_0^2 = (2\\omega + 3)^{-1}$.] Will \\cite{will94} has studied the quantitative constraints on the coupling parameter $\\alpha_0$ of Jordan-Fierz-Brans-Dicke theories that could be brought by gravitational-wave observations. His result is that in most cases the bounds coming from gravity-wave observations will be comparable to presently known bounds coming from solar-system experiments (namely, $\\alpha_0^2 < 10^{-3}$). This result of Ref.~\\cite{will94} seems to suggest that gravitational-wave-based {\\it strong-field} tests of gravity do not go really beyond the solar-system {\\it weak-field} tests of gravity. We wish, however, to emphasize that this seemingly pessimistic conclusion is mainly due to having restricted one's attention to the special, one-parameter Jordan-Fierz-Brans-Dicke theory. Indeed, in this theory the strength of the coupling of the scalar field $\\varphi$ to matter is given by the constant quantity $\\alpha_0$ independently of the state of condensation of the gravitational source. As a consequence, the predictions of the theory differ from those of Einstein's theory by a fraction of order $\\alpha_0^2$ in all situations: weak-field ones or strong-field ones, alike. By contrast, it has been emphasized in Refs.~\\cite{def93,def96} that the more generic tensor-scalar theories in which the strength of the coupling of $\\varphi$ to matter, namely \\begin{equation} \\alpha (\\varphi) \\equiv \\frac{\\partial \\ln A (\\varphi)}{\\partial \\, \\varphi} \\, , \\label{eq1.1} \\end{equation} depended on the value of $\\varphi$, allowed for the existence of genuine {\\it strong-field effects}, by which the presence of a highly condensed source, such as a neutron star, could generate order-unity deviations from general relativity even if the weak-field limit of $\\alpha (\\varphi)$ is arbitrarily small. [More precisely, these non-perturbative strong-field effects take place when $\\beta (\\varphi) \\equiv \\partial \\, \\alpha (\\varphi) / \\partial \\, \\varphi$ is negative.] This led us to consider, instead of the Jordan-Fierz-Brans-Dicke model $A(\\varphi) = \\exp(\\alpha_0 \\varphi)$, the class of theories defined by \\begin{equation} A(\\varphi) = \\exp \\left( \\frac{1}{2} \\, \\beta_0 \\varphi^2 \\right) \\, . \\label{eq1.2} \\end{equation}  This class of theories contains two arbitrary (dimensionless) parameters: $\\beta_0$ appearing in Eq.~(\\ref{eq1.2}), and $\\varphi_0$, the asymptotic value of $\\varphi$ at spatial infinity. [By contrast, in Jordan-Fierz-Brans-Dicke theory, $\\varphi_0$ has no observable effects.] Equivalently, the two parameters in these theories can be defined as being the strength of the linear coupling of $\\varphi$ to matter in the weak-field limit $(\\varphi \\approx \\varphi_0)$, \\begin{equation} \\alpha_0 = \\alpha (\\varphi_0) = \\frac{\\partial \\ln A(\\varphi_0)}{\\partial \\, \\varphi_0} = \\beta_0 \\varphi_0 \\, , \\label{eq1.3} \\end{equation} and the non-linear coupling parameter \\begin{equation} \\beta_0 = \\frac{\\partial \\, \\alpha (\\varphi_0)}{\\partial \\, \\varphi_0} = \\frac{\\partial^2 \\ln A(\\varphi_0)}{\\partial \\, \\varphi_0^2} \\, . \\label{eq1.4} \\end{equation}  A convenient feature of the two-parameter family of theories (\\ref{eq1.2}) is that they have just the amount of generality needed both to parametrize the most general boost-invariant weak-field (``post-Newtonian'') deviations from Einstein's theory, and to encompass nonperturbative strong-field effects. Indeed, on the one hand, the two theory-parameters $(\\alpha_0 , \\beta_0)$ determine the two well-known Eddington-Nordtvedt-Will parameters, \\begin{equation} \\overline{\\gamma} \\equiv \\gamma_{\\rm Eddington} - 1 = -2 \\, \\alpha_0^2 / (1 + \\alpha_0^2) \\, , \\label{eq1.5} \\end{equation} \\begin{equation} \\overline{\\beta} \\equiv \\beta_{\\rm Eddington} - 1 = \\frac{1}{2} \\, \\beta_0 \\, \\alpha_0^2 / (1 + \\alpha_0^2)^2 \\, , \\label{eq1.6} \\end{equation} which measure the most general, boost-invariant, deviations from general relativity at the first post-Newtonian level. (See \\cite{def92,def2pn} for the generalization of the Eddington parameters to the second post-Newtonian level.) On the other hand, when $\\beta_0 \\lesssim -4$ non-perturbative strong-field effects develop in the theories defined by Eq.~(\\ref{eq1.2}).  All existing gravitational experiments can be interpreted as constraints on the two-dimensional space of theories defined by Eq.~(\\ref{eq1.2}). In other words, we can work within the common $(\\alpha_0 , \\beta_0)$ plane, and consider each gravitational experiment (be it of weak-field or strong-field nature) as defining a certain exclusion plot within that plane. In some recent work \\cite{def96}, we have constructed such exclusion plots, as defined by considering both solar-system experiments and binary-pulsar experiments. The present work will generalize these exclusion plots in several respects: (i)~we shall plot the regions of the $(\\alpha_0 , \\beta_0)$ plane probed by future gravitational-wave observations of neutron star--neutron star, and neutron star--black hole systems, (ii)~we shall improve our previous study of the probing power of binary-pulsar experiments by considering a sample of realistic nuclear equations of states (instead of the polytropic one used by us before) and by using updated pulsar data, and (iii)~we shall consider the individual constraints on $\\alpha_0$ and $\\beta_0$ brought by the main solar-system experiments (instead of using published combined limits on $\\overline{\\gamma}$ and $\\overline{\\beta}$).  This paper is organized as follows: In Sec.~II we summarize our (numerical) approach to computing the various form factors that describe the coupling of the scalar field $\\varphi$ to a neutron star. In Sec.~III we generalize Ref.~\\cite{will94} in discussing how gravitational wave observations can give quantitative tests of tensor-scalar gravity. In Sec.~IV we combine and compare gravitational-wave tests with binary-pulsar tests and solar-system ones. Our conclusions are presented in Sec.~V, while an Appendix discusses finite-size effects in tensor-scalar gravity. ",
        "conclusions": "The main conclusion of the comparison carried out in Figs.~\\ref{fig1}--\\ref{fig5} is that, in all cases, future LIGO/VIRGO observations of inspiralling compact binaries turn out not to be competitive with present binary-pulsar tests in their {\\it discriminating\\/} probing of the strong-field, and radiative, aspects of relativistic gravity. This conclusion may seem paradoxical. It should not be interpreted negatively against LIGO/VIRGO observations which, as shown in Figs.~\\ref{fig1}--\\ref{fig5}, will independently probe strong-field gravity and will exclude regions of parameter space allowed by solar-system experiments. Rather, it is simply a reminder that binary-pulsar experiments are superb tools for probing strong-field and radiative aspects of gravity. It is also somewhat a good news for gravitational wave data analysis (which promises to be already a very challenging task even if one {\\it a priori\\/} assumes the validity of general relativity; see, e.g., \\cite{dis98}). Indeed, our results Figs.~\\ref{fig1}--\\ref{fig5} indicate that our present experimental knowledge of the law of relativistic gravity is sufficient to justify using general relativity as the standard theory of gravitational radiation.  Note that this conclusion explicitly refers to the quantitative probing of plausible\\footnote{Within the presently accepted framework for low-energy fundamental physics, namely field theory, the only alternative (long-range) gravity theories which, (i) do not violate the basic tenets of field theory, and (ii) are not already necessarily extremely constrained by existing equivalence-principle tests, are the tensor-scalar gravity theories.} {\\it deviations\\/} {}from general relativity. At the {\\it qualitative\\/} level, and also at the {\\it non-discriminating\\/} quantitative level, LIGO/VIRGO observations will bring invaluable advances in our experimental knowledge of relativistic gravity. First, they will provide the first direct observation of gravitational waves far in the wave zone (while binary pulsar experiments prove the reality of the propagation with finite velocity of the gravitational interaction in the near zone of a binary system). Second, they will (hopefully) lead to superb additional confirmations of general relativity through the observation of the wave forms emitted during the inspiral and coalescence of neutron stars or black holes; see, e.g., \\cite{reviews,bs9495}.  Independently of the comparison between LIGO/VIRGO tests and binary-pulsar tests, the present paper has provided the first systematic study of the influence of the nuclear equation of state on the theoretical probing power of binary-pulsar tests. In particular, Fig.~\\ref{fig2} shows that if the equation of state were as soft as predicted by the simple Pandharipande model, binary-pulsar tests quantitatively supersede solar-system ones in all the region $\\beta_0 \\lesssim -1$ of parameter space. Even if we consider the less constraining stiff equations of states, the present work confirms the limit \\begin{equation} \\beta_0 > -4.5 \\label{eq5.1} \\end{equation} found (modulo $10\\%$) in \\cite{def96}. We recall that this limit can be interpreted as a limit on the ratio of the two weak-field post-Einstein parameters \\begin{equation} \\frac{\\beta_{\\rm Edd} - 1}{\\gamma_{\\rm Edd} - 1} \\approx - \\frac{1}{4} \\, \\beta_0 < 1.1 \\, . \\label{eq5.2} \\end{equation}  Finally, we pointed out that the discovery of a binary pulsar with a black-hole companion has the potential of providing a superb new probe of relativistic gravity. The discriminating power of this probe might supersede all its present and foreseeable competitors in measuring $\\alpha_0$ down to the level $\\alpha_0^2 < 10^{-6}$."
    },
    "9803/astro-ph9803289_arXiv.txt": {
        "abstract": "We propose a model for the source of the X-ray background (XRB) in which low luminosity active nuclei ($L\\sim 10^{43}\\ergps$) are obscured ($N\\sim 10^{23}\\, {\\rm cm}^{-2}$) by nuclear starbursts within the inner $\\sim 100$~pc. The obscuring material covers most of the sky as seen from the central source, rather than being distributed in a toroidal structure, and hardens the averaged X-ray spectrum by photoelectric absorption. The gas is turbulent with velocity dispersion $\\sim {\\rm few}\\times 100\\, \\kmps$ and cloud-cloud collisions lead to copious star formation.  Although supernovae tend to produce outflows, most of the gas is trapped in the gravity field of the starforming cluster itself and the central black hole.  A hot ($T\\sim 10^6-10^7$~K) virialised phase of this gas, comprising a few per cent of the total obscuring material, feeds the central engine of $\\sim 10^7\\, \\Msun$ through Bondi accretion, at a sub-Eddington rate appropriate for the luminosity of these objects. If starburst-obscured objects give rise to the residual XRB, then only 10 per cent of the accretion in active galaxies occurs close to the Eddington limit in unabsorbed objects. ",
        "introduction": "The flat spectrum of the X-ray background (XRB) above 1~keV (Marshall et al 1980; Gendreau et al 1995; Chen, Fabian \\& Gendreau 1997) is not simply accounted for by the integration of known classes of source, which generally have much steeper spectra. A currently popular model invokes intrinsic absorption in active galaxies with which to flatten the observed X-ray spectrum (Setti \\& Woltjer 1989; Madau, Ghisellini \\& Fabian 1994; Celotti et al 1995; Comastri et al 1995). The intrinsic absorption must range from column densities of about $10^{22} - 10^{24}\\psqcm$, and must cover most (at least $\\sim 2/3$)\\footnote{Since the XRB has a flat spectrum right down to 1~keV, the fraction of sources with absorption exceeding $10^{21}\\, \\pcmsq$ must approach 90 per cent; see Section 3} of the sky as seen by the source itself. The combination of different levels of absorption and redshift of the objects can then lead to the observed power-law background spectrum in the 1--10~keV band. The rollover in the spectrum at about 30~keV is due redshift acting on an intrinsic 100~keV break, such as is seen in nearby active galaxies (Zdziarski et al 1995).  The geometry of the obscuring material within a typical, X-ray background-contributing, active galaxy is unclear. The nucleus must be mostly surrounded by a typical column density of say $10^{23}\\psqcm$. If the material is freely orbiting the nucleus, then it should soon collide with itself, dissipate, and flatten into a disk, unless either a) the orbits are carefully arranged or b) there is sufficient energy to continually throw matter into a wide range of orbital inclinations. We note that what evidence there is for a torus of molecular material around active galaxies does not indicate that its high column density part covers much of the Sky. Indeed, models where the obscuring material is extended on scales $\\sim 100\\, \\pc$ (Granato et al 1997) account for the infrared properties of Seyfert galaxies more successfully than ones which postulate a thick compact ($<1\\, \\pc$) torus (Pier \\& Krolik, 1992, 1993).  The first case a) may be accounted for by a warped disk. Pringle (1996) and Maloney, Begelman \\& Pringle (1996) have shown that the outer parts of irradiated accretion disks are unstable to warping, the final result of which is that much of the sky is covered by the warp. The details are currently uncertain as to whether most of the sky can be covered and whether such high column densities can be attained in the warped material. We do not pursue that further here. The second case b) may be accounted for by a nuclear starburst, the supernovae of which can provide the energy to push clouds around and so obscure the nucleus. It is that possibility we explore further here.  There is indeed much empirical evidence for a connection between nuclear starbursts and active galactic nuclei (Terlevich \\& Melnick 1985, Terlevich, D\\'{\\i}az \\& Terlevich 1990; Perry \\& Williams 1993, Heckman et al 1995, 1997; Maiolino et al 1997). The extra featureless ultraviolet continuum in Seyfert 2 galaxies (FC2) appears to be due to a nuclear starburst (Terlevich \\& Cid-Fernandes 1995; Heckman et al 1997). Turner et al (1997, 1998) in a discussion of the absorption properties of Seyfert 2 galaxies propose that the starburst may be involved there. The striking similarity between the evolution of the luminosity density due to star formation (Madau et al 1996) and that due to QSOs (Boyle \\& Terlevich 1998, Dunlop 1998) also suggests a close link between star formation processes and the fueling of QSO accretion. In addition, some sort of starburst activity is inevitable when an AGN forms, since vast amounts of material (e.g., triggered by merging) are funneled into the core of the galaxy giving rise ultimately to the central engine.  It is then likely that this extended obscuring starburst is a common feature of all AGN during at least some phase, although for the most massive and luminous objects it might last only for a short time.  We make here the proposal that the formation of massive stars in the inner $\\sim 100$ pc of a central massive black hole is instrumental in both fuelling the central engine by accretion and obscuring it by distributing cold clouds all around that engine. Since most of the X-ray accretion power in the Universe emerges in the XRB (see e.g. Chen et al 1997) then it is in the obscured starburst mode that most of it takes place.  An obscuring starburst also accounts for the optical, narrow-line, appearance of a population of faint hard X-ray sources (NLXGs; Boyle et al 1995; Roche et al 1995; Carballo et al 1995; Griffiths et al 1996; McHardy et al 1998; Hasinger 1996; Almaini et al 1996; Romero-Colmenero et al 1996) which dominate at faint flux levels and so may provide most of the XRB (see, however, Hasinger et al 1998 and Schmidt et al 1998 who cast some doubts on the reality of these objects as a class different to AGN). The broad-line Seyfert spectrum will not be detectable, unless we have a favourable line of sight, or observe in the mid-infrared in an object where the column density is not too high (Granato, Danese \\& Franceschini 1997). The combination of a starburst spectrum and a Seyfert narrow-line spectrum will make classification of the objects ambiguous (Iwasawa et al 1997b).  There are various examples of galaxies whose spectrum is dominated by a starburst/LINER component at all wavelengths except at hard X-rays where the AGN shows up, obscuration preventing its detection at lower energies. Among them, NGC4945 is obscured by a column $\\sim 5\\times 10^{24}\\, \\pcmsq$ (Iwasawa et al 1993) and NGC 6240 by a column $> {\\rm few}\\times 10^{24}\\, \\pcmsq$ (Iwasawa \\& Comastri 1998). ",
        "conclusions": "Nuclear starbursts can play a major role in obscuring low-luminosity AGN.  The emerging X-ray spectrum will be moderately to highly absorbed by the cold and warm gas in the star cluster, thus providing the flat spectral shape that is needed to produce the XRB.  The column of absorbing gas is likely to vary vith viewing direction in a `random' way due to supernovae produced in the starburst.  The average high fraction of the sky that the majority of the faint X-ray sources have to see covered by absorbing gas arises in material related to the star cluster rather than in a hypothetical torus, whose thickness to radius ratio must be large. Arguments favouring a geometry for the obscuring material more isotropic than toroidal have been put forward before. Turner et al (1998)  find no obvious dependence of the X-ray properties of NLXGs with the orientation of the host galaxy. Iwasawa et al (1995) also suggested that the obscuring material in the Seyfert 2 galaxy $IRAS$ 18325-5926, which is obscured by a column in excess of $10^{22}\\, \\pcmsq$ but shows no X-ray reflected component, is likely to be isotropically distributed around the central engine.  Indeed the question remains, as in all obscured AGN models for the XRB, as to why the distribution of obscuring material along different lines of sight in the population of these objects is exactly tuned in such a way that it gives rise to the featureless spectrum of the XRB. There is no physical principle which leads to this and some fine-tuning of this distribution is required to produce a detailed model for the XRB.  The optical/UV properties of the highly absorbed AGN will be dominated by the starburst itself.  Collisions between cold clouds, occuring at several $\\times 100 \\kmps$ will lead to radiative shocks with some emission at soft X-ray energies.  Therefore, the obscuring material will also contribute to the observed soft X-ray emissivity, making some of these objects visible in the ROSAT PSPC band in spite of being absorbed.  Near infra-red spectroscopy of NLXGs also lends support to this hybrid (i.e., accretion onto a black hole plus an obscuring starburst) hypothesis. While a number of NLXGs show unambiguous evidence for hidden AGN (eg. highly ionized coronal lines), the non-detection of broad Paschen $\\alpha$ requires obscuring columns with $N_H > 10^{23}$~cm$^{-2}$ (Almaini et al 1998). Such thick columns are inconsistent with the large X-ray luminosities observed in ROSAT, unless there is an additional source of soft X-ray flux. A contribution from the activity in the obscuring starburst provides a natural explanation.  Gas columns along typical lines of sight in these objects are modest and for normal dust to gas ratios the absorbing material will be optically thin in the mid- and far-infrared (Granato et al 1997) where most of UV and soft X-ray reprocessed radiation will be emitted in a rather isotropic fashion.  Surveys at these wavelengths (particularly at $10-15\\, \\mu$m) should be especially efficient in finding these objects. Barcons et al (1995) measured the average ratio of X-ray (2-10~keV) to mid-infrared emission (the {\\it IRAS} 12$\\mu$m band) for 12$\\mu$m-selected type 1 and type 2 AGN.  The starburst obscured AGN are likely to have $f(5\\, \\keV)/f(12\\, \\mu{\\rm m})\\sim 2\\times 10^{-7}$, flux ratio that was found for Seyfert 2 galaxies, and therefore an obscured AGN with a 2-10~keV X-ray flux of $10^{-14}\\, {\\rm erg}\\, {\\rm cm}^{-2}\\, {\\rm s}^{-1}$ is likely to emit at a level of $\\sim 3$~mJy at 10-15$\\mu$m.  The {\\it European Large Area Infrared Space Observatory Survey} (ELAIS, Oliver et al 1997) is reaching a sensitivity of $\\sim 2$~mJy at 15$\\mu$m over a survey area of 13${\\rm deg}^2$ (some 1000-2000 objects are expected in total). Even if at a 2-10~keV flux of $10^{-14}\\, \\ergpcmsqps$ only a minor fraction of the sources are starburst obscured AGN, they will certainly show up in large numbers in the ELAIS survey.  The recently commissioned Sub-mm Common User Bolometer Array (SCUBA) at the James Clerk Maxwell Telescope in Hawaii offers an ideal opportunity to detect obscured AGN at high redshifts. The re-radiation of nuclear energy in the far infra-red is expected to produce a thermal emission peak at $\\lambda \\sim 60-100\\mu$m, which is redshifted into the sub-millimetre region. A starburst-obscured AGN at redshift $z=2$ with a 2-10~keV X-ray flux of $5\\times10^{-15}\\, \\ergpcmsqps$ is expected to give a 350$\\mu$m flux of $\\sim 20$mJy. Deep SCUBA observations of hard X-ray selected fields (eg. with ASCA, AXAF or XMM) should reveal these obscured objects.  As stated, we expect that the very massive black holes will appear as quasars. Sometimes however a massive black hole ($M_{BH}>10^8\\Msun$) may have a more massive surrounding star cluster in which case our scenario may apply and the active nucleus will be hidden. This may explain the low level of X-ray emission from IRAS galaxies in general and the hyperluminous galaxies (e.g. IRAS F15307+3252) in particular (Fabian et al 1996). The X-ray upper limit to that object requires less than $2\\times 10^{-4}$ of its power to emerge as X-rays which means that any quasar-like nucleus must be smothered from all directions.   \\begin{figure} \\centerline{\\psfig{figure=obscured.fig2.ps,width=0.5\\textwidth,angle=270}} \\caption{The energy content of the XRB per unit logarithmic energy interval (solid line) as a function of photon energy.  The dashed curve shows the contribution of unabsorbed AGN, where it is assumed that they make 50 per cent of the XRB at 1~keV. There is probably a roll-over in their contribution at energies $> 30-50\\, \\keV$ which is the exponential cutoff observed in local AGN at $\\sim 100\\, \\keV$ redshifted out to $z\\sim 2$. The dotted line shows the total energy produced by all objects, before obscuration is taken into account.} \\end{figure}  A geometrically thin disk structure is expected in obscured AGN in our model on scales of $<1\\pc$ where angular momentum begins to be important. Compton scattering, in some cases in an optically-thick medium, is necessary to explain the X-ray emission of some Seyfert 2 galaxies (see Iwasawa et al 1997a for a detailed study of NGC 1068 and Turner et al 1997 for Mrk 3 among other examples). In the case of NGC 1068 the pc-scale disk structure has been mapped at radio wavelengths (Gallimore, Baum \\& O'Dea 1997). Our model can therefore account for both Compton-thin and Compton-thick Seyfert 2 galaxies. Much of the extended torus attributed to these objects is just the starburst obscuration, but a planar Compton-thick inner part lies along our line of sight in some of them.   If the model proposed here describes correctly the sources of the residual XRB, it has then important implications on how accretion occurs in the Universe.  Fig.~2 shows the distribution of the energy in the XRB as a function of photon energy in the range 1 to 100 keV (Fabian \\& Barcons 1992). Unobscured AGN (QSOs and Seyfert 1s) make $\\sim 50$ per cent of the XRB at 1~keV (Hasinger 1996, McHardy et al 1998, Hasinger et al 1998, Schmidt et al 1998), with an energy spectral index $\\sim 1$. If the `shoulder' in the XRB (Fig.~2)  is due to obscuration (mostly photoelectric absorption of the photons with energies $< 30\\, \\keV$), the energy content at 30 keV provides a measure of the total energy produced by accretion in AGN, for the same underlying power law spectrum.  What is remarkable then, is that only 10 per cent of the accretion occurs in unobscured objects (assuming that before it is absorbed by the surrounding material, the spectrum generated by accretion in both cases is similar). The remaining 90 per cent occurs at sub-Eddington accretion rates, and a relevant fraction is absorbed and re-radiated at longer wavelengths, particularly in the infrared. This means that most of the sky as seen from an average active nucleus is obscured.  Moreover, given the requirement that these low luminosity AGN are at least one order of magnitude fainter than typical broad line objects (see Section 2), we predict that starburst obscured AGN may outnumber brighter, unobscured AGN by a factor of 100 or more if they are to account for the remainder of the XRB. This estimate is in good agreement with the number count predictions for QSOs and NLXGs based on extrapolating deep ROSAT observations (Almaini \\& Fabian 1997).  If most of accretion in the Universe is highly obscured, then the amount of emitted power per unit galaxy based on optical or UV QSO luminosity functions (So{\\l}tan 1982, Phinney 1997), and therefore the mass in black holes in AGN, might have been underestimated. This is due to the fact that 90 per cent of the accretion power would be obscured and re-radiated in the infrared.  Also, if most AGN are Advection Dominated Accretion Flows (ADAFs), as it has been proposed by Di Matteo \\& Fabian (1997) as the source of the XRB, the implied low mass to energy conversion efficiency also means that the black hole masses would have to be larger, in agreement with the local estimates."
    },
    "9803/astro-ph9803130_arXiv.txt": {
        "abstract": "\\noindent High resolution, deep radio images are presented for two distant radio galaxies, 3C324 ($z=1.206$) and 3C368 ($z=1.132$), which are both prime examples of the radio--optical alignment effect seen in powerful radio galaxies with redshifts $z \\gta 0.6$. Radio observations were made using the Very Large Array in A--array configuration at 5 and 8\\,GHz, and using the MERLIN array at 1.4 and 1.65\\,GHz. Radio spectral index, radio polarisation, and rotation measure maps are presented for both sources.  \\noindent Radio core candidates are detected in each source, and by aligning these with the centroid of the infrared emission the radio and the optical\\,/\\,infrared images can be related astrometrically with 0.1 arcsec accuracy. In each source the radio core is located at a minimum of the optical emission, probably associated with a central dust lane. Both sources also exhibit radio jets which lie along the directions of the bright strings of optical knots seen in high resolution Hubble Space Telescope images. The northern arm of 3C368 shows a close correlation between the radio and optical emission, whilst along the jet direction of 3C324 the bright radio and optical knots are co--linear but not co--spatial. These indicate that interactions between the radio jet and its environment play a key role in producing the excess ultraviolet emission of these sources, but that the detailed mechanisms vary from source to source.  \\noindent 3C368 is strongly depolarised and has an average rest--frame rotation measure of a few hundred rad\\,m$^{-2}$, reaching about 1000\\,rad\\,m$^{-2}$ close to the most depolarised regions.  3C324 has weaker depolarisation, and an average rest--frame rotation measure of between 100 and 200\\,rad\\,m$^{-2}$. Both sources show large gradients in their rotation measure structures, with variations of up to 1000\\,rad\\,m$^{-2}$ over distances of about 10\\,kpc. ",
        "introduction": "\\label{intro}  The discovery that the optical, ultraviolet and line emission of powerful high redshift radio galaxies is elongated and aligned along the direction of the radio axis (McCarthy \\etal\\ 1987, Chambers \\etal\\ 1987)\\nocite{mcc87,cha87}, a phenomenon known as the `alignment effect', indicates a close association between radio source activity and the optical emission of these galaxies. A number of processes have been proposed to account for these radio--optical correlations (see McCarthy 1993 for a review),\\nocite{mcc93} the most promising models falling into two broad categories: interaction models in which the correlations arise directly through interactions between the radio jet and its environment, and illumination models in which radiation from a partially obscured active galactic nucleus (AGN) both photo--ionises the emission line gas within an ionisation cone and is scattered by dust grains or electrons producing the optical and ultraviolet alignment.  One of the earliest models for the alignment effect, which has remained popular, is star--formation induced by the passage of the radio jets (e.g. Rees 1989).\\nocite{ree89} This process has been observed in jet--cloud interactions at low redshifts, for example in Minkowski's object \\cite{bro85,bre85} and in the radio lobe of 3C285 \\cite{bre93}. There is quite convincing evidence that it also occurs in at least some distant radio galaxies (e.g. Best \\etal\\ 1997a, Dey \\etal\\ 1997)\\nocite{bes97b,dey97} but, as yet, direct evidence for the presence of young stars in the aligned structures is limited.  A large proportion of the powerful distant radio galaxies which have been studied using spectropolarimetry have optical emission which is polarised at up to the 15\\% level, and scattered broad lines in their spectra (e.g. Dey \\etal\\ 1996, Cimatti \\etal\\ 1996, 1997, and references therein)\\nocite{cim96,dey96,cim97a}.  A significant proportion of the aligned emission of these galaxies must therefore be scattered light from an obscured AGN. In an optically unbiased sample of distant radio galaxies, however, Tadhunter \\etal\\ \\shortcite{tad97} detected significant polarisation in only 40\\% of the galaxies although most showed large ultraviolet excesses.  Nebular continuum emission provides another important alignment mechanism. Dickson \\etal\\ \\shortcite{dic95} showed that a combination of free--free, free--bound, and two--photon continuum, as well as the Balmer forest, can contribute a significant proportion (5--40\\%) of the ultraviolet emission from the nuclear regions of powerful radio galaxies, and this may be even higher in the extended aligned emission regions (e.g. Stockton \\etal\\ 1996)\\nocite{sto96a}.  Since either or both of photo--ionisation by the AGN and shock--ionisation by the radio jet may be responsible for exciting the gas, distinguishing between the jet--cloud interaction model and the illumination model is made more difficult.  The structure and properties of the radio emission from these distant radio galaxies can offer important clues to the nature of the radio--optical correlations.  Hubble Space Telescope (HST) observations of 28 radio galaxies at redshifts $z \\sim 1$, selected from the revised 3CR catalogue of Laing \\etal\\ \\shortcite{lai83}, have shown that for the majority of small radio sources the alignment effect manifests itself in the form of strings of bright knots tightly aligned between the two radio lobes (Best \\etal\\ 1996, 1997b).\\nocite{bes96a,bes97c} Deep, high resolution radio observations of these sources provide a critical test of the interaction\\,/\\,illumination models: if the locations of the radio jets and the strings of optical activity are co--linear, then the influence of the radio jets must play a key role in the alignment effect. This would not, however, provide direct proof of the jet--induced star formation hypothesis: the emission may be nebular continuum emission from warm gas shocked by the radio jets, or may be associated with an increased scattering efficiency in these regions, for example by jet shocks breaking up cold gas clouds and exposing previously hidden dust grains, thereby increasing the surface area for scattering along the jet axis \\cite{bre96b}.  One of the major problems in interpreting the radio--optical structures of distant radio galaxies has been accurate relative astrometry of the radio and optical\\,/\\,infrared frames of reference. Radio observations are automatically in the International Celestial Reference Frame (ICRF), with positional errors of about 0.01 arcsec, but the absolute positions of the HST images are uncertain at the 1 arcsecond level \\cite{lat97}. Such astrometric uncertainties mean that bright radio knots may be either correlated or anti--correlated with the position of the luminous optical emission regions. This problem can be alleviated somewhat by multi--frequency radio observations of sufficient sensitivity to detect a flat--spectrum central radio core: this core is expected to be coincident with the nucleus of the galaxy, which itself is located roughly at the centroid of the infrared image of the galaxy. The optical (rest--frame ultraviolet) emission cannot be used for this procedure, since it is often severely affected by both dust extinction (e.g. de Koff \\etal\\ 1996, Best \\etal\\ 1997b)\\nocite{bes97c,dek96} and the aligned emission. The infrared images also contain a component of aligned emission (e.g. Eisenhardt \\& Chokshi 1990; Rigler \\etal\\ 1992; Dunlop \\& Peacock 1993; Best \\etal\\ 1997b,1998)\\nocite{eis90,rig92,dun93,bes98d} but at a significantly lower level, and are seen to show a sharp central peak which can be astrometrically aligned with the radio core.  It should be noted that this process can never be 100\\% reliable, owing to the ambiguity of identifying radio core components based upon spectral index information alone; the case of 3C356 (Fernini \\etal\\ 1993, Best \\etal\\ 1997a) is a case in point. For the radio galaxies in the current paper, however, further radio--optical correlations are observed when this process is employed (see Sections 3 and 4), indicating that the core identifications are most likely correct.  The properties of the radio emission also provide a probe of the large--scale physical environment of the radio source. There has been growing evidence that powerful radio galaxies at large redshifts lie in rich (proto--) cluster environments, based upon measures of galaxy cross--correlation functions (e.g. Yates \\etal\\ 1989)\\nocite{yat89}, detections of powerful extended X--ray emission from the vicinity of distant radio sources (e.g. Crawford and Fabian 1996)\\nocite{cra96b}, detections of companion galaxies in narrow--band images and with spectroscopy (e.g. McCarthy 1988, Dickinson \\etal\\ 1997)\\nocite{mcc88,dic97a}, and the very bright infrared magnitudes and large characteristic sizes of the radio galaxies themselves, suggesting that they are as large and luminous as brightest cluster galaxies (see Best \\etal\\ 1998 for a review)\\nocite{bes98d}. Radio data provide further evidence in support of this hypothesis. At low redshifts, radio galaxies in rich cluster environments, such as Cygnus A, tend to have very large rotation measures, $RM \\gta 1000$\\,rad\\,m$^{-2}$ \\cite{dre87,tay94}. Carilli \\etal\\ \\shortcite{car97} have studied a sample of 37 radio galaxies with redshifts $z > 2$, and detect rotation measures exceeding 1000\\,rad\\,m$^{-2}$ in 19\\% of them. They consider this percentage to be a conservative lower limit, and interpret these large rotation measures as arising from hot, magnetised (proto--)cluster atmospheres surrounding the sources, with field strengths $\\sim 1$\\,nT.  Detailed studies of individual sources with very high rotation measures support this interpretation \\cite{car94a,pen97}. In comparison with these very distant sources, the polarisation properties of radio galaxies with redshift $z \\sim 1$ have not been well studied at high angular resolution (although see Pedelty \\etal\\ 1989, Liu and Pooley 1991, Fernini \\etal\\ 1993, and Johnson \\etal\\ 1995)\\nocite{ped89a,joh95,fer93,liu91a}.  We have selected two radio galaxies from the 3CR sample, with redshifts in excess of one: 3C324 ($z=1.206$) and 3C368 ($z=1.132$). These galaxies are both prime examples of the alignment effect; their Hubble Space Telescope (HST) images show highly extended optical morphologies with strings of bright knots close to the radio axis \\cite{lon95}. We have carried out sensitive, multi--frequency, polarimetric radio observations of these radio sources at high angular resolution using the VLA and MERLIN. A description of the observations and the data reduction techniques is given in Section~\\ref{observs}. In Section~\\ref{sec3c324}, we present the radio data for 3C324, and make a comparison with the optical and infrared images. This is repeated for 3C368 in Section~\\ref{sec3c368}. These sources are compared in Section~\\ref{concs}, in which we also summarise our conclusions.  Throughout the paper we assume $H_0 = 50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $\\Omega = 1$, and all positions are given in equinox J2000 coordinates. ",
        "conclusions": "\\label{concs}  We summarise below the main observational results of this work:  \\begin{enumerate} \\item We have detected radio core candidates in each of 3C324 and 3C368. Aligning these relatively flat spectrum knots with the centre of the infrared images enables the radio and the optical\\,/\\,infrared frames of references to be astrometrically aligned to an accuracy of about 0.1 arcsec.  \\item The radio cores of each source lie at a minimum of the optical emission, probably associated with obscuration by a central dust lane. Dust may also play an important role in determining the extended morphology of the optical emission.  \\item Radio jets are detected in both sources, and in each case are co-linear with the strings of bright optical knots.  \\item These observations sort out the disparity between lower angular resolution observations of the polarisation structures of these radio galaxies \\cite{ped89a,fer93}. The higher angular resolution reveals very large gradients in the rotation measure structures, reaching up to 1000\\,rad\\,m$^{-2}$ over a distance of order 10\\,kpc, and regions of strong depolarisation.  \\item There is a tight correlation between the depolarisation measures and the gradients in the rotation measure, suggesting that the depolarisation is external, and caused by variations within each beamwidth of the Faraday depth of material in the vicinity of the radio source. \\end{enumerate}  The striking co-linearity of the radio jet with the bright optical knots in the two sources indicate that interactions of the radio jet itself, rather than just the AGN or the radio cocoon, must play a critical role in the origin of the aligned optical emission. Of interest is that, given these strong interactions between the radio jets and their environment, more enhanced radio emission is not seen in these regions, in contrast to that observed in bright radio knots at the sites of jet--cloud interactions at low redshifts (e.g. 4C29.30, van Breugel \\etal\\ 1986). This suggests that the more powerful radio jets in these high redshift sources are not as strongly disrupted by their interactions.  Despite the many similarities between the two sources, there are equally important differences. The northern arm of 3C368 shows a tight correlation between the regions of radio and optical emission, whilst in 3C324 although these are co-linear they are not co-spatial. The nature of the optical (rest--frame ultraviolet) emission of the two galaxies is also very different. 3C324 has a high percentage of spatially extended optical polarisation, $P \\approx 11\\%$, indicating that illumination by the central AGN must be important in this source \\cite{cim96}. In contrast, HST and Keck observations of 3C368 have detected no polarised optical emission \\cite{bre96c}.  3C368 is an exceptional source, being probably the most aligned of all of the 3CR radio galaxies and possessing the highest emission line flux. The extent of its radio emission is 73\\,kpc, only slightly greater than that of the optical emission. Its structure may be interpreted in terms of it being a relatively young radio source (a few times $10^6$ years), in which on--going interactions between the radio jet and the interstellar medium remain the dominant effects. In the northern region of this source, where the bright optical knots lie coincident with the radio jet, the nebular continuum emission is very luminous and the emission line velocities are high \\cite{dic95,sto96a}.  Interactions between the radio jet and its environment, possibly including a companion galaxy (see Section~\\ref{368radopt}) may brighten the radio jet and accelerate and shock--ionise the warm emission line gas, explaining all of these features. The strong correlation between the radio and optical structures cannot be explained if the gas responsible for the nebular continuum is photoionised by the AGN. Jet--induced star formation in this region cannot be excluded, but there is no direct evidence for this. Any small contribution of scattered light might not be detectable due to the dilution of the polarisation by the very strong nebular continuum contribution.  The radio emission of 3C324 is more extended and reaches well beyond the region of bright optical emission. In this source, we may be observing the residual effects after the radio jets have forced a passage through the host galaxy. In the aftermath of the jet shocks the nebular continuum contribution has decreased, and now some scattered light is detected. Cimatti \\etal\\ \\shortcite{cim96} calculate that at most 30 to 50\\% of the optical flux density can be associated with scattered radiation, with the rest likely to be nebular continuum emission or light from a young starburst. They suggest that the knot close to the western lobe may currently be undergoing a burst of star formation at a rate of $70 M_{\\odot}$\\,yr$^{-1}$. Most of the optical emission from 3C324 arises from the bright knotty structures along the radio axis, with no evidence for the scattered emission being evenly distributed over an ionisation cone structure. Therefore, the scattering material in this source must also be preferentially located along the jet direction.  Two possible mechanisms for this latter effect would the production of dust in a region of star formation induced by the radio jet, or the disintegration of cooled clumps of optically thick gas by the jet, and the exposure of previously hidden dust grains at their centre \\cite{bre96b}.  \\smallskip  At low redshifts, rotation measure values and variations in excess of a few hundred rad\\,m$^{-2}$ are seen in two types of radio source, compact steep spectrum radio sources in which the radio lobes lie within the host galaxy (e.g. Taylor \\etal\\ 1992)\\nocite{tay92}, and FR\\,II radio sources located in clusters with luminous X--ray cooling flows. For the former, Garrington and Akujor \\shortcite{gar96} have suggested that the rotation measure correlates with linear size, with the dense gas located towards the centre of the host galaxy being responsible for the Faraday rotation. For the latter, Taylor \\etal\\ \\shortcite{tay94} found a correlation between the rotation measure and the cooling flow rate, suggesting that the intracluster medium is responsible.  We have argued that the large rotation measures observed in the southern lobe of 3C368 are due to emission line gas associated with the host galaxy. The northern lobe of this source, however, and both lobes of 3C324, lie distant from the host galaxy. Dickinson \\shortcite{dic97a} has reported the detection of extended X--ray emission at a luminosity of $L_{\\rm X} = (8.1 \\pm 1.6) \\times 10^{44}$\\,erg\\,s$^{-1}$ from 3C324, a value comparable to that of the Coma cluster. He associates this with cooling intracluster gas. Similarly, 3C368 has been detected using the ROSAT PSPC, with a signal--to--noise ratio of 5, giving $L_{\\rm X} \\sim 1.7 \\times 10^{44}$\\,erg\\,s$^{-1}$ \\cite{cra95}. If these X--ray luminosities are associated with cooling flows, then the cooling flow rates and rotation measures lie roughly upon the correlation of Taylor \\etal\\ \\shortcite{tay94}, suggesting that the high rotation measures of 3C324 and of the regions of 3C368 away from the emission line gas are indeed associated with their intracluster gas. These rotation measure values are not as extreme as, for example, the 20000\\,rad\\,m$^{-2}$ seen in 3C295, at the centre of a very rich cluster of galaxies at $z = 0.461$ \\cite{per91}, but the large values and gradients provide further support for the hypothesis that powerful distant radio galaxies lie in at least moderately rich young cluster environments."
    },
    "9803/chao-dyn9803022_arXiv.txt": {
        "abstract": "The energy spectrum and the nolinear cascade rates of MHD turbulence is not clearly understood. We have addressed this problem using direct numerical simulation and analytical calculations. Our numerical simulations indicate that Kolmogorov-like phenomenology with $k^{-5/3}$ energy spectrum, rather than Kraichnan's $k^{-3/2}$, appears to be applicable in MHD turbulence. Here, we also construct a self-consistent renomalization group procedure in which the mean magnetic field gets renormalized, which in turns yields $k^{-5/3}$ energy spectrum. The numerical simulations also show that the fluid energy is transferred to magnetic energy. This result could shed light on the generation magnetic field as in dynamo mechanism. ",
        "introduction": "The fluid parcels in a turbulent flow have random motion. However, the random flow velocities obeys certain properties. Since most of the talks in the conference dealt with chaos which also yields random signals, it is important to contrast the difference between chaos and turbulence. Chaos can occurs in a nonlinear system with a few (3-10) degrees of freedom, but a turbulent system has many (order millions or more) degrees of freedom. Also, a turbulent system may be chaotic; may not be chaotic such as in structures (e.g., vortex street); or it may have both chaos and structure coexisting with each other.  Turbulence is ubiquitous and is of very practical importance. Some of the major applications are in mixing, aeroplane and high speed vehicle design, atmospheric flows and weather prediction, astrophysical objects like stars, jets, interplanetary medium etc. Even with its wide industrial, practical, and theoretical importance, the understanding of turbulence is very weak. There are many empirical laws from experiments and simulations. There are some phenomenologies, the most famous among these is by Kolmogorov. There only a few mathematically rigorous calculations, primarily Kraichnan's direct interaction approximation, calculations based on renormalization group etc. In this paper we will focus on some of the statistical properties of velocity and magnetic fields in a turbulent magnetohydrodynamic (MHD) plasma. We have attempted to review in a concise manner some of the phenomenological, numerical, analytical work, and observations from the solar wind, with an emphasis on our work (with D. A. Roberts, M. L. Goldstein, J. K. Bhattacharjee, V. Eswaran).  Outline of the paper is as follows. Since the existing MHD turbulence phenomenologies are motivated by Kolmogorov's fluid turbulence phenomenology, we review Kolmogorov's arguments in section 2. In section 3 we review the existing MHD turbulence phenomenologies. Section 4 contains the numerical results, which are compared with the predictions of the turbulence phenomenologies. In section 5 we briefly report the cascade rates of fluid and magnetic energies. Section 6 contains a renormalization group scheme which provides a self-consistent procedure to obtain the effective mean magnetic field. Section 7 contains conclusions. ",
        "conclusions": "In this paper we have reviewed some of the phenomenological, observational, numerical, and analytic work in the statistical theory of MHD turbulence, with emphasis on our studies. We find that the solar wind observations and the numerical results are inconsistent with the predictions of the existing MHD turbulence phenomenologies. The energy spectrum and the cascade rates appear to be closer the predictions of Kolmogorov-like phenomenology even when the mean magnetic field or the magnetic field of the largest eddies is large compared to the inertial range velocity and magnetic field fluctuations (a region where Kolmogorov-like theory is not expected to hold).  We have attempted to obtain Kolmogorov-like energy spectrum in MHD turbulence in presence of arbitrary $B_0$ by postulating that the effective $% B_0$ is scale dependent, unlike what has been taken by Kraichnan and Dobrowolny et al. We have constructed a renormalization group scheme and shown that the self consistent $B_0(k)\\propto k^{-1/3}$ and $E(k)\\propto k^{-5/3}$. This analysis has been worked out when $E^{+}=E^{-}$ and $r_A=1.$ The generalization to arbitrary parameters, or at least to the limiting cases, is also planned. We will be able to the get the Kolmogorov's constants for MHD\\ turbulence analytically using this procedure; these constants will be useful for the large-eddy-simulations (LES) of MHD turbulence.  We have also carried out a preliminary study of the cascade rates of fluid and magnetic energies. We find that there is a net transfer of fluid energy to magnetic energy. We are investigating the parameter regimes in which these cascade rates are large enough (more than dissipation rates) so that the total magnetic energy increases with time, what is found in dynamo theory.  Lastly, we would like to mention that the MHD turbulence theories are very important for modelling various astrophysical phenomena and plasma processes. We have estimated turbulent heating, nonclassical viscosity and resistivity of the solar wind using the MHD turbulence phenomenologies \\cite {Verma:sw}. In this light, search for a satisfactory theory of MHD turbulence appears quite important.  We thank all our collaborators, D. A. Roberts, M. L. Goldstein, J. K. Bhattacharjee, and V. Eswaran. MKV\\ also thanks V. Subrahmanyam, M. Barma, V. Ravishankar, and D. Sa for numerous useful discussions."
    },
    "9803/nucl-th9803012_arXiv.txt": {
        "abstract": "s{We discuss models to calculate one-- and two--neutron capture reactions on light nuclei. These are applied to calculate the reaction rates of $^{15}$N(n,$\\gamma$)$^{16}$N, $^{16}$N(n,$\\gamma$)$^{17}$N and $^4$He(2n,$\\gamma$)$^6$He. The possible astrophysical importance is discussed.} ",
        "introduction": "Neutron capture reactions on light nuclei play a role in various astrophysical scenarios. In the framework of Inhomogeneous Big Bang Models a high neutron flux can bridge the mass 5 and mass 8 gaps. Subsequent neutron capture reactions may trigger a primordial r--process \\cite{rau94}. Another site for neutron capture reactions is the high--entropy bubble formed during a type II supernova \\cite{woo92,mey92,woo94}. Due to the photodisintegration at very high temperatures an $\\alpha$--rich environment is created. When the temperature has dropped heavier elements can be built up mainly by $\\alpha$-- and neutron--capture reactions. Again a critical question is how the mass 5 and mass 8 gaps can be bridged.  In Section 2 we describe our model to calculate one--neutron capture reactions. This model is applied to the reactions $^{15}$N(n,$\\gamma$)$^{16}$N and $^{16}$N(n,$\\gamma$)$^{17}$N. In Section 3 we will describe the theory of calculating a two--neutron capture reaction in a three--body model. We calculate the reaction rate of $^4$He(2n,$\\gamma$)$^6$He and compare the result with other works. In the final chapter we discuss and summarize our results. ",
        "conclusions": "We calculated reaction rates for one--neutron capture on $^{15}$N and $^{16}$N and two--neutron capture on $^4$He. The rate for $^{15}$N(n,$\\gamma$)$^{16}$N is in good agreement with both previous calculations and experimental data. In general the neutron capture rates on stable targets are quite well known. The reaction $^{16}$N(n,$\\gamma$)$^{17}$N is an example of a capture reaction on an unstable target. We find a considerable enhancement to a previous calculation. This enhancement can be also be observed at various other capture reactions on unstable targets in this mass range \\cite{her98}. This fact could influence the reaction path both in the nucleosynthesis of Inhomogeneous Big Bang Models and in the alpha--rich freeze--out of type II supernovae.  For the reaction $^4$He(2n,$\\gamma$)$^6$He we also find a strong enhancement of the rate compared to the calculation of Fowler \\cite{fowler75}. But from the calculation of the inverse photodisintegration rate by Danilin {\\it et al}~\\cite{danilin93} we deduce an even much higher rate. This is probably due to the fact that this rate also includes the simultaneous capture of two neutrons. But even with this enhanced rate the path via $^4$He(2n,$\\gamma$)$^6$He is dominated by the reaction $^4$He($\\alpha$n,$\\gamma$)$^9$Be at conditions typical for the alpha--rich freeze--out. This is due to the effective destruction of $^6$He through photodisintegration."
    },
    "9803/astro-ph9803183_arXiv.txt": {
        "abstract": "According to a currently popular paradigm, nuclear activity in quasars is sustained {\\it via } accretion of material onto super-massive black holes located at the quasar nuclei. A useful tracer of the gravitational field in the vicinity of such central black holes is available in the form of extremely dense gas clouds within the broad emission-line region (BLR) on the scale of $\\sim 1~$parsec. Likewise, the radio sizes of the lobe-dominated radio sources are believed to provide a useful statistical indicator of their ages.  Using two homogeneously observed (and processed) sets of lobe-dominated radio-loud quasars, taken from literature, we show that a positive correlation exists between the radio sizes of the quasars and the widths of their broad $H\\beta$ emission lines, and this correlation is found to be significantly stronger than the other well known correlations involving radio size. This statistical correlation is shown to be consistent with the largest (and, hence, very possibly the oldest) radio sources harboring typically an order-of-magnitude more massive central engines, as compared to the physically smaller and, hence, probably much younger radio sources. This inference is basically in accord with the \"accreting central engine\" picture for the radio-loud quasars. ",
        "introduction": "Being too compact to be resolved with even the most advanced optical telescopes, the structure and kinematics of the broad emission line region (BLR), a prime feature of quasars, continues to be a major enigma in the AGN research (see, e.g., Brotherton, 1996; Marziani et al., 1996; Corbin, 1992).  Nonetheless, it is a key ingredient to the theoretical models that seek to explain various observable properties of quasars, e.g., their intense $\\gamma$-ray emission (e.g., Dermer \\& Schlickeiser, 1993; Ghisellini \\& Madau, 1996).  Deciphering the BLR geometry has therefore been a key objective of many observational programmes.  One strategy for this is the so called `reverberation mapping' (e.g., Peterson, 1993; Koratkar \\& Gaskell, 1991), from which a positive dependence of the BLR size on the bolometric luminosity has been inferred:  $r~\\sim 0.06 L_{46}^{0.5}~pc$ (Netzer \\& Laor 1993), where $L_{46}$ is the luminosity expressed in the units of $10^{46} erg.s^{-1}$. \\begin{table*} \\label{ The sample of lobe dominated radio-loud quasars (48 quasars)} \\begin{tabular}{lrrcccrrc} \\multicolumn {9}{c}{\\bf Table 1: Sample of lobe-dominated quasars (42 quasars)}\\\\ \\multicolumn {9}{c}{}\\\\ \\hline \\hline \\multicolumn {1}{c}{QSO}&\\multicolumn {1} {c}{z} & \\multicolumn {1}{c} {LAS}& \\multicolumn {1}{c}{log($f_c$)}& \\multicolumn {1}{c} {$\\rm log(R_v)$}&\\multicolumn {1}{c} {$\\rm M_{v}$}&\\multicolumn {2}{c}{W($\\rm H\\beta$) (km s$^{-1}$)}&\\multicolumn {1}{c}{ $\\xi$}\\\\ \\cline {7-8} &&\\multicolumn {1}{c} {(arcsec)}&&&&B(96)&JB(91)&\\\\ \\hline 0003+158  &0.450&36.0(1)&$-$0.38(a)&1.95&-26.0&4760&...&...\\\\ 0042+101  &0.583&59.0(3)&$-$0.50(b)&1.79&-24.8&...&17774&...\\\\ 0044+030  &0.624&18.6(5)&$-$0.42(a)&0.83&-27.2&5100&...&...\\\\ 0110+297  &0.363&76.2(2)&$-$0.80(b)&1.59&-24.8&...&8702&...\\\\ 0115+027  &0.670&13.1(2)&$-$0.50(b)& 2.25&-26.5&5000&7591&0.66\\\\ 0118+034  &0.765&45.0(4)&$-$1.20(b)&2.17&-25.1&...&22403&...\\\\ 0133+207  &0.425&68.0(2)&$-$1.00(b)&2.51&-24.0&...&17403&...\\\\ 0134+329  &0.367&1.3(3)&$-$1.14(a)&2.50&-25.7&3800&5863&0.65\\\\ 0405$-$123&0.575&31.7(6)&$-$0.23(a)&2.36&-28.4&4800&...&...\\\\ 0414$-$060&0.781&36.4(4)&$-$0.50(a)&1.68&-27.8&8200&16602&0.49\\\\ 0518+165  &0.759&0.8(9)&$-$0.90(b)&3.94&-24.1&...&4876&...\\\\ 0538+498  &0.545&0.2(7)&$-$1.40(b)&3.43&-24.2&...&3456&...\\\\ 0710+118  &0.768&48.0(2)&$-$1.85(a)&1.09&-27.1&20000&17774&1.13\\\\ 0800+608  &0.689&25.0(10)&$-$0.60(b)&2.65&-24.2&...&6912&...\\\\ 0837$-$120&0.200&169.0(2)&$-$0.70(a)&1.80&-24.7&6060&...&...\\\\ 0838+133  &0.680&10.0(7)&$-$0.50(b)&3.16&-25.4&3000&4197&0.71\\\\ 0903+169  &0.410&50.0(8)&$-$1.40(c)&1.70&-24.5&4400&...&...\\\\ 0952+097  &0.298&12.5(2)&$<$$-$0.50(c)&$<$1.90&-24.1&3800&...&...\\\\ 0955+326  &0.530 &1.0(3)&$-$0.07(a)&2.23&-27.1&1380&...&...\\\\ 1004+130  &0.240&115.0(2)&$-$1.68(a)&0.43&-25.7&6300&9998&0.63\\\\ 1007+417  &0.613&32.0(2)&$-$0.39(a)&2.17&-27.0&3560&6912&0.52\\\\ 1048$-$090&0.334&83.0(2)&$-$1.31(a)&1.66&-24.9&5620&...&...\\\\ 1100+772  &0.311&30.0(1)&$-$0.91(a)&1.63&-25.8&6160&7961&0.77\\\\ 1103$-$006&0.423&21.0(2)&$-$0.14(a)&2.20&-25.8&6560&...&...\\\\ 1111+408&0.730&13.2(2)&$-$1.80(b)&1.89&-25.3&...&9134&...\\\\ 1137+660&0.652&44.2(2)&$-$1.07(a)&1.83&-27.1&6060&8702&0.69\\\\ 1223+252&0.268&67.0(2)&$-$1.59(b)&0.99&-23.8&...&9319 &...\\\\ 1250+568&0.321&1.5(2)&$-$1.62(a)&2.01&-23.6&4560&6295&0.72\\\\ 1305+069&0.599&46.5(4)&$<$$-$1.20(a)&$<$1.52&-26.1& 6440&...&...\\\\ 1351+267&0.310&190.0(2)&$-$0.27(a)&1.72&-24.3&8600&...&...\\\\ 1425+267&0.366&230.0(1)&$-$0.40(a)&1.03&-26.2&9410&...&...\\\\ 1458+718&0.905&2.1(2)&$-$1.00(d)&2.74&-27.3&3000 &...&...\\\\ 1512+370&0.371&54.0(1)&$-$0.71(a)&1.88&-25.6&6810&...&...\\\\ 1545+210&0.264&70.0(1)&$-$1.32(a)&1.74&-24.4&7030&...&...\\\\ 1618+177&0.555&48.0(3)&$-$0.67(a)&1.98&-26.5&7000&...&...\\\\ 1622+238&0.927&21.7(2)&$-$1.72(d)&1.63&-26.7&7100&...&...\\\\ 1704+608&0.371&55.0(1)&$-$1.91(a)&0.73&-26.6&6560&...&...\\\\ 1742+617&0.523&40.0(11)& $-$2.00(b)&1.61&-24.0&...&9751&...\\\\ 1828+487&0.691&14.0(11)&$-$0.50(b)&4.05&-26.2&...&9998&...\\\\ 2135$-$147&0.201&150.0(2)&$-$1.13(a)&1.97&-24.9&7300&11479&0.63\\\\ 2251+113&0.323&9.8(2)& $-$1.52(a)&1.00&-25.8&4160&8702&0.48\\\\ 2308+098&0.432&108.0(1)&$-$0.78(a)&1.41&-26.3&7970&...&...\\\\ \\hline \\multicolumn {9}{l}{{\\bf References for LAS:} 1 : Kellermann et al. (1994), 2 : Nilsson et al. (1993), 3 : Singal (1988), }\\\\ \\multicolumn {9}{l}{4 : Kapahi (1995), 5 : Price et al. (1993), 6 : Morganti et al.(1993), 7 : Bogers et al. (1994),}\\\\ \\multicolumn {9}{l}{ 8 : Bridle et al (1994),9 : Akujor et al. (1993), 10 : Jackson et al. (1990), 11 : Reid et al. (1995)}\\\\ \\multicolumn {9}{l}{{\\bf References for $f_c$:} a : Wills \\& Browne (1986), b : Jackson \\& Browne (1991a), }\\\\ \\multicolumn {9}{l}{c : Brotherton (1996), d : Wills et al. (1992), e : Kellermann et al. (1994)}\\\\ \\end{tabular} \\end{table*}  An early indication about the BLR geometry came from an empirical study of a heterogeneously selected sample of radio-loud quasars with the characteristic core-lobe type radio structure. The study revealed a statistically significant anti-correlation between the {\\it prominence} of the radio core relative to the lobe, and the FWHM of the $H\\beta$ broad emission line ($W$). It was thus inferred that the BLR clouds are predominantly confined to a rotating disk-shaped region surrounding the quasar nucleus and oriented roughly perpendicular to the jet axis (Wills \\& Browne, 1986; hereafter WB86). The significance of this correlation was found to be considerably lower in a subsequent study, however (Jackson \\& Browne, 1991b). In a recent work, it has been proposed that the absolute {\\it visual} magnitude, $M_v$, of the quasar (plus its host galaxy) provides a more reliable  measure of the intrinsic power of the central engine, and therefore the beamed radio core flux normalized by $M_v$ is a better indicator of the orientation of the jet relative to the line-of-sight (Wills \\& Brotherton, 1995; Brotherton, 1996). Adopting this new parameter for the core-prominence and employing a larger sample of quasars , these authors have found a conspicuous anti-correlation between the core-prominence and the $W(H\\beta)$, thereby supporting the disk-like BLR geometry inferred earlier by WB86 (at least for the $H\\beta$ emitting clouds).  It is now widely believed that the large widths of the BLR emission lines are a manifestation of the deep gravitational potential at the centers of quasars, possibly due to super-massive black-holes. The accretion of the material postulated for sustaining the quasar luminosity is expected to steadily increase the mass of the central engine during the lifespan of the nuclear activity. If this conjecture is basically right, the question arises:  {\\sl do we see in the data any evidence for the postulated growth of the central mass} (e.g, in the dynamics of the BLR clouds) ? In the present study we shall examine this issue by employing published measurements of radio-loud, steep-spectrum quasars. While the needed reliable estimator of the age is generally not available for individual sources, the overall radio size can serve as a useful statistical measure of age (since most radio sources are believed to steadily grow in size with time; cf. Sect. 3). Thus, we wish to examine here the nature of any relationship between the observed linear sizes of steep-spectrum radio quasars and the widths of their broad H$\\beta$ emission lines. ",
        "conclusions": ""
    },
    "9803/astro-ph9803057_arXiv.txt": {
        "abstract": "The method of obtaining confidence intervals on a subset of the total number of parameters ($p$) of a model used for fitting X-ray spectra is to perturb the best-fitting model until, for each parameter, a range is found for which the change in the fit statistic is equal to some critical value. This critical value corresponds to the desired confidence level and is obtained from the $\\chi^{2}$ distribution for $q$ degrees of freedom, where $q$ is the number of interesting parameters. With the advent of better energy-resolution detectors, such as those onboard {\\it ASCA} ({\\it Advanced Satellite for Cosmology and Astrophysics}) it has become more common to fit complex models with narrow features, comparable to the instrumental energy resolution. To investigate whether this leads to significant non-Gaussian deviations between data and model, we use simulations based on \\asca data and we show that the method is still valid in such cases. We also investigate the weak-source limit as well as the case of obtaining upper limits on equivalents widths of weak emission lines and find that for all practical purposes the method gives the correct confidence ranges. However, upper limits on emission-line equivalent widths may be over-estimated in the extreme Poisson limit. ",
        "introduction": "\\label{intro}  The procedure for generating confidence intervals for the parameters of models used to fit X-ray spectra with the \\chisq  minimization technique is well established (e.g. Lampton, Margon and Bower 1976; Avni 1976). One uses the fact that if $\\chi ^{2} _{\\rm true}$ is the value of the statistic from a given experiment, calculated for the `true model' (known only to nature) and $\\chi_{\\rm min}^{2}$ is the value of the statistic obtained from the best-fitting model, then from many repetitions of the same experiment $\\Delta \\chi^{2} = \\chi^{2}_{\\rm true} - \\chi^{2} _{\\rm min}$ has a probably distribution like $\\chi ^{2}$ with $p$ degrees of freedom ($\\equiv \\chi^{2}_{p}$) where $p$ is the total number of free parameters. Confidence intervals generated for the $p$ parameters from the $\\chi^{2}_{p}$ distribution are then {\\it joint} confidence intervals for all $p$ parameters. If we are interested only in a subset $q$ of the $p$ parameters then the confidence intervals are generated from the $\\chi^{2}_{q}$ distribution. The crucial difference between $p$-parameter and, say, one-parameter errors is as follows. In the former case the $P\\%$ confidence intervals from $\\chi^{2}_{p}$ will {\\it simultaneously} enclose the true values of all parameters in $P\\%$ of all experiments. In the latter case, the $P\\%$ intervals from $\\chi^{2}_{1}$ will enclose the true values of any of the parameters in $P\\%$ of all experiments, {\\it but it will be a different $P\\%$ subset of experiments for each parameter}. In a particular case, the number of `interesting' parameters (i.e. $q$) is determined by the scientific problem being posed.  In principle one generates the $\\Delta \\chi ^{2}$ space by stepping through a $q$-dimensional grid of parameter values. The use of $\\chi ^{2}$ in model-fitting requires that there are a sufficient number of photons per energy bin for the statistical variations to be Gaussian. However, with the advent of better energy resolution X-ray detectors it is increasingly becoming the case that the Gaussian approximation cannot be made unless the spectrum is binned, sacrificing valuable information. In such cases the statistical variation in counts per bin is Poisson and one must use the maximum likelihood ratio (hereafter, `$C$ statistic') to optimize the model parameters (see Cash 1979). Cash (1979) demonstrated that the $C$ statistic can be used to generate confidence intervals in an analogous manner to $\\chi ^{2}$ since $\\Delta C$ has a probability distribution like $\\chi ^{2}$ except for terms of order $\\alpha/ \\sqrt{n}$ where $\\alpha$ depends on the model and $n$ is the number of photons carrying information about the parameter in question. Hereafter, we shall use $C$ for the sake of generality.  Lampton \\etal (1976) warned that (at that time), there was no general proof that the technique for projection of the subset of $q$ parameters did not depend explicitly on model linearity. Then, Avni (1976), using some non-linear models in  simulations of {\\it Uhuru} data, showed that the $\\chi^{2}_{q}$ region worked in these particular cases, but there was still no general proof. Such a proof was presented by Cash (1976), showing that the $\\chi^{2}_{q}$ region worked for any data set, {\\it provided that the deviations are Gaussian}. We must remember that X-ray detectors now have much better energy resolution and sensitivity than they did then and accordingly the models have become much more complex. Both Lampton \\etal (1976) and Avni (1976) used {\\it Uhuru} spectral responses with seven energy bins between 1 and 7 keV, and simple power-law or plasma models. It is not clear whether, for example, a model including a narrow emission line whose intrinsic width is comparable to the energy resolution of the detector, would introduce non-Gaussian deviations between model and data. The purpose of this paper is to check whether the method for parameter estimation currently in use gives the correct confidence intervals even with the new generation of improved energy-resolution instrumentation. We are particularly interested in models with features such as emission lines and absorption edges (which occur frequently in a wide range of X-ray sources), whose widths in energy space are comparable to the instrumental energy resolution. We also wish to check the Poisson regime, when the source count rate is low, and when upper limits must be obtained on the equivalent width of weak or non-detected emission-line features.  The structure of the paper is as follows: \\S \\ref{simulations} describes basic simulations and models used in the investigation; \\S \\ref{results} describes the basic results; \\S \\ref{lineew} demonstrates the equivalence of emission-line equivalent width and intensity confidence regions; \\S \\ref{weaksource} describes results for the extreme Poisson limit and \\S \\ref{weakline} describes results pertaining to measuring upper limits on weak emission-line features. Our conclusions are stated in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have verified that for all practical purposes, the method of generating confidence intervals for a subset, $q$, of $p$ model parameters, using the $\\chi^{2}_{q}$ distribution can still be used even if the model has components which are narrow compared to the instrumental energy resolution (such as emission lines and absorption edges). We have also investigated the weak-source and weak emission-line limits and find the method to work, except for the extreme Poisson limit when there may be one or less total counts per energy bin. In this case, equivalent widths of emission lines may be somewhat over-estimated.  It must be remembered, however, that the $\\chi^{2}_{q}$ confidence ranges can say nothing of the {\\it simultaneous} confidence ranges of the other $p-q$ parameters. For example, suppose one observes an active galaxy or X-ray binary and measures the magnitude of an X-ray reflection continuum component (due to Compton-thick scattering) and the equivalent width of an iron-K line and quotes, as is common practice, 90\\% confidence errors for one interesting parameter ($\\Delta \\chi^{2} = 2.7$). Now, the relation between the iron-K line equivalent width and the strength of the reflection continuum can be predicted from a theoretical model, so one can determine whether the measured values are consistent with such a model, within the errors. However, one-parameter errors (as used , for example, in Zdziarski \\etal 1996) are inappropriate since these confidence ranges are not simultaneous. One must use two-parameter errors in such a case.  \\vspace{2cm}  Much of this work was done during an extended stay at the Institute of Space and Astronautical Science (ISAS), Japan, during the summer of 1993 and two weeks in November 1996. The author thanks everyone in the X-ray astronomy group at ISAS for their great hospitality. The author would also like to thank Peter Serlemitsos, Richard Mushotzky, Andy Fabian, and Andy Ptak for useful discussions, and Keith Arnaud for generally maintaining XSPEC and fixing bugs promptly. The author is very grateful to Dr. W. Cash for correcting some serious errors in an earlier version of the paper and is also indebted to an anonymous referee for making some extremely important points which led to a complete revision of the paper.  \\newpage"
    },
    "9803/astro-ph9803327_arXiv.txt": {
        "abstract": "We present {\\it Faint Object Camera (FOC)} ultraviolet images of the central $14 \\times 14 \\arcsec$ of Messier 31 and Messier 32.  The hot stellar population detected in the composite UV spectra of these nearby galaxies is partially resolved into individual stars, and their individual colors and apparent magnitudes are measured.  We detect 433 stars in M~31 and 138 stars in M~32, down to detection limits of $m_{F275W}$~=~25.5~mag and $m_{F175W}$~=~24.5~mag. We investigate the luminosity functions of the sources, their spatial distribution, their color-magnitude diagrams, and their total integrated far-UV flux. Comparison to {\\it IUE} and {\\it HUT} spectro-photometry and {\\it WFPC2} stellar photometry indicates consistency at the 0.3~mag level, with possible systematic offsets in the {\\it FOC} photometry at a level less than this. Further calibrations or observations with the {\\it Space Telescope Imaging Spectrograph (STIS)} will be necessary to resolve the discrepancies.  Our interpretation rests on the assumption that the published {\\it FOC} on-orbit calibration is correct.  Although M~32 has a weaker UV upturn than M~31, the luminosity functions and color-magnitude diagrams of M~31 and M~32 are surprisingly similar, and are {\\it inconsistent} with a majority contribution from any of the following: PAGB stars more massive than 0.56~$M_\\odot$ (with or without associated planetary nebulae), main sequence stars, or blue stragglers.  Both the luminosity functions and color-magnitude diagrams are {\\it consistent} with a dominant population of stars that have evolved from the extreme horizontal branch (EHB) along tracks with masses between 0.47 and 0.53 $M_\\odot$.  These stars are well below the detection limits of our images while on the zero-age EHB, but become detectable while in the more luminous (but shorter) AGB-Manqu$\\acute{\\rm e}$ and post-early asymptotic giant branch (PEAGB) phases.  The {\\it FOC} observations require that only a very small fraction of the main sequence population (2\\% in M~31 and 0.5\\% in M~32) in these two galaxies evolve though the EHB and post-EHB phases, with the remainder evolving through bright PAGB evolution that is so rapid that few if any stars are expected in the small field of view covered by the {\\it FOC}. A model with a flat EHB star mass distribution reproduces the {\\it HUT} and {\\it IUE} spectra of these two galaxies reasonably well, although there is some indication that an additional population of very hot (T$_{\\rm eff} > 25000$ K) EHB stars may be needed to reproduce the {\\it HUT} spectrum of M~31 near the Lyman limit, and to bring integrated far-UV fluxes of M~31 and M~32 into agreement with {\\it IUE}.  In addition to the post-EHB population detected in the {\\it FOC}, we find a minority population ($\\sim$ 10\\%) of brighter stars that populate a region of the CMD that cannot be explained by canonical post-HB evolutionary tracks. The nature of these stars remains open to interpretation.  The spatial distributions of the resolved UV-bright stars in both galaxies are more centrally concentrated than the underlying diffuse emission, implying that stellar populations of different age and/or metallicity might be responsible for each component. ",
        "introduction": "The spectra of elliptical galaxies and spiral galaxy bulges exhibit a strong upturn shortward of 2700~\\.{A}, dubbed the ``UV upturn.'' At the time of its discovery, the existence of a hot stellar component went against the traditional picture of early-type galaxies. The canonical view of ellipticals held that these galaxies contained a cool, passively evolving population of old stars.  The pioneering UV observations of ellipticals -- with the {\\it Orbiting Astronomical Observatory (OAO)} (Code \\& Welch 1979\\markcite{CW79}) and the {\\it International Ultraviolet Explorer (IUE)} (Bertola, Capaccioli, \\& Oke 1982\\markcite{BCO82}) -- could only sample the Rayleigh-Jeans tail of the hot UV flux, with poor signal-to-noise and resolution. Early explanations for the source of the UV upturn covered a wide range of candidates, including massive young stars, hot horizontal branch stars, planetary nebula nuclei, and several binary scenarios (see Greggio \\& Renzini 1990\\markcite{GR90} for a complete review). The presence of young stars would imply ongoing star formation in early-type galaxies, while the evolved candidates suggested that old stellar populations could be efficient UV emitters.  Characterized by the $m_{1550}-V$ color, the UV upturn shows surprisingly strong variation (ranging from 2.05--4.50 mag) in nearby quiescent early-type galaxies (Bertola et al.\\ 1982\\markcite{BCO82}; Burstein et al.\\ 1988\\markcite{B88}), even though the spectra of ellipticals at longer wavelengths are qualitatively very similar.  A large sample of UV measurements demonstrated that the UV upturn is positively correlated with the strength of Mg$_2$ line absorption in the $V$ band, in the sense that the $m_{1550}-V$ color is bluer at higher line strengths, opposite to the behavior of optical color indices (Burstein et al.\\ 1988\\markcite{B88}). Opposing theories have been devised to explain this correlation. For example, Lee (1994\\markcite{L94}) and Park \\& Lee (1997\\markcite{PL97}) suggest that the UV flux originates in the low metallicity tail of an evolved stellar population with a wide metallicity distribution.  Their reasoning is that the more massive ellipticals formed earlier, and that it is actually the {\\it mean} metallicity that is higher (and driving the optical indices) in the older and bluer galaxies. In contrast, several groups (Brocato et al.\\ 1990\\markcite{BMMT90}; Bressan, Chiosi, \\& Fagotto 1994\\markcite{BCF94}; Greggio \\& Renzini 1990\\markcite{GR90}; Horch, Demarque, \\& Pinsonneault 1992\\markcite{HDP92}; Brown et al.\\ 1997\\markcite{B97}; Yi, Demarque, \\& Oemler 1997\\markcite{YDO97}) argue that metal-rich horizontal branch stars and their progeny are responsible for the UV flux.  Under the metal-rich hypothesis, the canonical trend in horizontal branch morphology (i.e., redder HBs with higher metallicity) is reversed at high metallicity, due to increased helium abundance and possibly a higher mass loss rate on the red giant branch, resulting in the production of hot UV-efficient stars on the extreme horizontal branch (EHB).  These hypotheses lead to different ages for the stellar populations in these galaxies.  Ages exceeding those of Galactic globular clusters are required under the low-metallicity Park \\& Lee (1997\\markcite{PL97}) hypothesis, while ages as low as 8 Gyr are allowed in the Bressan et al.\\ (1994\\markcite{BCF94}) model. In these two scenarios, the EHB stars are drawn from either tail of the metallicity distribution.  However, it is also possible, indeed perhaps more likely, that the EHB stars arise from progenitors near the peak of the metallicity distribution (cf.\\ Dorman, O'Connell, \\& Rood 1995\\markcite{DOR95}), but represent a relatively rare occurrence. The correlation of $m_{1550}-V$ with the global metallicity of the galaxy might indicate that this rare path of stellar evolution becomes less so at high metallicity and helium abundance.  The {\\it Hopkins Ultraviolet Telescope (HUT)}, designed for medium resolution ($\\approx$~3~\\.{A}) spectroscopic observations of faint extended UV sources down to the Lyman limit at 912~\\.{A}, offered a new perspective on these populations. With observations of two galaxies on the Astro-1 mission (M~31 and NGC~1399), Ferguson et al.\\ (1991\\markcite{F91}) demonstrated that young, massive stars cannot be a significant contributor to the UV upturn.  There is a lack of strong \\ion{C}{4} absorption expected from such stars, and the continuum flux decreases from 1050~\\.{A} down to the Lyman limit.  Such a decrease implies that the UV flux is dominated by stars with temperatures $\\leq 25000$~K and is incompatible with emission by a population of young stars having a normal initial mass function. Ferguson et al.\\ (1991\\markcite{F91}) also suggested that a bimodal distribution on the horizontal branch was needed to reproduce the shape of spectra from the near-UV to the far-UV, otherwise the spectra would be flatter than observed. Six more galaxies were observed on the Astro-2 mission (M~49, M~60, M~87, M~89, NGC~3115, and NGC~3379), and with these data Brown et al.\\ (1997\\markcite{B97}) demonstrated that a two-component population of high-metallicity post-HB stars could reproduce the UV light seen in nearby ellipticals.  In this model, most ($> 80$\\%) of the UV-producing stars were undergoing post-asymptotic-giant-branch (PAGB) evolution, and the remainder were evolving along AGB-Manqu$\\acute{\\rm e}$ paths from the extreme horizontal branch.  Although in the minority, these AGB-Manqu$\\acute{\\rm e}$ stars can produce the majority of the flux, because their lifetimes are orders of magnitude longer than those of the PAGB stars.  Because even the brightest of these galaxies are faint and extended in the UV, studies of the UV upturn have focused mostly on the composite spectral energy distributions of ellipticals. However, the UV imaging capabilities of {\\it HST} have now opened the possibility of studying the resolved UV population, at least in the nearest galaxies. Attempts to do this prior to the {\\it HST} refurbishment were undertaken by King et al.\\ (1992\\markcite{K92}), for M~31, and by Bertola et al.\\ (1995\\markcite{BBB95}), for M~31 and M~32.  King et al.\\ (1992\\markcite{K92}) obtained a pre-COSTAR {\\it Faint Object Camera (FOC)} observation of a $44 \\times 44 \\arcsec$ field in the center of M~31, using the F175W filter and the F/48 relay. They found more than 100 sources that they identified as PAGB stars.  Intermediate-mass ($M > 0.6~M_{\\odot}$) PAGB stars are short-lived, and given a population size constrained by the fuel consumption theorem, the large number of detected stars implied that these were low-mass PAGB stars.  King et al.\\ (1992\\markcite{K92}) estimated that these stars accounted for approximately 20\\% of the UV light, with the rest unresolved, presumably coming from EHB stars and their AGB-Manqu$\\acute{\\rm e}$ descendants, which could account for this unresolved light without violating fuel consumption constraints.  Bertola et al.\\ (1995\\markcite{BBB95}) used the {\\it FOC} to image M~31, M~32, and NGC~205 with the combination of the F150W and F130LP filters on the F/48 relay. Although the optical luminosity enclosed by the {\\it FOC} field was higher in M~32 than in M~31, they found far fewer UV sources in the M~32 field.  Because M~32 has a weaker UV upturn and lower metallicity, the UV light was expected to originate in PAGB stars of higher mass (and shorter lifetimes) than those in M~31; so, the relative numbers of detected sources were in line with these expectations.  However, the luminosity functions in M~31 and M~32 appeared similar, in contrast to expectations when comparing a population of less massive PAGB stars to a population of more massive ones.  Such a puzzle may be partly explained if the PAGB stars are enshrouded in dust during the early part of their evolution away from the AGB.  Both previous {\\it FOC} studies faced daunting challenges in untangling the uncertainties in the {\\it FOC} sensitivity calibration and the red leak of the filters.  King et al.\\ (1992\\markcite{K92}) adopted the best in-flight calibration at the time, and used ground calibrations of the filter and photocathode response to assess the effects of red leak. They also determined that the pre-COSTAR PSF required a huge aperture correction of 2.6~mag. Bertola et al.\\  (1995\\markcite{BBB95}) derived an independent calibration based upon {\\it IUE} observations of NGC~205, M~31, and M~32.  In the Bertola et al.\\ calibration, the nominal {\\it FOC} F150W+F130LP efficiency curve required multiplication by factors of 0.21--0.92 (varying with wavelength) in order to produce agreement between the {\\it FOC} and {\\it IUE}. Checking their revised calibration against common stars in the King et al.\\ (1992\\markcite{K92}) F175W images, Bertola et al.\\ determined that the F175W efficiency curve also required revision, such that the peak in the efficiency curve was at 28\\% of its nominal value, but with increased red leak from longer wavelengths.  We demonstrate in \\S\\ref{secf48} that the pre-COSTAR {\\it FOC} data calibration was seriously in error.  We have used the refurbished, recalibrated {\\it FOC} to follow up these earlier studies with deeper UV images of the M~31 and M~32 cores, in order to further characterize the evolved stellar populations in these galaxies.  Our observations use the F175W and F275W filters to determine color-luminosity relationships in these galaxies, and to compare them to the predictions of stellar evolutionary theories. ",
        "conclusions": "\\label{secdis}  Since the discovery of the UV upturn phenomenon, a great deal of research has tried to characterize the stellar population producing the UV light in spiral bulges and ellipticals.  The range of candidates has included young massive stars, hot horizontal branch stars, planetary nebula nuclei (PAGB stars), and several binary scenarios (see Greggio \\& Renzini 1990\\markcite{GR90} for a complete review). Earlier work relied primarily upon {\\it IUE} and {\\it HUT} spectra of composite stellar populations, and pre-COSTAR {\\it FOC} observations (King et al.\\ 1992\\markcite{K92}; Bertola et al.\\ 1995\\markcite{BBB95}) of the brightest stars. Our current understanding of the populations in ellipticals and spiral bulges predicts that a significant population of hot HB and post-HB stars should be evolving from the extreme horizontal branch and following either AGB-Manqu$\\acute{\\rm e}$ or post-early-AGB evolution (see Brown et al.\\ 1997\\markcite{B97} and references therein). The prediction is mainly based upon fuel consumption constraints that rule out PAGB stars as the sole UV producers; the PAGB stars are just too short-lived to efficiently produce the required UV light.  Our {\\it FOC} observations confirm the existence of hot post-EHB stars in the centers of M~31 and M~32, and also confirm that PAGB stars are not the predominant component of the UV-bright population.  The existence of these EHB stars also implies that the horizontal branches in both M~31 and M~32 are at least somewhat extended, because ``red clump'' horizontal branches are unlikely to produce the stars seen in our color-magnitude diagrams. Future observations, given the improved UV capabilities of HST with the installation of {\\it STIS}, should be able to reach the HB and thus characterize the entire evolved population more directly.  We note that significant systematic uncertainties may exist in our {\\it FOC} photometry, at the 0.3 mag level.  Our results should be considered with these uncertainties in mind; however, the uncertainties are not large enough to allow a completely different interpretation of the detected stellar population.  As discussed in \\S\\ref{secehbcmd}, we find a minority population ($\\sim$ 10\\%) of brighter stars that cannot be explained by canonical post-HB evolutionary tracks. Although PEAGB or very low-mass PAGB stars evolve through this region of the CMD, their evolution during this phase is too rapid to produce the 35 stars seen in M~31 while maintaining consistency with fuel consumption constraints and observed integrated spectra of M~31. The nature of these stars remains unexplained as of this writing.  The fraction of light in the resolved UV population in the center of M~31 and in the center of M~32 is consistent with expectations from the {\\it HUT} and {\\it IUE} spectra of these galaxies.  The far-UV light in M~31 can be explained by a main sequence population where 98\\% of the stars channel through intermediate-mass PAGB evolution and 2\\% through EHB evolution.  The far-UV light in M~32 can be explained by a population where 99.5\\% of the stars channel through intermediate-mass PAGB evolution and 0.5\\% through EHB evolution. Our simple model populations can account for most of the far-UV light, although there is room in both galaxies for a contribution from very low-mass hot EHB stars, which would be below our detection limits even in the AGB-Manqu$\\acute{\\rm e}$ phase.  We found that the stellar populations in M~31 and M~32 do not appear remarkably different in our {\\it FOC} UV images. Fewer stars appear to be entering the UV-efficient EHB paths in M~32, and this explains both the smaller absolute number of detected stars and the smaller fraction of resolved UV light. The similarities of the luminosity functions of M~31 and M~32 were unexpected, given the dramatically different $m_{1550}-V$ colors of the two galaxies. This finding suggests that, while the fractional mass in EHB stars is probably sensitive to the properties of the overall stellar population, (e.g., metallicity, age, or helium abundance), the mass distribution on the EHB may not be as sensitive to these parameters.  Certainly, the differences can not be investigated without deeper images that can resolve more of the evolved population.  {\\it STIS} observations planned for late 1998 will reach the HB in M~32 and provide direct evidence for the horizontal branch distribution."
    },
    "9803/astro-ph9803111_arXiv.txt": {
        "abstract": "We have obtained high-resolution echelle spectra of 18 solar-type stars that an earlier survey showed to have very high levels of \\ion{Ca}{2} H and K emission.  Most of these stars belong to close binary systems, but 5 remain as probable single stars or well-separated binaries that are younger than the Pleiades on the basis of their lithium abundances and \\ha\\ emission.  Three of these probable single stars also lie more than 1 magnitude above the main sequence in a color-magnitude diagram, and appear to have ages of 10 to 15 Myr.  Two of them, HD 202917 and HD 222259, also appear to have a kinematical association with the pre-main sequence multiple system HD 98800. ",
        "introduction": "Two years ago we published a survey of \\ion{Ca}{2} H and K emission strengths in more than 800 southern solar-type stars (\\cite{h96} 1996; hereafter Paper I), determined from low-resolution ($R\\approx2,000$) spectra obtained at CTIO.  The purpose of the survey was to provide at least rough estimates of the ages of the individual stars, and to examine the distribution of emission strengths in a large and unbiased sample.  The very existence of a simple or single relationship between chromospheric emission (CE) and age is debatable, but there is ample evidence that CE declines steadily with age (\\cite{sk72} 1972; \\cite{sdj} 1991).  In other words, we are confident of a general decline of CE with age because of observations of stars in clusters, and also because CE is so intimately tied to stellar rotation and we know that rotation declines with age in solar-type stars.  But we also know that the CE levels of individual stars vary due to rotational modulation of active regions, long-term cycles, and other phenomena, and we also know that not all stars reach the Zero-Age Main Sequence (ZAMS) with the same angular momentum, and thus we are not so sure that there is a unique CE-age relation that applies to all stars.  The survey of Paper I included stars from a G-dwarf sample defined using the combination of two-dimensional spectral types (\\cite{hk1} 1975; \\cite{hk2} 1978; \\cite{hk3} 1982; \\cite{hk4} 1988) and Str\\\"omgren photometry (\\cite{o88} 1988, \\cite{o93} 1993), and it turned up two groups of stars that we have studied further.  The first group we called ``Very Active'' because they exhibited CE levels well above any seen in the field stars of the earlier survey of \\cite{vp80} (1980).  The second group is called ``Very Inactive'' and consists of stars that appear to have activity levels well below the Sun's.  The Very Inactive stars will be the subject of a future paper.  Here we concentrate on the Very Active stars, and we are led to examine them in detail for several reasons.  First, stars with very high levels of activity are that way because they rotate rapidly, and they rotate rapidly either because they are very young (and have not yet lost much of their initial angular momentum), or because they are in a close binary system (where the companion's tidal forces lead to synchronous rotation).  Both types of systems are interesting, and both types offer laboratories for the study of the rotation-activity relation.  A second reason for undertaking this detailed study is to find very young stars in the immediate solar neighborhood (i.e., within about 50 pc).  Even if they were evenly mixed in the Galaxy, very young stars would be rare just because their ages ($\\la100$ Myr) represent such a small fraction of the age of the Galactic disk.  But such stars are {\\it not} evenly mixed because they form in discrete regions and take time to be dispersed into the field.  There may be a few stars near the Sun that are as young as, say, the Pleiades, but they are few indeed.  However, even small numbers matter since they imply that many more such stars lie in the much vaster volume of the greater solar neighborhood (i.e., within $\\sim100$ pc).  Some of these very young field stars, such as the HD 98800 system, may be examples of the elusive post-T Tauri star class; see, e.g., \\cite{s98} (1998).  The moderate-resolution spectra obtained for Paper I were centered on the \\ion{Ca}{2} H and K lines to determine the level of CE.  We have obtained higher-resolution echelle spectra to confirm their high activity levels (at \\ha), to test for youth (with Li), and to observe each star several times to search for radial velocity changes indicative of close companions. ",
        "conclusions": "We present high resolution echelle spectroscopy of 18 Very Active southern solar-type stars identified in the \\ion{Ca}{2} H and K survey of \\cite{h96} (1996).  We find evidence from line doubling or radial velocity variations that 13 of these are members of short-period, close binary systems.  Activity in these stars may thus be due to the presence of a close companion, rather than youth; however, it is entirely possible that some of the objects (HD 106506, 119022, 155555AB, and HD 222259B) are young close binaries.  Just a few exposures of high resolution but very modest S/N seems to be a highly effective and efficient means of identifying active close binary systems.  Based on our \\ha\\ observations, we confirm that the remaining five of the eighteen stars are also Very Active, and find no evidence that such activity is due to membership in a close binary system.  Four of these stars also have significant Li abundances, comparable to or larger than Pleiades stars of similar color.  We consider these five stars to be young candidates, i.e., stars whose chromospheric activity seems associated solely with youth.  We note that two of the young candidates (HD 202917 and HD 222259) appear to have the same $U,V$ velocities as HD 98800, a rare example of a post-T Tauri star in the field, according to \\cite{s98} (1998).  While the formation site of these stars is unclear, the three stellar systems may be part of a small group of very young stars sharing a common history and origin.  The $UVW$ velocities of HD 175897 are in excellent agreement with those of the Pleiades.  Two other interesting objects are HD 106506 and HD 119022, which were noted above to be possible young objects excluded from our final best young candidate list due to their duplicity.  Both of these objects demonstrate sharp features in the broader Na D lines and stronger features of other metals.  We do not observe any radial velocity shifts of these sharp features.  This would suggest an interstellar or circumstellar origin for the sharp D-line features. These two objects are among the most distant in our sample: the Hipparcos-based distances are ${\\sim}125$ pc, but seems too near for an interstellar origin.  There is thus the intriguing possibility that these features arise from circumstellar material, perhaps not unlike that surrounding ${\\beta}$ Pictoris or early-type shell stars similarly inferred from the presence of sharp features in a broader stellar absorption line. Comparison of the spectral type and photometry for HD 119022 suggests a moderate-sized reddening of perhaps 0.1 to 0.15 mag in $E(B-V)$.  Higher quality spectra and high resolution imaging of these objects would be of great interest.  Finally, Li abundances or upper limits were derived for the sharp-lined inactive \\ha\\ standards employed in our program.  We report detectable Li for the solar \\teff\\ star HD 76151 and for the late F-stars HD 38393 and 45067. There may be a Li abundance difference in the two similar wide components of HD 158614; better knowledge of their fundamental parameters from spatially resolved photometry and spectroscopy is needed."
    },
    "9803/astro-ph9803039_arXiv.txt": {
        "abstract": "Multi-epoch VLBA observations of the maser in NGC 4258 have yielded a 4\\% geometrical distance to the galaxy.  The potential scientific payoffs of finding similar objects at large distances, in the Hubble flow, are considerable.  In this contribution, I discuss the plausibility of detecting high-redshift water masers, and describe a search strategy that we have implemented to realize this objective. ",
        "introduction": "VLBA observations of the water masers in NGC 4258 reveal a nearly edge-on, slightly warped, extremely thin disk in nearly perfect Keplerian rotation around a central binding mass of $3.5\\times10^{7}$\\,M$_{\\odot}$ (Watson \\& Wallin 1994; Greenhill {\\it et al.} 1995; Miyoshi {\\it et al.} 1995; Moran {\\it et al.} 1995; Herrnstein, Greenhill, \\& Moran 1996). The VLBA observations of the NGC 4258 maser provide insight into the structure and kinematics of the accretion disk, and NGC 4258 is an exceptional laboratory for the study of AGN accretion phenomenon and the connection between accretion disks and jet emission (Herrnstein {\\it et al.} 1997; Herrnstein {\\it et al.} 1998a). They can also be used to derive a precise geometric distance to NGC 4258.  The upper panel of Figure~1 is a schematic representation of the best-fitting NGC~4258 disk model, as derived from the positions and line-of-sight (LOS) velocities of the masers.  As the disk rotates, the 'systemic' masers along the near edge of the disk drift in position and LOS velocity by about 30\\,$\\mu$as\\,yr$^{-1}$ and 9\\,km\\,s$^{-1}$, respectively, with respect to the apparently stationary 'high-velocity' masers in the plane of the sky (Herrnstein 1997a\\&b). The expressions at the bottom of Figure~1 illustrate that {\\it both} the LOS accelerations ($\\dot{v}_{LOS}$) {\\it and} the proper motions ($\\dot{\\theta}_{x}$) can be used to derive a purely geometric distance ($D$) to NGC 4258.  Because the actual space velocities ($v_{rot}$) and angular radii ($\\theta_{R}$) of the systemic masers cannot be measured directly, these acceleration and proper motion distances are model dependent. Fortunately, however, most of the model dependence resides in $\\theta_{R}$, and the two distance estimates together provide a geometric distance estimate that is largely model independent. With this in mind, we have observed NGC 4258 with the VLBA at 3--4 month intervals for the last three years.  The LOS accelerations and proper motions provided by the first five epochs of these data yield a geometric distance of $7.3\\pm0.3$\\,Mpc (Herrnstein 1997a; Herrnstein {\\it et al.} 1998b).  \\begin{figure}[p]     % \\plotone{herrnf1.eps} \\caption {The top panel is a schematic representation of the sub-parsec molecular disk in NGC 4258.  The bottom panels show the expected acceleration and proper motion vectors superposed on actual VLBA data of the systemic masers taken in May of 1995.} \\end{figure} ",
        "conclusions": ""
    },
    "9803/astro-ph9803333_arXiv.txt": {
        "abstract": "We present the results of adaptive-optics imaging of the $z=0.2923$ QSO PG\\,1700+518 in the $J$ and $H$ bands.  The extension to the north of the QSO is clearly seen to be a discrete companion with a well-defined tidal tail, rather than a feature associated with the host galaxy of PG\\,1700+518 itself.  On the other hand, an extension to the southwest of the QSO (seen best in deeper, but lower-resolution, optical images) does likely comprise tidal material from the host galaxy. The SED derived from images in $J$, $H$, and two non-standard optical bands indicates the presence of dust intermixed with the stellar component. We use our previously reported Keck spectrum of the companion, the SED found from the imaging data, and updated spectral-synthesis models to constrain the stellar populations in the companion and to redetermine the age of the starburst. While our best-fit age of 0.085 Gyr is nearly the same as our earlier determination, the fit of the new models is considerably better. This age is found to be remarkably robust with respect to different assumptions about the nature of the older stellar component and the effects of dust. ",
        "introduction": "PG\\,1700+518 ($z=0.2923$), one of the more luminous low-redshift QSOs, shows a bright, arc-like structure extending about 2\\arcsec\\ to the north of the QSO (Hutchings, Neff, \\& Gower \\markcite{hut92b}1992; Stickel \\etal\\ \\markcite{sti95}1995).  We recently presented a spectrum of this extension, showing that it is dominated in the optical by a stellar population $\\sim10^8$ years old (Canalizo \\& Stockton \\markcite{can97}1997; hereinafter Paper 1).  We argued that it is plausible that this starburst and the QSO activity were both triggered by a recent interaction, and that spectroscopic dating of starbursts in QSO hosts and strongly interacting companions could lead to the development of an empirical evolutionary sequence for QSOs.  From the imaging data we had on PG\\,1700+518 at that time, we were unsure whether the extension was a companion galaxy or a feature associated with the host galaxy of the QSO.  Here we describe the results of adaptive-optics (AO) imaging and show that the object is indeed a discrete companion galaxy undergoing strong interaction with the QSO host.  We combine these data with previous imaging in two optical bandpasses to give a spectral-energy distribution (SED) from 0.42 \\micron\\ to 1.28 \\micron\\ in the rest frame in order to place additional constraints on the stellar populations.  We then use these constraints and fits of new spectral synthesis models to the Keck LRIS spectrum to obtain a more reliable age for the post-starburst component. ",
        "conclusions": "\\subsection{The Morphology of the Companion Galaxy} The AO images of PG\\,1700+518 are shown in Fig.\\ 1, in original, PSF-subtracted, and {\\it plucy}-restored versions. The extension to the north of the QSO, which has the appearance of an arc-like or ``boomerang'' shape in the best previous ground-based images (Hutchings \\& Neff \\markcite{hut92}1992; Stickel \\etal\\ \\markcite{sti95}1995; Paper 1), retains that general appearance in our higher-resolution images, but it now also shows discrete condensations within this overall structure.  It is also clearly a separate companion galaxy:  there is a rather abrupt dropoff in surface brightness just south of the brightest condensation, whereas we would expect to see more continuity with the inner regions if this were a tail associated with the host galaxy.  Such a connection should stand out in spite of our using the QSO to model the PSF, since we expect to be sensitive to any non-elliptically-symmetric features.  The companion, though distinct, is apparently in the process of merging with the QSO host galaxy.  The main features of the companion are consistent between the $J$ and $H$ images and must be real.  The bright condensation $a$ is probably the nucleus of the companion.  The apparent tidal tail, curving to the north and east, contains another condensation, $b$, which may be a bright star-forming region or even a dwarf galaxy forming from the tidal debris (\\eg\\ Duc \\& Mirabel \\markcite{duc94}1994; Hunsberger, Charlton, \\& Zaritsky \\markcite{hun96}1996).  The nature of features at lower surface brightnesses is less certain because of the effect of noise on the deconvolution.  There is clearly material to the east of $a$, looking like another condensation in the $J$ image, but more like an arc in the $H$ image.  The peak to the south of $a$ on the $J$ image is apparently an artifact of inadequacies in the PSF model, but there is some evidence in these images for bridge-like material between the companion and the QSO.  If most of the luminous material north of the QSO is associated with the companion, is there any evidence for tidal debris from the QSO host?  Stickel \\etal\\ \\markcite{sti95}1995 noted a possible faint extension to the SW of the QSO, which can also be seen in Fig.\\ 2 of Paper 1.  Here we show this feature in two optical bandpasses in Fig.\\ 2, where we have slightly oversubtracted the wings of the PSF profile to show the SW extension more clearly.  We suggest that this is likely a counter-tidal feature from the QSO host.  The inner, higher-surface-brightness contours of the host galaxy appear to be aligned nearly E--W (Paper 1).  As we were completing this {\\it Letter,} we were informed that Hines \\etal\\ \\markcite{hin98}1998 had obtained a Hubble Space Telescope NICMOS image of PG\\,1700+518.  They also conclude that the extension to the N is a companion, but they see what we have described as a tidal tail rather as part of a ring.  \\subsection{The SED of the Companion and the Age of the Starburst}  In Paper 1, we presented a spectrum of the companion, corrected for contamination from the QSO. We modeled the SED as a superposition of two simple stellar populations (instantaneous bursts). Our age estimates, based on a $\\chi^2$ fit of the Bruzual \\& Charlot \\markcite{bru93}(1993) models to the spectrum of the companion, gave 0.09 (+0.04, $-0.03$) Gyr and 12.25 Gyr, respectively, for the two components.  Two recent developments encourage us to try to refine these estimates: better models are now available (Bruzual \\& Charlot \\markcite{bru98}1998), and we now have images in bandpasses covering a wide spectral range, which can additionally constrain the stellar populations. We use the same Keck Low-Resolution Imaging Spectrograph (LRIS) data from Paper 1. Because we need fairly high resolution for detailed fitting to features in the spectrum, we restrict ourselves to the solar-metallicity spectral-synthesis models based on the Gunn \\& Stryker \\markcite{gun83}(1983) and Jacoby, Hunter, \\& Christian \\markcite{jac84}(1984) spectra.  We first discuss fits to the Keck LRIS spectrophotometry, considered in isolation; then we show how the models must be modified to take into account the SED over a wider wavelength region. Although the fit to the spectrum shown in Fig.\\ 3 is from our final model, it is typical of the quality of fit we find for other models we discuss here.  If we try a similar approach to that of Paper 1, using only simple stellar populations from the newer models (Bruzual \\& Charlot \\markcite{bru98}1998), we obtain nearly the same age for the younger population (0.10 Gyr), but a much younger age for the older population (1.8 Gyr).  While this new model fits the Keck spectrum much better than did the model given in Paper 1, and we could no doubt improve the fit even more by adding a third, older population, the morphological evidence suggests another approach.  The presence of a strong tidal tail indicates a dominant, pre-existing disk component, and the evidence for recent star formation indicates a gas-rich system.  Such a galaxy would be expected to have been forming stars over the entire lifetime of the disk.  We therefore consider another range of models:  those in which there has been an exponentially-decaying rate of star formation on which is superposed a single recent starburst. Specifically, we assume that the younger population can be modelled as a burst (there is no appreciable difference in this case whether the burst is taken to be instantaneous or has a finite duration of several Myr) and that the older population has an exponentially-decaying star formation beginning 10 Gyr ago.  We find that the age of the younger population and the quality of the fit are insensitive to the decay rate of star formation in the older component over a reasonable range (time constants $\\sim3$--10 Gyr), but that we cannot get as good fits to the observed spectrum for either an old instantaneous burst or a constant star-formation rate.  For definiteness, we assume an exponential time constant of 5 Gyr.  Considering only the Keck spectrophotometry, we find the best $\\chi^2$ fit by combining this underlying population with a starburst with an age of 114 Myr.  We now deal with additional constraints from the SED at longer wavelengths.  We determine the SED of the companion over a rest-frame range from 0.42 \\micron\\ to 1.28 \\micron\\ from the AO $J$ and $H$-band images and images in two non-standard optical bands centered at 5442 and 7248 \\AA, with FWHM of 1002 and 1260 \\AA, respectively (see Fig.\\ 2; these bandpasses have been designed to avoid strong emission lines at the redshift of PG\\,1700+518). Figure 4 shows the flux densities in these four filters for a 1\\arcsec\\ diameter aperture centered 0\\farcs4 E and 2\\farcs3 N of the QSO.  This region is far enough from the QSO that the photometry should not be very sensitive to the chosen scaling or other errors in the PSF subtraction process.  While it does not exactly coincide with the region covered by the Keck LRIS spectrum, and there could be small differences in the two SEDs, we attempt to find a model that will fit both. We have explored a wide variety of SEDs involving superpositions of both simple and composite stellar populations based on the Bruzual \\& Charlot \\markcite{bru98}(1998) models; none of these gives a satisfactory fit both to the Keck LRIS spectrum and to the wide-band SED. A pure old population comes moderately close to matching the overall SED, but it fails completely to reproduce either the strong Balmer absorption spectrum or the continuum shape of the restframe UV---blue spectrum.  On the other hand, the ``reddest'' reasonable model we can find that adequately fits the Keck spectrum and the optical photometry falls well below the $J$ and $H$ points (see Fig.\\ 4). Even this latter model is rather unphysical, comprising a pure young population (to fit the UV---blue spectrum) and a pure old stellar population (to attempt to fit the IR photometry), with nothing in between; however, the addition of any intermediate-age component only makes the fit worse.  Clearly one plausible way to raise the relative flux at longer wavelengths while retaining an early-type spectrum at shorter wavelengths is to include some sort of reddening due to dust.  However, simply applying a Galactic reddening law (\\ie\\ assuming a dust screen between the object and the observer) is neither appropriate nor very effective.  Screen-like reddening sufficient to force a fit to the overall SED results in very strong variation in extinction across the range of the Keck spectral fit, making it difficult to fit simultaneously the lines and continuum. One needs a means of effecting a substantial reddening of the IR flux with respect to the optical flux without changing the slope in the UV---blue region too much.  Witt, Thronson, \\& Capuano \\markcite{wit92}(1992) have shown that dust distributions that are more-or-less coextensive with the stellar distribution, in addition to being more realistic, can provide exactly this sort of reddening.  Such models differ from the standard reddening law (due to intervening dust) in two main respects:  (1) some of the blue light removed along the sightline is compensated by scattering of blue light emitted in other directions into the line of sight, and (2) optical depth effects ensure that most of the light appearing at short wavelengths has suffered little extinction, while, at longer wavelengths, a larger portion of the total stellar distribution contributes to the emergent flux.  We have used the family of ``dusty galaxy'' models calculated by Witt \\etal\\ \\markcite{wit92}(1992) to produce sets of modified Bruzual \\& Charlot \\markcite{bru98}(1998) models. The dust and the stars are both assumed to have constant density within a sphere, and the only variable is the optical depth to the center at a specific wavelength. While these models are highly artificial, they are sufficient to show the general nature of the reddening due to embedded dust, and they give us a good fit to both the LRIS spectrophotometry and the wide-band SED.  We have included in Fig.\\ 4 the model we have found to give the best fit to both the Keck spectral data and the overall SED. This same model is shown in Fig.\\ 3, both as fit to the Keck spectrum and as individual components:  the starburst, represented by an 85-Myr-old instantaneous burst, and the underlying population having star formation beginning 10 Gyr ago and exponentially decaying with a time constant of 5 Gyr. Both components are reddened by the curve given by the Witt \\etal\\ \\markcite{wit92}(1992) ``dusty galaxy'' model with a $V$-band central optical depth of 6.  The fit of the model to the Keck LRIS spectrophotometry (Fig.\\ 3) is remarkably good, reproducing most individual features as well as the continuum slopes quite accurately.  We can account for most of the deviations.  Poor fits to H$\\beta$ (and, to a lesser extent, H$\\gamma$) absorption are due to distortion of these profiles by general emission around the QSO not specifically associated with the companion:  there is both a positive component and a negative component (at a higher velocity, from the region on the opposite side of the QSO that was used to subtract off scattered QSO light; see Paper 1 for details).  Similarly, the excess peak near 3868 \\AA\\ is due to [\\ion{Ne}{3}] emission.  The broad dip between 4500 and 4600 \\AA\\ is due to excess subtraction of \\ion{Fe}{2} emission, which is apparently spatially variable (Paper 1).  We have found that we can improve the fit of the \\ion{Ca}{2} $K$ line by including a secondary burst $\\sim2$ Gyr ago, which may be evidence for a previous close passage of the companion, but this evidence alone is slender enough that we have chosen not to include this added complexity to the model.  The only significant discrepancy that we cannot explain is the poor fit to H10 $\\lambda3798$; this may simply be due to a glitch in our observed spectrum.  \\subsection{Towards an Evolutionary Sequence for QSOs}  Although the SED model we have presented is the best fit we have found to the data, subject to the constraints of simplicity and astrophysical reasonableness that we have chosen, we can find other models that fit nearly as well.  However, in exploring various models, we have found that, for reasonably plausible combinations of young and old stellar populations and reddening curves (\\ie\\ those that come close to fitting both the LRIS spectrophotometry and the overall SED), there is very little spread in the age of the young population.  We consistently obtain ages in the range from 75 to 100 Myr, and we believe that this is a reasonable estimate of the uncertainty in our determination, subject only to any remaining uncertainty in the Bruzual \\& Charlot \\markcite{bru98}(1998) models themselves.  We regard this robustness in the age determination as a hopeful sign for our ongoing efforts to develop an age sequence for triggering events for QSOs lying in the transition region between ultraluminous IR galaxies and the classical QSO population in the far-IR two-color diagram (Paper 1; Canalizo \\& Stockton \\markcite{can98}1998; Stockton \\markcite{sto98b}1998).  This enthusiasm is only slightly dampened by the fact that the AO imaging shows clearly that our age determination is for the interacting companion to PG\\,1700+518, rather than for the host galaxy itself.  Obtaining similar quality spectrophotometry for the extension of the host galaxy to the SW, which is both considerably fainter and closer to the QSO, would be extremely difficult.  Nevertheless, its colors appear to be quite similar to those of the companion, and models of starbursts in interacting pairs (\\eg\\ Mihos \\& Herquist \\markcite{mih96}1996) indicate that star formation peaks at times of closest passage in both participants.  The projected $\\sim2$\\arcsec\\ ($\\approx7$ kpc) separation between the QSO and companion and the projected $\\sim5$ kpc tail length of the companion are entirely consistent with a close passage 85 Myr ago, if projection factors are $\\sim0.5$ and mutual velocities are $\\sim150$ km s$^{-1}$.  We therefore believe that the age of starburst in a close, clearly interacting companion is a good surrogate for an age determined from the QSO host galaxy itself for purposes of attempting to define an evolutionary sequence."
    },
    "9803/astro-ph9803043_arXiv.txt": {
        "abstract": "A preliminary 1.4 GHz RLF at redshift of about 0.14 is derived from the {\\it Las Campanas Redshift Survey} (LCRS) and the NVSS radio data. No significant evolution has been found at this redshift in comparison to the 'local' RLF. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803275_arXiv.txt": {
        "abstract": "We consider the possibility that cosmic magnetic field, instead of being uniformly distributed, is strongly correlated with the large scale structure of the universe. Then, the observed rotational measure of extra-galactic radio sources would be caused mostly by the clumpy magnetic field in cosmological filaments/sheets rather than by a uniform magnetic field, which was often assumed in previous studies. As a model for the inhomogeneity of the cosmological magnetic field, we adopt a cosmological hydrodynamic simulation, where the field is passively included, and can approximately represent the real field distribution with an arbitrary normalization for the field strength. Then, we derive an upper limit of the magnetic field strength by comparing the observed limit of rotational measure with the rotational measure expected from the magnetic field geometry in the simulated model universe. The resulting upper limit to the magnetic field in filaments and sheets is ${\\bar B}_{fs} \\la 1 \\mu{\\rm G}$ which is $\\sim10^3$ times higher than the previously quoted values. This value is close to, but larger than, the equipartition magnetic field strength in filaments and sheets. The amplification mechanism of the magnetic field to the above strength is uncertain. The implications of such a strength of the cosmic magnetic field are discussed. ",
        "introduction": "The strength and morphology of the intergalactic magnetic fields remain largely unknown, because it is intrinsically difficult to observe them (for recent reviews, see Kronberg 1994; Beck et al. 1996; Zweibel \\& Heiles 1997). While the magnetic fields inside typical galaxies are observed to be of order of $3-10\\mu{\\rm G}$, recent observations indicate that the magnetic fields of the similar strength are also common in core regions of rich clusters (for a recent discussion see En{\\ss}lin \\etal 1997). On the other hand, the upper limit for a large-scale field strength placed by the observed rotation measure (RM) of quasars is about $B_{IGM} \\la 10^{-9} L_{\\rm rev,Mpc}^{-1/2}{\\rm G}$ (Kronberg 1994). Here, $L_{\\rm rev,Mpc}$ is the reversal scale of the magnetic fields in units of Mpc in comoving coordinates. It was assumed the fields are uniform in direction within the reversal scale and varied as a comoving passive field in the expansion of the universe. The above value is close to the limit placed by the observed anisotropy in the cosmic microwave background radiation is $B_{IGM} < 6.8 \\times 10^{-9} (\\Omega_o h^2)^{1/2} {\\rm G}$ (Barrow, Ferreira \\& Silk 1997). Here, a uniform magnetic field was assumed. $h$ is $H_o$ in unit of $100~{\\rm km/s/Mpc}$.  According to a recent popular view based on both observational and theoretical cosmology, the dominant nonlinear structure in the universe is a web-like network of filaments (\\eg Bond, Kofman \\& Pogosyan 1996; Shectman \\etal 1996). At a lower density contrast, however, bubbly walls whose intersections are in fact filaments are the dominant structures. Such large scale structures form via gravitational instability, and then turbulent motions inside the structures as well as streaming motions along the structures necessarily exist (\\eg Kang \\etal 1994). Hence, magnetic fields in the early universe, if existed, should have been modified and amplified by those flow motions (\\eg Kulsrud \\& Anderson 1992; Kulsrud \\etal 1997). As a result, we expect that {\\it the ``geometry'' of cosmic magnetic field, which means the spatial distribution of the field strength and its orientation, should be correlated with the large-scale nonlinear structures of the universe}. In other words, the field strength increases with the matter density and its orientation tends to align with the sheets and filaments, while its random component is associated with the local turbulent motions. The magnetic fields should be strong along the walls of cosmic bubbles, even stronger along the filamentary superclusters, and the strongest inside the clusters of the galaxies, while they are very weak inside the voids.  Based on the proposition that the large scale magnetic field is correlated with the large scale structure of the universe, we consider a way to estimate the field strength which relies on the observational data.  The cosmic magnetic field together with free electrons in the intergalactic medium (IGM) induces the Faraday rotation in polarized radio waves from extra-galactic sources. The observational RM data of quasars show a systematic growth of RM with redshift, $z$, and limit RM to $\\sim 5~{\\rm rad~m^{-2}}$ or less at $z=2.5$ (Kronberg \\& Simard-Normandin 1976; Kronberg 1994 and references therein). Adopting a model for the distribution of the large scale magnetic field, these data can be used to constrain the strength of the magnetic field. For example, the upper limit estimated by Kronberg (1994) was based on a model in which the field is uniform within the reversal scale, the orientation is random, and there are $N=l_s/L_{rev}$ reversals along the line of sight to a source $l_s$ Mpc away from us.  In the present paper, we re-derive the upper limit by taking a new model in which the intergalactic magnetic field is mostly confined within filaments and sheets but very weak inside voids, rather than the simple model of uniform field. Of course the field would be strongest inside clusters, but their contribution is usually excluded in the observed RM. If we take only the geometric consideration that the field distribution along the lines of sight to radio sources has a small filling factor, high only inside filaments and sheets but low inside voids, then it is obvious that the expected magnetic field strength in filaments and sheets should be higher than the value derived from the uniform field model. But this geometric model is too simple, since the electron distribution as well as the field direction are uncertain. Here, we take a more practical approach by adopting numerical data in a simulated universe which includes a magnetic field, which is evolved during the large scale structure formation. With the magnetic field distribution, specific to this simulation model, we argue that the magnetic field strength in filaments and sheets is limited to ${\\bar B}_{fs} \\la 1~\\mu{\\rm G}$. We should emphasize that {\\it this limit should be approximately valid regardless of details of any model for the origin of the large scale magnetic field, provided that it is correlated with the large scale structure of the universe}.  In the next section, we describe the model simulation and the resulting magnetic field geometry. In \\S3, we explain the procedure to calculate the upper limit of the magnetic field strength in filaments and sheets constrained by the observed RM. In the final section, we discuss the implications of the possible existence of $\\sim 1~\\mu{\\rm G}$ or less magnetic field in filaments and sheets. ",
        "conclusions": "Our result indicates that, with the present value of the observed limit in RM, the existence of magnetic field of up to $\\sim 1~\\mu{\\rm G}$ in filaments and sheets can not be ruled out, if the cosmic magnetic field is preferentially distributed in these nonlinear structures, rather than uniformly distributed in the intergalactic space.  {\\it We emphasize that this result is independent of the details of how the magnetic field originated.} However, it is not certain if the dynamo processes such as the one considered in Kulsrud \\etal (1997) can amplify a week seed field to the above field strength.  Interestingly, there is  confirmation of such a magnetic field in one case. Kim \\etal (1989) reported the possible existence of a intercluster magnetic field of $0.3-0.6~\\mu{\\rm G}$ in the plane of the Coma/Abell 1367 supercluster. They reached the result from the observation using high dynamic range synchrotron images at $327~{\\rm MHz}$. Clearly, this result should be confirmed in other regions of the sky with the newly available possibilities of low radio frequency interferometry such as with the Giant Meterwave Radio Telescope (GMRT) in India, or the new receiver systems on the Very Large Array (VLA) in the USA.  We may compare the above upper limit with the strength of the magnetic field whose energy is in equipartition with the thermal energy of the gas in filaments and sheets. The equipartition magnetic field strength can be written as \\begin{equation} B = 0.33 h \\sqrt{T \\over 3 \\times 10^6 {\\rm K}} \\sqrt{\\rho_b \\over 0.3 \\rho_c}~\\mu{\\rm G}. \\label{b_equip} \\end{equation} The fiducial values $T = 3 \\times 10^6 {\\rm K}$ and $\\rho_b = 0.3 \\rho_c$ may considered to be appropriate for filaments (see, Kang \\etal 1994 and Figure 3). The values appropriate for sheets should be somewhat smaller. So, our upper limit set by RM is close to, but several times larger than, the equipartition magnetic field strength.  Filaments and sheets with the above fiducial temperature and gas density can contribute radiation at X-ray wavelengths to the cosmological background. Using the above values and a typical size of $10h^{-1}$Mpc, the luminosity of filaments and sheets in the soft X-ray band ($0.5{-}2\\,$keV) is expected to be $\\la 3\\times 10^{41}{\\rm erg~s^{-1}}$. Actually, Soltan \\etal (1996) and Miyaji \\etal (1996) reported the detection of extended X-ray emission from structures of a comparable size and a luminosity of $\\approx 2.5 \\times 10^{43}{\\rm erg~s^{-1}}$, which is correlated to Abell clusters. These observations suggest that the temperature and gas density outside clusters could even be considerably larger than the fiducial values assumed above. In any case, if these observations can be further confirmed, and the equipartition of magnetic field and gas energies is assumed, it would imply the existence of a $\\sim 1\\mu{\\rm G}$ magnetic field on large scales outside galaxy clusters.  The possible existence of strong magnetic field of $1~\\mu{\\rm G}$ or less in filaments and sheets has many astrophysical implications, some of which we briefly outline in the following:  The large-scale accretion shocks where seed magnetic fields could be generated are probably the biggest shocks in the universe with a typical size $\\ga$ a few (1-10) Mpc and very strong with a typical accretion velocity $\\ga$ a few $1000~{\\rm km}~{\\rm s}^{-1}$. The accretion velocity onto the shocks around clusters of a given temperature, or a give mass to radius ratio $M_{cl}/R_{cl}$, is smaller in a universe with smaller $\\Omega_o$, and is given as $v_{acc} \\approx 0.9-1.1 \\times 10^3~{\\rm km~s^{-1}} [(M_{cl}/R_{cl})/(4 \\times 10^{14}{\\rm M}_{\\odot}/{\\rm Mpc})]^{1/2}$ in model universes with $0.1 \\le \\Omega_o \\le 1$ (Ryu \\& Kang 1997a). With up to $\\sim 1~\\mu{\\rm G}$ or less magnetic field around them, the large-scale accretion shocks could serve as possible sites for the acceleration of high energy cosmic rays by the first-order Fermi process (Kang, Ryu \\& Jones 1996; Kang, Rachen \\& Biermann 1997). Although the shocks around clusters would be the most efficient sites for acceleration, those around filaments and sheets could make a significant contribution as well (Norman, Melrose \\& Achterberg 1995). With the particle diffusion model in quasi-perpendicular shocks (Jokipii 1987), the observed cosmic ray spectrum near $10^{19}{\\rm eV}$ could be explained with reasonable parameters if about $10^{-4}$ of the infalling kinetic energy can be injected into the intergalactic space as the high energy particles (if an $E^{-2}$ spectrum of cosmic rays is assumed; for a slightly steeper spectrum as suggested by radio relic sources, the efficiency is closer to 0.1, En{\\ss}lin \\etal 1998).  The discoveries of several reliable events of high energy cosmic rays at an energy above $10^{20} {\\rm eV}$ raise questions about their origin and path in the universe (a recent review is P. Biermann 1997), since their interaction with the cosmic microwave background radiation limits the distances to their sources to less than 100 Mpc, perhaps within our Local Supercluster. The Haverah Park and Akeno data indicate that their arrival directions are in some degree correlated with the direction of the Supergalactic plane (Stanev \\etal 1995; Hayashida \\etal 1996; Uchihori \\etal 1996). In Biermann, Kang \\& Ryu (1996), we noted that if the magnetic field of $\\sim 1\\mu{\\rm G}$ or less exists inside our Local Supercluster and there exist accretion flows infalling toward the supergalactic plane, it is possible that the high energy cosmic rays above the so-called GZK cutoff ($E> 5\\times 10^{19}$ eV) can be confined to the supergalactic plane sheet, an effect analogous to solar wind modulation.  In each case a shock front pushes energetic particles upstream as seen from its flow. Obviously, this effect would occur only for a small part of phase space, namely those particles with a sufficiently small initial momentum transverse to the sheet. This would explain naturally the correlation between the arrival direction of the high energy cosmic rays and the supergalactic plane. Also, confinement means that for all the particles captured into the sheets, the dilution with distance $d$ is $1/d$ instead of $1/d^2$, increasing the cosmic ray flux from any source appreciably with respect to the three-dimensional dilution. So we may see sources to much larger distances than expected so far.  On the other hand, particles with a larger initial momentum transverse to the sheet would be strongly scattered, obliterating all source information from their arrival direction at Earth.  With the magnetic field in the intergalactic medium, charged particles passing it would not only experience deflection. It would also smear out their arrival direction as well as delay their arrival time. Plaga (1995) suggested that an exhibition of this latter effect would be the delay of the arrival times of $\\gamma$-rays from a cascade caused by photons from highly time-variable extragalactic sources. However, with strong fields of $\\sim 1~\\mu{\\rm G}$ or less in filaments and sheets intervening between the sources and us, the expected consequence would be a strong smearing of the sources rather than the delay of arrival times (Kronberg 1995). But, for details, calculations following the propagation of photons and charged particles should be done.  Another interesting exploration is to study the radiation emission arising from filaments and sheets strong magnetic field of $1~\\mu{\\rm G}$ or less. We noted that the original estimate of the magnetic field in the plane of the Coma/Abell 1367 supercluster was based on a synchrotron radio continuum measurement (Kim \\etal 1989). We also noted that the thermal Bremsstrahlung emission in the soft X-ray band  may provide an interesting check on the work presented here. Whether these structures can contribute at other wavelengths, such as $\\gamma$-ray energies, to the cosmological background remains an an unanswered but challenging question at this time.  These issues will be discussed elsewhere."
    },
    "9803/astro-ph9803153_arXiv.txt": {
        "abstract": "We report $J$, $H$, and $K$ photometry of 86 stars in 40 fields in the northern hemisphere. The fields are smaller than or comparable to a 4$\\times$4~arcmin field-of-view, and are roughly uniformly distributed over the sky, making them suitable for a homogeneous broadband calibration network for near-infrared panoramic detectors. $K$ magnitudes range from 8.5 to 14, and $J-K$ colors from -0.1 to 1.2. The photometry is derived from a total of 3899 reduced images; each star has been measured, on average, 26.0 times per filter on 5.5 nights. Typical errors on the photometry are $\\sim$~0\\fm012. ",
        "introduction": "The widespread availability of near-infrared (NIR) panoramic detectors has rendered possible many scientific programs which were unfeasible with single-element photometers. New and more sophisticated data acquisition and reduction techniques have successfully exploited the capabilities of the new technology while, at the same time, photometric calibration has generally relied on older pre-existing standard networks based on ``one-pixel'' photometry. Such networks include the SAAO system of Glass (\\cite{glas}), expanded and rationalized by Carter (\\cite{carter90}, \\cite{carter95}), comprising probably the most comprehensive and best-observed list available; the MSO system defined by Jones \\& Hyland (\\cite{jone80}, \\cite{jone82}), now more or less supplanted or absorbed into the AAO system (Allen \\& Cragg \\cite{alle}); the ESO system (Engels \\cite{enge}; Bouchet, Schmider, \\& Manfroid \\cite{bouchet}); and the system which appears to the greatest extent to have inherited the original mantle of the photometry of H. L. Johnson in the 1960s, the CIT (Caltech/Cerro-Tololo) system of Frogel et al. (\\cite{frog}) updated by Elias et al. (\\cite{elia82}). This last forms the basis of the unpublished, but widely used, hybrid standard star set maintained at the 3.8~m UK Infrared Telescope (UKIRT).  All these comprise relatively bright stars suitable for photometry at small- and medium-sized telescopes with instruments which do not have the limited well capacities of array elements. However, when observed with array detectors even on moderate-sized telescopes, stars with  K $\\lesssim$ 7.5 produce saturated pixels unless defocussed or observed in non-standard modes with extremely short on-chip integration times, neither of which stratagem is conducive to precise and homogeneous calibration.  The ``UKIRT Faint Standards'' (Casali \\& Ha-warden  \\cite{casa}) upon which the calibration of this work is based comprise a new set of much fainter stars chosen and observed at UKIRT to facilitate observations with panoramic detectors with limited dynamic range. However, they are relatively few in number and isolated, occur mostly around the celestial equator, and many of the stars are too faint to be useful for small- or medium-sized telescopes. Only preliminary results are currently available for the UKIRT Faint Standards, although this situation is actively being remedied by the expansion and reobservation of the list at UKIRT.  We present here a set of $J$, $H$, and $K$ photometric measurements, obtained with the Arcetri NICMOS3 camera, ARNICA. The photometry comprises 86 stars in 40 fields observable from the northern hemisphere. The selection of the standard fields is described in Section~2, followed by a discussion of observing and data reduction techniques in Section~3. Section 4 presents the photometry and a comparison with other photometric systems. ",
        "conclusions": "We have presented new NIR photometry for 86 stars in 40 fields. The sky coverage is relatively uniform, and ideal for observatories with $\\delta\\,\\gtrsim\\,30^\\circ$. On average, stars have been observed on more than five different nights, and typical errors on the photometry are 0\\fm012. We find some indication of a color transformation between the ARNICA (NICMOS) photometry reported here and the original UKIRT (InSb) system, but a definitive statement awaits a larger data set designed specifically to determine such a transformation."
    },
    "9803/astro-ph9803015_arXiv.txt": {
        "abstract": "The BeppoSAX satellite has recently opened a new way towards the solution of the long standing gamma-ray bursts' (GRBs) enigma, providing accurate coordinates few hours after the event thus allowing for multiwavelength follow-up observational campaigns.  The BeppoSAX Narrow Field Instruments observed the region of sky containing GRB970111 16 hours after the burst. In contrast to other GRBs observed by BeppoSAX no bright afterglow was unambiguously observed. A faint source (1SAXJ1528.1+1937) is detected in a position consistent with the BeppoSAX Wide Field Camera position, but unconsistent with the IPN annulus. Whether 1SAXJ1528.1+1937 is associated with GRB970111 or not, the X-ray intensity of the afterglow is significantly lower than expected, based on the properties of the other BeppoSAX GRB afterglows. Given that GRB970111 is one of the brightest GRBs observed, this implies that there is no obvious relation between the GRB gamma-ray peak flux and the intensity of the X-ray afterglow. ",
        "introduction": "The comprehension of the nature of the Gamma-Ray Bursts (GRBs) is a long-standing problem of a world-wide scientific community since the announcement of their discovery   (Klebesadel \\etal 1973). Many observational (Fishman \\& Meegan 1995) and theoretical (Lamb 1995; Paczynski 1995) efforts did not succeed in understanding the origin of GRBs. The launch of the BeppoSAX satellite (Boella \\etal 1997a) revolutionized the field, opening a new observational window soon after the GRB event. Due to its Gamma Ray Burst Monitor (GRBM, 40--700~keV, Frontera \\etal 1997a; Feroci \\etal 1997a) and its Wide Field Cameras (WFCs, 2--26~keV, Jager \\etal 1997) this satellite is capable of detecting GRBs in the gamma-ray band and accurately localizing them in X-rays through a coded mask proportional counter.  Five GRBs, amongst those simultaneously detected by the GRBM and the WFCs, were promptly analyzed, allowing multiwavelength follow-up observational campaigns. The first result is the BeppoSAX discovery of the X-ray afterglow of GRB970228 (Costa \\etal 1997, Costa \\etal 1997a) and the discovery of a related optical transient by ground-based telescopes (van Paradijs \\etal 1997). Further results have been achieved with the detection of the X-ray afterglows of GRB970402 (Feroci \\etal 1997b, Piro \\etal 1997a), GRB970508 (Costa \\etal 1997c, Piro \\etal 1997b) and GRB971214 (Heise \\etal 1997a, Antonelli \\etal 1997). From GRB970508 an indication of an extragalactic origin has been derived through the detection of an optical transient (Bond, 1997; Djorgovski \\etal 1997) and the measurement of its optical spectrum (Metzger \\etal 1997), providing a lower limit of 0.835 for the redshift of the possible GRB optical afterglow.  One of the most intriguing mysteries of GRB emitters is possibly solved, but the overall picture is far from clear. In fact, out of the five events for which BeppoSAX performed rapid follow up searches of a X-ray counterpart, one (GRB970111) has given a result that is significantly different from the other four. The celestial location of GRB970111 was observed by BeppoSAX just 16 hours after the GRB event, and no unambiguous evidence for an X-ray afterglow was found. A new faint source (1SAX J1528.1+1937) was detected at a flux level that is much lower than that expected on the basis of the properties of the other GRBs later observed by BeppoSAX. Here we present this detection, discuss its association with GRB970111 and the diversity from the other four BeppoSAX GRBs. ",
        "conclusions": "The BeppoSAX follow-up observation of the error box of GRB970111 was the first prompt follow-up observation of a GRB ever performed by an X-ray satellite. Before BeppoSAX the time-scale of a possible X-ray emission from GRB remnants was completely unknown. This first basically non-detection, therefore, could only be interpreted as an upper limit to the time-scale of the decline of an X-ray afterglow or to its flux. Now, with the detection of the X-ray afterglows of GRB970228, GRB970402, GRB970508 and GRB971214, BeppoSAX has set up a new scenario in which GRB970111 seems misplaced. Also the detection of the X-ray afterglow of a GRB (GRB970828, Remillard \\etal 1997; Murakami \\etal 1997) by the RossiXTE and the ASCA satellites supports the general framework for the GRBs' afterglow built by BeppoSAX.  GRB970228, GRB970402 and GRB970828 showed a similar behavior, with a fading X-ray counterpart continuously decaying from the GRB main emission into the afterglow following an approximate $t^{-1.3}$ law. In the case of GRB970228, the spectral analysis (Frontera \\etal 1997b) confirms the continuity between the latest GRB emission and the X-ray counterpart detected after few hours. This temporal behaviour could be explained in the framework of the fireball model (Cavallo \\& Rees 1978; Rees \\& Meszaros 1992) as a highly radiative expansion of a relativistic shell. GRB970508 has shown a X-ray counterpart whose decay is more complicated than the above three. The above model could still account for this different behavior, but it needs to invoke a non-uniform surrounding medium, with a density scaling as $r^{-2}$  (Vietri 1997).  Whether 1SAX J1528.1+1937 is related to GRB970111 or not, this gamma-ray burst had a much faster decay than observed for any of the others. In order to make this clear, we compare a hypothetic power-law decay of GRB970111 with the ``typical'' power-law decay of GRB970228 reported in Costa \\etal (1997a). Therefore, in Fig. 3 we assume that the new X-ray source is associated with the GRB and impose a power-law decay of the afterglow starting from the WFC mean flux at a time centred on the GRB X-rays duration. The needed power-law index is -1.5.  \\begin{figure} \\epsfxsize=\\hsize   \\centerline{\\epsffile{fig_3.ps}} \\vspace{-0.5cm} \\caption[]{X-ray (2--10~keV) decay law of the candidate counterpart of GRB970111, compared to GRB970228. The dot-dashed and the solid horizontal lines are the mean X-ray flux for GRB970228 and GRB970111, respectively. The inclined dashed line is the decay law suggested by Costa et al. (1997a) for GRB970228. The inclined solid line shows the power-law index, 1.5, needed for connecting the WFC GRB970111 mean flux and the 1SAXJ 1528.1+1937 flux} \\label{fig:Decay} \\end{figure}  Alternatively, assuming that 1SAX J1528.1+1937 is not related to GRB970111 we can derive a lower limit to the power-law index by using the upper limit of the MECS flux in the region of sky defined by the error box, to obtain a value very similar to the 1.5 value given above.  Trying to extract GRB970111 from the group as an intrinsically different event, we note that its gamma-ray fluence is about more than three times larger than the largest of the other three. On the other hand, even if this GRB is of the ``No High Eenergy'' type (that is, it shows only weak emission above 300 keV, Pendleton et al. 1997), the ratio between the X-ray (2--10~keV) and gamma-ray (40--700~keV) fluences is about 4\\%, to be compared to 20\\% (2--10~keV) for GRB970228 (Frontera \\etal 1997b), 5\\% (2--10~keV) for GRB970402 (Nicastro \\etal 1997) and 40\\% (2--26~keV) for GRB970508 (Piro \\etal 1997b). GRB970111 appears therefore as the one (together with the April event) with the less efficient low X-rays mechanism for energy release. Furthermore, no optical source was found in the WFC error box changing its intensity more than 0.5 magnitudes at a level  of B=23 and R=22.6 from about 19 hours to about one month later (Castro-Tirado \\etal 1997; Gorosabel \\etal 1998). A radio search at 1.4 GHz (Frail \\etal 1997) and at millimetric wavelength (Smith \\etal 1997) did not find a counterpart to 1SAXJ1528.1+1937. These results support the idea that the optical, radio and millimetric channels are unefficient as well. Since GRB970111 was one of the brightest events ever detected in gamma-rays, one may conclude that its gamma-ray channel was efficient enough to dissipate most of the energy generated in the burst.  Alternative interpretations of the lack of X\\-/optical\\-/radio afterglow of the GRB970111 may be either a very rapidly evolving afterglow, with a decay law faster than observed in the other BeppoSAX GRB afterglows, or the absence of an afterglow source. The former hypothesis would be in agreement with the model by Tavani (1997) of a decay behavior represented by a power law with index $-21/8$ due to the observation in a fixed energy band (2--10~keV) of a synchrotron emission spectrum with a rapidly evolving critical frequency. Alternatively, the latter situation could be due, as an example, to the scenario in which the event that caused the GRB occurred in a region in which the interstellar medium density is low enough (perhaps the external regions of a host galaxy) to justify the absence of an external shock, possibly responsible for the afterglow emission in the other cases (Katz \\& Piran 1997)."
    },
    "9803/gr-qc9803028_arXiv.txt": {
        "abstract": " ",
        "introduction": "The energy range between the grand unification scale $M_{\\rm GUT}\\sim 10^{16}$ GeV and the Planck scale $\\mpl\\simeq 1.22\\times 10^{19}$ GeV is crucial for fundamental physical questions and for testing current ideas about grand unification, quantum gravity, string theory. Experimental results  in this  energy range are  of course very difficult to obtain. From the particle physics point of view, there are basically two important experimental results that can be translated into statements about this energy range (see e.g.~\\cite{rev} for recent reviews): (i) the accurate measurement of gauge coupling constants at LEP that, combined with their running with energy, shows the unification of the couplings at the scale $\\mgut$, provided that the running is computed including the supersymmetric particles in the low energy spectrum. And (ii) the negative results on proton decay; the lower limit on the inverse of the partial decay width for the processes $p\\ra e^+\\pi^0$ and $p\\ra K^+\\bar{\\nu}$ are $5.5\\times 10^{32}$ yr and $1.0\\times 10^{32}$ yr, respectively, and imply a lower bound  $\\mgut \\gsim 10^{15}$ GeV, which excludes non-supersymmetric SU(5) unification. Further improvement is expected from the SuperKamiokande experiment, which should probe lifetimes $\\sim 10^{34}$ yr.  From the cosmological point of view, informations on this energy range can only come from particles which decoupled  from the primordial plasma at  very early time. Particles which stay in thermal equilibrium down to a decoupling temperature $T_{\\rm dec}$ can only carry informations on the state of the Universe at $E\\sim T_{\\rm dec}$. All informations on physics at higher energies has  in fact been obliterated by the successive interactions.  The condition for thermal equilibrium is that the rate $\\Gamma$ of the processes that mantain equilibrium be larger than the rate of expansion of the Universe, as measured by the Hubble parameter $H$~\\cite{KT}. The rate is given by $\\Gamma =n\\sigma |v|$ where $n$ is the number density of the particle in question, and for massless or light particles in equilibrium at a temperature $T$, $n\\sim T^3$;  $|v|\\sim 1$ is the typical velocity and $\\sigma$ is the cross-section of the process. Consider for instance the weakly interacting neutrinos. In this case the equilibrium is mantained, e.g., by electron-neutrino scattering, and at energies below the $W$ mass $\\sigma\\sim G_F^2\\langle E^2\\rangle \\sim G_F^2T^2$ where $G_F$ is the Fermi constant and $\\langle E^2\\rangle$ is the average energy squared. The Hubble parameter during the radiation dominated era is related to the temperature by $H\\sim T^2/\\mpl$. Therefore~\\cite{KT} \\be \\left(\\frac{\\Gamma}{H}\\right)_{\\rm neutrino} \\sim \\frac{G_F^2T^5}{T^2/\\mpl}\\simeq\\left( \\frac{T}{1\\rm MeV}\\right)^3\\, . \\ee Even the weakly interacting neutrinos, therefore, cannot carry informations on the state of the Universe at temperatures larger than approximately 1 MeV. If we repeat the above computation for gravitons, the Fermi constant $G_F$ is replaced by Newton constant $G=1/\\mpl^2$ (we always use units $\\hbar =c=k_B=1$) and at energies below the Planck mass \\be \\left(\\frac{\\Gamma}{H}\\right)_{\\rm graviton} \\sim \\left(\\frac{T}{\\mpl}\\right)^3\\, . \\ee The gravitons are therefore decoupled below the Planck scale. (At the Planck scale the above estimate of the cross section is not valid and nothing can be said without a quantum theory of gravity). It follows that relic gravitational waves  are a potential source of informations on very high-energy physics. Gravitational waves produced in the very early Universe have not lost memory of the conditions in which they have been produced, as it happened to all other particles, but still retain in their spectrum, typical frequency and intensity, important informations on the state of the very early Universe,  and therefore on physics at correspondingly high energies, which cannot be accessed experimentally in any other way. It is also clear that the property of gravitational waves that makes them so interesting, i.e. their extremely small cross section, is also responsible for the difficulties of the experimental detection.\\footnote{Thinking in terms of cross-sections, one is lead to ask  how comes that gravitons could be detectable altogheter, since the graviton-matter cross section is smaller than the neutrino-matter cross section, at energies below the $W$-mass, by a factor $G^2/G_F^2\\sim 10^{-67}$ and neutrinos are already so difficult to detect. The answer is that gravitons are bosons, and therefore their occupation number per cell of phase space can be $n_k\\gg 1$; we will see below that in interesting cases, in the relic stochastic background we can have $n_k\\sim 10^{40}$ or larger, and the squared amplitude for exciting a given mode ot the detector grows as $n_k^2$. So, we will never really detect gravitons, but rather classical gravitational waves. Neutrinos, in contrast, are fermions and for them $n_k\\leq 1$.}  With the very limited experimental informations that we have on the very high energy region,  $\\mgut \\lsim E\\lsim \\mpl$, it is unlikely that  theorists will be able to foresee all the interesting sources of relic  stochastic background, let alone to compute their spectra. This is particularly clear in the Planckian or string theory domain where, even if we succeed in predicting some interesting physical effects, in general we cannot compute them reliably. So, despite the large efforts that have been devoted to understanding possible sources, it is still quite possible that, if a relic background of gravitational waves will be detected, its physical origin will be a surprise. In this case a model-independent analysis of what we can expect might be useful.  In this paper we discuss  whether, from the experience gained with various specific computations of relic backgrounds,  it is possible to extract statements or order of magnitude estimates which are as much as possible model-independent. These estimates would constitute a sort of minimal set of naive expectations, that could give some orientation, independently of the uncertainties and intricacies of the specific cosmological models. We discuss  typical values of the frequencies involved and of the  expected intensity of the background gravitational radiation, and we try to distinguish between statements that are relatively model-independent and results  specific to given models.  The paper is written having in mind a reader interested in gravitational-wave detection but not necessarily competent in early Universe cosmology nor in physics at the string or Planck scale, and a number of more technical remarks are relegated in footnotes and in an appendix. We have also tried to be self-contained and we have attempted to summarize and occasionally clean up many formulas and numerical estimates appearing in the literature. The organization of the paper is as follows. In sect.~2 we introduce the variables most commonly used to describe a stochastic background of gravitational waves. We give a detailed derivation of the relation between exact formulas for the signal-to-noise ratio, and approximate but simpler characterizations of the characteristic amplitude and of the noise. The former variables are convenient in theoretical computations while the latter are commonly used by experimentalists, so it is worthwhile to understand in some details their relations. In sect.~3 we apply these formulas to compute the sensitivity to a stochastic background that could be obtained with a second Virgo interferometer correlated with the first, and we compare with various others detectors. We find that in the Virgo-Virgo  case the  noise which would give the dominant limitation to the measurement of a stochastic background is the mirror thermal noise, and we give the sensitivity for different forms of the relic GW spectrum. In sect.~4 and 5 we discuss estimates of the typical frequency scales. We examine the statements leading to the conclusion that Virgo/LIGO will explore the Universe at temperatures $T\\sim 10^7$ GeV (sect.~4), and in the appendix we discuss some qualifications to this statement. In sect.~5  we discuss the possibility to reach much higher energy scales, including the typical scales of grand unification and quantum gravity. In sect.~6 we discuss different scenarios, depending on how the inflationary paradigm is implemented. Characteristic values of the intensity of the spectrum and existing limits and predictions are discussed in sect.~7. Sect.~8 contains the conclusions. ",
        "conclusions": "Present GW experiments have not been designed especially for the detection of  GW backgrounds of cosmological origin. Nevertheless, there are chances that in their frequency window there  might be a cosmological signal. The most naive estimate of the frequency range for signals from the very early Universe singles out the GHz region, very far from the region accessible to ground based interferometers, $f<$ a few kHz, or to resonant masses. To have a signal in the accessible region, one of these two conditons should be met: either we find a spectrum with a long low-frequency tail, that extends from the GHz down to the kHz region, or we have some explosive production mechanism much below the Planck scale. As we have discussed, both situations seem to be not at all unusual, at least in the examples that have been worked out to date. The crucial point is the value of the intensity of the background.  With a very optimistic attitude, one could hope for a signal, present just in the 10Hz-1~kHz band, with the maximum  intensity compatible with the nucleosynthesis bound,  $\\hogw\\sim$ a few $\\times 10^{-6}$ (or even $10^{-5}$, stretching all parameters to the maximum limit). Such an option is not excluded, and the fact that such a background is not predicted by the mechanisms that have been investigated to date is probably  not a very strong objection, given our theoretical ignorance of physics at the Planck scale and the rate at which new production mechanisms have been proposed in recent years, see the reference list. However, with more realistic estimates, on general grounds it appears difficult to  predict  a background that in the kHz region exceeds the level $\\hogw\\sim$ a few $\\times 10^{-7}$, independently of the production mechanism. This should be considered the minimum detection level at which a significant search can start. Such a level is beyond the sensitivity of first generation experiments, unless a ground based interferometer is correlated with a second interferometer, located at a distance small enough so that a significant correlation is possible, and large enough to decorrelate local noises. A few tens of kilometers would probably be the right order of magnitude. In this case we could detect a signal at the level $2\\times 10^{-7}$ in one year of integration time, with SNR=1.65, and the level of a few $\\times 10^{-8}$ could be reached with longer integration time (and possibly allowing for a slightly worse confidence level; this could make some sense in a stochastic search because the SNR increases with time, and the hint of a signal at low SNR would provide a strong motivation for pursuing the search). The difficulties of such a detection are clear, but it should be stressed that the payoff of a positive result would be enormous, opening up a window in the Universe and in fundamental high-energy physics that will never be reached with particle physics experiments.   \\vspace*{5mm}  {\\bf Acknowledgments.} I am very grateful to Adalberto Giazotto for many interesting discussions and stimulating questions, which prompted me to write down this paper. I thank Valeria Ferrari and Raffaella Schneider for discussing with me their unpublished results on astrophysical backgrounds.  I also thank for useful discussions or comments on the manuscript Pia Astone, Danilo Babusci, Carlo Baccigalupi, Alessandra Buonanno, Massimo Cerdonio, Eugenio Coccia, Stefano Foffa, Maurizio Gasperini, Marco Lombardi, Emilio Picasso, Riccardo Sturani and  Andrea Vicer\\`e.  \\appendix"
    },
    "9803/astro-ph9803290_arXiv.txt": {
        "abstract": "\\noindent We describe a method of determining the system parameters in non-eclipsing interacting binaries. We find that the extent to which an observer sees the shape of the Roche-lobe of the secondary star governs the amount of distortion of the absorption line profiles.  The width and degree of asymmetry of the phase-resolved absorption line profiles show a characteristic shape, which depends primarily on the binary inclination and gravity darkening exponent. We show that, in principle, by obtaining high spectral and time resolution spectra of quiescent cataclysmic variables or low mass X-ray binaries in which the mass-losing star is visible, fitting the shape of absorption line profiles will allow one to determine not only the mass function of the binary, but also the binary inclination and hence the mass of the binary components. ",
        "introduction": "The determination of the binary inclination in non-eclipsing interacting binaries has been problematic for many years. One method widely used to determine the binary inclination in dwarf novae and the soft X-ray transients is to measure the ellipsoidal variations of the late-type star (Warner 1995; van Paradijs \\& McClintock 1995). These variations are caused by the companion star presenting differing aspects of its distortion to the observer, giving rise to a double-humped modulation whose amplitude is strongly dependent on the binary inclination. However, one problem in measuring the ellipsoidal variations of the secondary star is that the accretion disc can contribute a significant amount of flux at optical wavelengths (e.g. typically about 10--50 per cent in the soft X-ray transients). This contribution must be accounted for if the binary inclination is to be determined using optical light curves (see Charles 1996 and references within).  Here we describe a technique which uses the information available about the shape of the Roche-lobe of the secondary star, and its effect on the shape of the absorption line profiles. The parameters we measure are insensitive to the disc contribution, but dependent on the binary inclination and gravity darkening exponent. We will first give a brief description of the model and then describe the effect of the mass ratio, inclination, limb and gravity darkening on the shape of the line profiles. Finally, we will simulate data and proceed to fit it using the model. ",
        "conclusions": "By obtaining high spectral and time resolution spectra of cataclysmic variables or low mass X-ray binaries, systems in which the absorption features of the secondary star can be seen, one can perform a radial velocity study of the secondary star and also determine the binary mass ratio (see Marsh, Robinson \\& Wood 1994 for a full description of this kind of study). However, as pointed out in this paper, it is also possible to extract the binary inclination by fitting the absorption lines profiles of the secondary star.  Some of the assumptions inherent in the model which predicts the shape of the line profiles should be noted. The main assumption in the model is that the observed surface of constant optical depth coincides with the Roche potential surface. However, we find that this only changes the shape of the line profiles by less than 1 per cent. Zonal flow patterns and/or dark- hot-spots can have an appreciable effect on the shape of the line profiles. However, the extent of these effects depends very much on the magnitude of the zonal flows and the size of the spots. Simulations show that only very low velocities are predicted for zonal flows (Martin \\& Davey 1995).  The characteristic modulations in the width of the line profiles are, in principle, similar to the ellipsoidal modulations, which are due to the observer seeing differing aspects of the gravitationally distorted secondary star as it orbits a compact object. However, this method uses the velocity information across the star as well as the projected area. It should also be noted that unlike modelling the ellipsoidal variations in the optical, where the accretion disc contribution must be taken into account, this method is independent of the disc contribution (as long as the disc contribution is not such that it totally swamps the secondary star features). This implies all emission and continuum sources, i.e. from a bright spot or a gas stream, will only affect the determination of the veiling factor and not the shape of the absorption lines. However, care must be taken in choosing the lines to use for this kind of analysis; the lines must be clear of any weak emission features arising from the disc or accretion flow.  The determination of the binary inclination is essential if one wants to obtain the mass of the binary components in non-eclipsing binaries. The method described here can be applied to bright cataclysmic variable stars or low mass X-ray binaries, in which one can resolve the absorption lines of the secondary star, and where the amount of X-ray heating is small. The effect of heating implies that one cannot use temperature sensitive absorption lines such as the Na $\\sc i$ 8183-8184 \\AA\\ doublet in the modelling (see Friend et al. 1990 and references within for the effects of irradiation). Such systems are the dwarf novae and the soft X-ray transients. This method allows one to determine all the system parameters from a high spectral resolution (few km~s$^{-1}$) spectroscopic study.  We have described a method of determining the binary inclination in non-eclipsing interacting binaries by fitting the shape of the absorption lines arising from the secondary star. We find that the amount of distortion of the absorption line profiles is primarily due to the extent to which an observer sees the shape of the Roche-lobe of the secondary star. We show that, in principle, by obtaining high spectral and time resolution spectra of quiescent dwarf novae or the soft X-ray transients, where the disc is low, fitting the shape of absorption line profiles will allow one to determine the binary inclination. Our simulations show that previous efforts to determine the inclination are flawed, and give systematically lower values for the inclination."
    },
    "9803/astro-ph9803259_arXiv.txt": {
        "abstract": "\\cite{peeb1} showed that in the gravitational instability picture galaxy orbits can be traced back in time from a knowledge of their current positions, via a variational principle. We modify this variational principle so that galaxy redshifts can be input instead of distances, thereby recovering the distances. As a test problem, we apply the new method to a Local Group model. We infer $M = 4$ to $8 \\times 10^{12} M_{\\sun} $ depending on cosmology, implying that the dynamics of the outlying Local Group dwarves are consistent with the timing argument. Some algorithmic issues need to be addressed before the method can be applied to recover nonlinear evolution from large redshift surveys, but there are no more difficulties in principle. ",
        "introduction": "Phase space is six--dimensional and therefore six numbers for each particle will specify the dynamics completely. The standard approach is to define initial conditions as six numbers (positions and velocities in three dimensions) and integrate those forward in time. This is the usual $N$-body approach to the problem.  Alternatively, it is not necessary to define those six numbers at only one time. It is equally well possible to split them, such that three numbers will be known at an initial time, and three at a final time, where the former are derived from physical arguments about the inital state, and the latter are provided by data on the current state of the system. This is the boundary value approach.  The usual boundary value for structure formation by hierarchical clustering at initial times is the gravitational instability requirement that initial peculiar velocities must vanish: therefore, it is possible to express the orbits as a sum of growing modes. This is what perturbation theory is designed to do, and linear perturbation theory and its extension, the Zel'dovich approximation are in wide use. However, as structure formation is non-linear on small scales, a non-linear formalism would be more useful. Non-linear perturbation theory exists (e.g. \\cite{nus}, \\cite{gram93a}, \\cite{gram93b}, \\cite{buch}) but is very complicated and has convergence problems, i.e. the dynamics at early times can be fitted to a high accuracy only at the expense of a good fit at later times.  A different method (cf \\cite{peeb1}) of addressing the boundary value problem is to use the variational principle. The basic method is to start with a parameterisation of the orbits which satisfies the boundary conditions and then adjust the parameters until a stationary point of the action is found. (A variant, suggested by \\cite{gia93} and implemented by \\cite{susp}, parameterises the density and velocity fields rather than the orbits.) This may seem like perturbation theory because it is a method which attempts to get better and better approximations of galaxy orbits but there are important differences. The main one is that perturbation theory attempts to fit early times even at the expense of accuracy at later times, whereas the variational principle spreads out the errors more uniformly over all times. The variational principle is also algorithmically more straightforward to do to higher orders.  The main disadvantage of Peebles' original method is that it requires the input of distances to recover redshifts. Redshifts, however, are easy to measure, whereas distances are not. We therefore change to recovering distances from redshifts. The recent nearly all-sky redshift surveys QDOT and PSC$z$ provide a strong motivation for variational methods in redshift space. \\cite{shay93} and \\cite{shay95} have pointed out that this could be achieved by treating the distance boundary condition as a parameter and fitting that until the recovered redshifts agree with the measured ones. Another approach is to modify the variational principle until the boundary conditions are of the desired form. We take such an approach and so does \\cite{whit98} but the details in his treatment differ from ours. The attractive feature of this approach is that it only requires small modifications of Peebles' elegant original method.  In this paper we develop a redshift space variational method, apply it to a small system --- the Local Group ---, and point out the algorithmic problems that need to be solved to extend to larger systems.  Various aspects of the Local Group dynamics have been studied by \\cite{peeb1}, \\cite{peeb2}, \\cite{peeb4}, and \\cite{dunn}. The new feature of our analysis is that we attempt to constrain the masses of the Milky Way, M31, and an underlying distribution of unclustered matter by using a likelihood approach.  Two points need to be made about these variational methods in general: Firstly, the true solutions of the variational action need not be a minimum even though the colloquial use of `least action' has stuck (cf \\cite{peeb2}). Secondly, the equations solved are the same as for $N$-body integration, only the approximations used to solve them are different. In neither of the two approaches do particles have to be galaxies -- they may well be samplers of some underlying distribution function and as \\cite{branch} and \\cite{dunn2} have pointed out, applying the variational method using galaxies only, neglecting biasing and unclustered matter, leads to wrong results. ",
        "conclusions": "The $N$-body check proves that the code works and is ready to be extended to larger systems. The main problem with a system like the Local Group is the occurrence of multiple solutions. We have no guarantee that the solutions we find are the real ones, even if they fit all redshifts perfectly and predict distances which are not too far out from the measured ones. In fact, some tests suggests that at least for certain masses, several possible solutions are very close to each other so it is easy to pick the wrong one. This problem will not occur in larger systems, so they should in fact be easier to deal with.  In future work, two issues will need to be addressed: First, we need an approximate and more efficient way of doing the force calculations in equations \\ref{reals} and \\ref{reds}. The direct sum that we have used needs to be replaced by a standard $N$-body method. Second, we need more efficient convergence. As mentioned above, the main difficulty lies in preventing the coefficients from finding two-cycles instead of fixed points of the iteration. A faster yet still robust method for dealing with this problem is needed.  Even when these problems are solved, the variational calculations will never have as many particles as $N$-body simulations. The reason for pursuing this method is that unlike the $N$-body calculation, which can only reproduce the current state of the system in a statistical sense, the variational method can fit the current state exactly.  {\\vskip 10pt}"
    },
    "9803/astro-ph9803284_arXiv.txt": {
        "abstract": "Ongoing searches for supernovae (SNe) at cosmological distances have recently started to provide a link between SN Ia statistics and galaxy evolution. We use recent estimates of the global history of star formation to compute the theoretical Type Ia and Type II SN rates as a function of cosmic time from the present epoch to high redshifts. We show that accurate measurements of the frequency of SN events in the range $0<z<1$ will be valuable probes of the nature of Type Ia progenitors and the evolution of the stellar birthrate in the universe. The {\\it Next Generation Space Telescope} should detect of order 20 Type II SNe per $4'\\times 4'$ field per year in the interval $1<z<4$. ",
        "introduction": "The remarkable progress in our understanding of faint galaxy data made possible by the combination of HST deep imaging (Williams \\etal 1996) and ground-based spectroscopy (Lilly \\etal 1995; Ellis \\etal 1996; Cowie \\etal 1996; Steidel \\etal 1996), has recently permitted to shed some light on the evolution of the stellar birthrate in the universe, to identify the epoch $1\\lta z\\lta 2$ where most of the optical extragalactic background light was produced, and to set important contraints on galaxy formation scenarios (Madau \\etal 1998; Steidel \\etal 1998). While one of the biggest uncertainties in our knowledge of the emission history of the universe is probably represented by the poorly constrained amount of starlight that was absorbed by dust and reradiated in the IR at early and late epochs, one could also imagine the existence of a large population of relatively old or faint galaxies still undetected at high-$z$, as the color-selected ground-based and {\\it Hubble Deep Field} samples include only the most actively star-forming young objects. It is then important at this stage to devise different observational strategies, free of some of the biases that plague current galaxy surveys, and to make testable predictions for future astronomical capabilities, such as SIRTF, FIRST and NGST.  Here, we shall focus our attention on the rate of supernova (SN) explosions in the universe. An obvious reason to consider SNe is purely observational, i.e. the fact that they are very bright objects, with luminosities as high as $10^{10}~L_\\odot$, and are point-like sources, making their detection possible even at very large distances/redshifts.  More in general, the evolution of the SN rate with redshift contains unique information on the star formation history of the universe, the initial mass function (IMF) of stars, and the nature of the binary companion in Type Ia events. All are essential ingredients for understanding galaxy formation, cosmic chemical evolution, and the mechanisms which determined the efficiency of the conversion of gas into stars in galaxies at various epochs (e.g. Madau \\etal 1996; Madau, Pozzetti, \\& Dickinson 1997; Renzini 1997). While the frequency of ``core-collapse supernovae'', SN~II and possibly SN~Ib/c, which have short-lived progenitors (e.g. Wheeler \\& Swartz 1993) is essentially related, for a given IMF, to the instantaneous stellar birthrate of massive stars, Type Ia SNe -- which are believed to result from the thermonuclear disruption of C-O white dwarfs in binary systems -- follow a slower evolutionary clock, and can then be used as a probe of the past history of star formation in galaxies (e.g. Branch \\etal 1995; Ruiz-Lapuente, Canal, \\& Burkert 1997; Yungelson \\& Livio 1998). The recent detection of Type Ia SNe at cosmological distances (Kim \\etal 1997; Garnavich \\etal 1998; Perlmutter \\etal 1998) allow for the first time a detailed comparison between the SN rates self-consistently predicted by stellar evolution models that reproduce the optical spectrophotometric properties of field galaxies, and the observed values.  In this {\\it Letter} we show how accurate measurements at low and intermediate redshifts  of the frequencies of Type II(+Ib/c) and Ia SNe could be used as an independent test for the star formation and heavy element enrichment history of the universe,  and significantly improve our understanding of the intrinsic nature and age of the populations involved in the SN explosions. A determination of the amount of star formation at early epochs is of crucial importance, as the two competing scenarios for galaxy formation, monolithic collapse -- where spheroidal systems formed early and rapidly, experiencing a bright starburst phase at high-$z$ (Eggen, Lynden-Bell, \\& Sandage 1962; Tinsley \\& Gunn 1976) -- and hierarchical clustering -- where ellipticals form continuosly by the merger of disk/bulge systems (White \\& Frenk 1991; Kauffmann \\etal 1993) and most galaxies never experience star formation rates in excess of a few solar masses per year (Baugh \\etal 1998) -- appear to make rather different predictions in this regard. We show how, by detecting Type II SNe at high-$z$, the {\\it Next Generation Space Telescope} should provide an important test for distinguishing between different scenarios of galaxy formation. ",
        "conclusions": "We have investigated the link between SN statistics and galaxy evolution. Using recent determinations of the star formation history of field galaxies from the present epoch to high-$z$, and a simple model for the evolutionary history of the binary system leading to a Type Ia event, we have computed the theoretical Type Ia and Type II SN rates as a function of cosmic time. While significant uncertainties still remain in these estimates, we believe the calculations presented in this {\\it Letter} offer a first, realistic glimpse to the evolution of the cosmic supernova rates with cosmic time. Our main results can be summarized as follows.  \\begin{itemize}  \\item At the present epoch, the predicted Type II(+Ib/c) frequency appears to match remarkably well the observed local value. The obvious caveat is the well known fact that the SN II rate is a sensitive function of the lower mass cutoff of the progenitors, $m_l$. Values as low as $m_l=6\\,\\msun$ (Chiosi, Bertelli, \\& Bressan 1992) or as high as  $m_l=11\\,\\msun$ (Nomoto 1984) have been proposed in the literature: adopting a lower  mass limit of 6 or 11 $\\msun$ would increase or reduce our Type II rates by a factor 1.5, respectively. Note also that rates obtained from traditional distant (beyond 4 Mpc) sample might need to be increased by a factor of 1.5--2 because of severe selection effects against Type II's fainter than $M_V=-16$ (Woltjer 1997).  \\item In the interval $0\\lta z\\lta 1$, the predicted rate of SN Ia is a sensitive function of the characteristic delay timescale between the collapse of the primary star to a WD and the SN event. Accurate measurements of SN rates in this redshift range will improve our understanding of the nature of SN~Ia progenitors and the physics of the explosions. Ongoing searches and studies of distant SNe should soon provide these rates, allowing a universal calibration of the Type Ia phenomenon.  \\item While Type Ia rates at $1\\lta z\\lta 2$ will offer valuable information on the star formation history of the universe at earlier epoch, the full picture will only be obtained with statistics on Type Ia and II SNe at redshifts $2<z<4$ or higher. At these epochs, the detection of Type II events must await the {\\it Next Generation Space Telescope} (NGST). A SN II has a typical peak magnitude $M_B\\approx -17$ (e.g. Patat \\etal 1994): placed at $z=3$, such an explosion would give rise to an observed flux of 15 nJy (assuming a flat cosmology with $q_0=0.5$ and $H_0=50\\,h_{50}\\kmsmpc$) at 1.8 \\micron. At this wavelength, the imaging sensitivity of an 8m NGST is 1 nJy ($10^4$ s exposure and $10\\sigma$ detection threshold), while the moderate resolution ($\\lambda/\\Delta \\lambda=1000$) spectroscopic limit is about 50 times higher ($10^5$ s exposure per resolution element and $10\\sigma$ detection threshold) (Stockman \\etal 1998). The several weeks period of peak rest-frame blue luminosity would be stretched by a factor of $(1+z)$ to few months. Figure 3 shows the cumulative number of Type II events expected per year per $4'\\times 4'$ field. Depending on the history of star formation at high redshifts, the NGST should detect between 7 (in the merging model) and 15 (in the monolithic collapse scenario) Type II SNe per field per year in the interval $2<z<4$. The possibility of detecting Type II SNe at $z\\gta5$ from an early population of galaxies has been investigated by Miralda-Escud\\'e \\& Rees (1997). By assuming these are responsible for the generation of all the metals observed in the Lyman-$\\alpha$ forest at high redshifts, a high baryon density ($\\Omega_bh_{50}^2=0.1$), and an average metallicity of $0.01Z_\\odot$, Miralda-Escud\\'e \\& Rees estimate the NGST should observe about 16 SN II per field per year with $z\\gta 5$. Note, however, that a metallicity smaller by a factor $\\sim 10$ compared to the value adopted by these authors has been recently derived by Songaila (1997). For comparison, the models discussed in this {\\it Letter} predict between 1 and 10 Type II SNe per field per year with $z\\gta 4$.  \\end{itemize}"
    },
    "9803/astro-ph9803237_arXiv.txt": {
        "abstract": "We present quantitative statistical evidence for a $\\gamma$-ray emission halo surrounding the Galaxy.  Maps of the emission are derived. EGRET data were analyzed in a wavelet-based non-parametric hypothesis testing framework, using a model of expected diffuse (Galactic $+$ isotropic) emission as a null hypothesis.  The results show a statistically significant large scale halo surrounding the center of the Milky Way as seen from Earth.  The halo flux at high latitudes is somewhat smaller than the isotropic $\\gamma$-ray flux at the same energy, though of the same order ($O(10^{-7}$--$10^{-6})$ ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$ above 1 GeV). ",
        "introduction": "The study of diffuse high-energy $\\gamma$-ray emission has proceeded along two complementary paths.  The first is via the use of parametric methods (see e.g. \\citeasnoun{hunter}), where a physically motivated model described by a small number of parameters is fit to the data. The alternate approach (see e.g. \\citeasnoun{chen} and \\citeasnoun{willis}) is non-parametric, where one attempts to find the flux distribution as a function of position without reference to a particular model. The two classes complement each other in the sense that the strengths of one are often the weaknesses of the other.  For instance, while a parametric fit usually gives some global measure of how good the fit is, but does not tell you where the model fails (e.g., ``the model did not account for this blob of flux over there'').  The non-parametric approach can give this information, but unlike the parametric analysis, the quantitative assessment of the results is complicated by the effects of statistical bias.  This latter point should be born in mind when examining the results presented in this paper.  Non-parametric analysis of photon-limited data is beset by certain difficulties, key of which is that Poisson noise is neither stationary nor additive.  Thus, the results of what would be a straightforward analysis (e.g. smoothing by a Gaussian kernel) for data contaminated by white noise are more difficult to interpret. We therefore apply a new wavelet-based technique which has the following characteristics: \\begin{itemize} \\item{Rigorous treatment of Poisson statistics.} \\item{Assessment (in some sense) of the statistical significance of the results.} \\item{Spatial adaptivity, in that structures at different size scales are recovered automatically.} \\end{itemize}  In this paper we apply non-parametric analysis to EGRET data taken during Phases 1-4 of the Compton Gamma Ray Observatory (CGRO) mission. Below we describe the non-parametric analysis method, and present results from EGRET data showing an extended halo of $\\gamma$-ray emission apparently surrounding the center of the Milky Way.  We show that this halo is statistically very strong in the data, and not obviously attributable to any systematic effect of the analysis or instrument, from which we conclude that it is most likely of astrophysical origin.  We conclude with some brief discussion on the possible origins and implications of the $\\gamma$-ray halo.  This work expands upon that first presented in \\citeasnoun{dixon1} and \\citeasnoun{dixon2}. ",
        "conclusions": "We have presented strong statistical evidence for a large scale anisotropic excess in high energy $\\gamma$-rays.  Examination of our maps indicates that this excess may originate in the Galactic halo, though our results in no way rule out a local origin.  The spectrum appears to be broadband; detailed investigation of the spectral properties will be the subject of future work.  The origin of the halo is unclear. Emission from inverse Compton in a large halo \\cite{moskalenko} appears to be a good candidate but it remains to be seen whether it can account for the entire observations in detail. A definitive answer on this topic will perhaps require a ``smoking gun'' in this or another energy band, or may have to wait for the GLAST mission.  Further, as noted by previous authors (\\citeasnoun{smialkowski}; \\citeasnoun{chen}; \\citeasnoun{strong1}), the existence of such an extended excess may impact estimates of the extragalactic $\\gamma$-ray background."
    },
    "9803/astro-ph9803001_arXiv.txt": {
        "abstract": "We carry out CCD photometry of galaxies in the $5.25$ square region centered on Coma cluster down to $M_R=-16.0$, beyond the limit of conventional morphological classification. We use the angular two-point correlation function as well as radial profiles in order to characterize the luminosity segregation. We find strong luminosity segregation for our total sample over the magnitude range of $-20 \\leq M_R \\leq -16$, which is not entirely accounted for in terms of the morphology-density relation that is known to exist only for bright galaxies. We use a single consistent parameter, the degree of luminosity concentration, to parameterize the morphology of galaxies over the wide magnitude range, where both giant and dwarf galaxies are included. Galaxies with high central concentration (HCC) show strong luminosity segregation, \\ie their clustering strength depends strongly on luminosity while those with low central concentration (LCC) show almost no luminosity segregation. Radial density profile shows that brighter HCC-type galaxies tend to more strongly concentrate near the cluster center while LCC-type galaxies do not show such a dependence on luminosity. We show that these results are tenable against the contamination by field galaxies and uncertainties in our method of classification. ",
        "introduction": "Some studies suggest that luminous galaxies are clustered more strongly than faint galaxies. This phenomenon is referred to as the luminosity segregation (hereafter LS). The LS can be interpreted as a result of either primordial effects or environmental effects. Explanations based on the primordial effect include the biased Cold Dark Matter (CDM) model in which galaxies form at high peaks in the density field (e.g., \\cite{kai84}; \\cite{dav85}). It is known that the LS is naturally predicted by this model. Valls-Gabaud, Alimi, \\& Blanchard (1989) pointed out that the correlation strength has a positive dependence on galaxy luminosity on the basis of the biased-CDM model. White et al. (1987) also predicted that clustering strength is a strong function of the circular velocity of galaxies, which in turn correlates with the luminosity as indicated by Tully-Fisher or Faber-Jackson relations. In terms of environmental effects, the LS can be regarded as a result of frequent merging or other dynamical mechanisms in the vicinity of the cluster core.  Observational evidence for the LS has been uncertain and controversial. Some found positive results (e.g., Capelato et al. 1980; Dom\\'inguez-Tenreiro \\& Pozo-Sanz 1988; Davis et al. 1988), and others found negative results (e.g., Phillips \\& Shanks 1987; Einasto 1991). The clustering property of galaxies is also correlated with their morphology, which is so-called the morphology segregation, or morphology-density relation (\\cite{dre80b}). If the morphology segregation is the fundamental correlation, the LS would be naturally expected. Early-type galaxies show a stronger degree of clustering than late-type galaxies, and at the same time, early-type galaxies are on the average brighter than late-type galaxies (e.g., \\cite{eep88}). Consequently, galaxies which are more strongly clustered are brighter than those which are less clustered, which is the LS. On the contrary, if the LS is an essential correlation, the morphology segregation would also be expected. Accordingly, it is critical to see if the LS is observed {\\it within a given morphological type} in order to disentangle the coupling of the LS and morphology segregation. Only a few such studies have been made so far (\\cite{ein91}; \\cite{lov95}).  There are also few studies on the LS among dwarf galaxies. Binggeli, Tammann, \\& Sandage ($1987$) found in the Virgo cluster that nucleated dwarf ellipticals (dEs) are more strongly concentrated towards the cluster center than nonnucleated dEs. Ferguson \\& Sandage ($1989$) found in the Virgo and Fornax clusters that the faint ($M_B>-13.3$) nonnucleated dEs have the distribution identical to that of the E/S0 galaxies and bright nucleated dEs. They also found that the faint nonnucleated dEs are more strongly concentrated on the cluster center than the bright nonnucleated dEs. Thompson \\& Gregory ($1993$) showed in the Coma cluster that both dEs and dwarf spheroidals have the same distributions as that of giant early-type galaxies. Morphological classifications of these studies are based on the eye inspection and Thompson \\& Gregory's criteria used for Coma dwarfs are slightly different from those of Sandage and collaborators used for Virgo dwarfs.  One of the reasons why the LS has not been examined systematically is the difficulty in sampling a large number of faint galaxies with known absolute magnitude. Clusters of galaxies are good targets to address the problem of the LS in high density environments. In  particular, in nearby clusters we can sample intrinsically faint galaxies, which are critically important to the study of the LS. However, nearby clusters have such a large apparent size that we cannot survey whole the cluster with a CCD which has a small physical size.  In this study, we present the angular two-point correlation function as well as the radial profiles of galaxies in the Coma cluster on the basis of a large homogeneous sample covering a wide magnitude range $-20 \\leq M_R \\leq -16$, where both giant and dwarf galaxies are included. We examine if there is a difference in these properties between the galaxies which have high central concentration of the surface brightness distribution and those with low central concentration. Our sample is made available by three new techniques; CCD mosaic, semi-automated data reduction/analysis software, and quantitative and objective classification of morphological type of galaxies based on surface photometry parameters.  In section $2$, we briefly explain our imaging observation of the Coma cluster. In section $3$, we describe the data reduction procedures, calibration methods which are special to our camera, and the method of constructing a homogeneous galaxy sample. In section $4$, we describe star-galaxy discrimination, evaluation of the number of contaminated field galaxies, and the method of classifying galaxies according to the surface brightness concentration. We present our results in section $5$. Finally, reliability of our results is discussed in section $6$ in terms of the effects of uncertainties in our classification and contamination by background galaxies. A comparison with previous studies are given in section $7$. ",
        "conclusions": "We have carried out a wide-field galaxy survey in the Coma cluster region with a mosaic CCD camera to study the clustering properties of cluster members. We have investigated the luminosity segregation (LS) quantitatively by measuring the angular two-point correlation function and radial distribution over the magnitude range of $-20 \\leq M_R \\leq -16$, where both giant and dwarf galaxies are included. Our analysis of the galaxy distribution based on the morphology-classified galaxy samples with the unprecedentedly deep limiting magnitude has yielded the following main results:  \\begin{enumerate}  \\item We have found that the galaxies with a high central concentration in surface brightness profile (the HCC type) have strong luminosity segregation while the galaxies with a low central concentration (the LCC type) show almost no luminosity segregation, \\ie the strength of clustering of the LCC-type galaxies does not depend on luminosity.  \\item We have found strong segregation in luminosity for the total sample of Coma cluster galaxies. This is because the majority of the total sample is comprised of the HCC-type galaxies which show strong luminosity segregation.  \\item Brighter HCC-type galaxies tend to more strongly concentrate near the cluster center than fainter HCC-type galaxies, while the LCC-type galaxies do not show such dependence on luminosity in the density profile.  \\end{enumerate}  We have shown that these results are tenable against the contamination by field galaxies and uncertainty in our method of morphological classification. Our results suggest that it would be necessary to consider this type-dependence of the LS in the study of clustering evolution with $\\acf$.  In order to see the universality of our results for clusters of galaxies in general, further studies for other clusters are necessary to investigate the luminosity and morphology dependence of galaxy distribution. It is also desired to assemble a large sample of spectroscopic data for clusters of galaxies to construct samples of confirmed cluster members.  Probably both of a primordial and an environmental effects would have influenced the clustering properties of galaxies in clusters. It is desirable to theoretically evaluate the LS due to each effect quantitatively. In addition, the observational investigation of the LS in high-z clusters would directly distinguish the effects, although morphological classification and field correction would become more difficult in such clusters."
    },
    "9803/astro-ph9803147_arXiv.txt": {
        "abstract": "We report on observations of $^{12}$\\coa emission from the chemically young starburst galaxy Mrk~109.  These observations were part of a study to constrain the O$_2$/CO ratio in metal--deficient galaxies, which were motivated by theoretical work that suggests the possible enhancement of the O$_2$/CO ratio in chemically young systems.  Five low metallicity ($Z \\leq 0.5 Z_{\\odot}$) IRAS galaxies at redshifts $z > 0.02$ (required to shift the 118.75 GHz $^{16}$O$_2$ line away from the atmospheric line) were searched for CO emission.  We detected the CO line in only Mrk~109.  From O$_2$ observations of Mrk~109, we achieved an upper limit for the O$_2$ column density abundance ratio of $N(\\oo)/N(\\co) < 0.31$. These results provide useful constraints for the theoretical models of chemically young galaxies.  We argue that either most of the molecular gas in Mrk~109 does not reside in dark clouds ($A_{V} \\ga 5$), or the standard equilibrium chemistry models are inadequate for metal--poor systems.  The molecular gas mass implied by the CO observations of Mrk~109 is $M({\\rm H}_2) \\simeq 4 \\times 10^{9} \\msun$, and the CO data are consistent with a central starburst induced by the interaction with a nearby companion. ",
        "introduction": "Although studying the molecular gas content of metal--deficient galaxies is challenging, such efforts are essential for our understanding of the formation and evolution of galaxies.  By studying chemically young galaxies in the local universe, we gain insight into the processes which occurred at early times ($z\\sim 1-5$) for the metal--rich spiral and elliptical galaxies found at the current epoch.  Unfortunately, chemically young galaxies tend to be dwarf galaxies and are difficult to detect in CO (Combes 1985; Arnault et al. 1988; Sage et al. 1992; Israel, Tacconi, \\& Baas 1995).  Due to the observational difficulties, very little is currently understood about the molecular gas content of chemically young galaxies.  One of the main uncertainties is the CO to H$_2$ conversion factor.  Wilson (1995) and Arimoto, Sofue, \\& Tsujimoto (1996) have found an empirical relationship indicating an increase in the CO to H$_2$ conversion factor with decreasing metallicity.  However, detailed studies of the molecular clouds in the LMC and SMC suggest other factors, besides metallicity, have an important role on the CO to H$_2$ conversion factor (Rubio 1997; Israel 1997).  Although significant questions remain in our understanding of the CO to H$_2$ factor, the molecular chemistry that occurs in metal--deficient galaxies is even more uncertain.  Molecular chemistry calculations are quite complicated and typically involve networks of hundreds to several thousand reactions (Graedel, Langer, \\& Frerking 1982; Herbst \\& Leung 1989 [HL89]; Langer \\& Graedel 1989 [LG89]; Bergin, Langer, \\& Goldsmith 1995 [BLG95]).  Since the chemistry models still have difficulties in explaining several of the observed abundance ratios in Galactic molecular clouds, very little modeling has been devoted to metal--poor galaxies.  One notable exception is the theoretical study of the chemistry in LMC and SMC molecular clouds by Millar \\& Herbst (1990) [MH90].  In general, the molecular chemistry models are sensitive to a variety of parameters, such as the density, temperature, and ionization field, but the dominant parameter for determining the relative molecular abundances of the carbon and oxygen species is the carbon to oxygen ratio (LG89).  The models predict that the O$_2$/CO ratio decreases exponentially with increases in the C/O ratio (LG89).  This could have interesting ramifications on the molecular abundances in chemically young galaxies.  Observations with the Hubble Space Telescope have shown that the C/O abundance ratio increases with increasing metallicity in {\\sc Hii} galaxies (Garnett et al. 1995).  Therefore, we could expect to find lower C/O ratios and correspondingly larger O$_2$/CO ratios within dark molecular clouds in chemically young galaxies.  Based on these simple ideas, the evolution of the O$_2$/CO abundance ratio as a function of metallicity in galaxies has been recently quantified for a variety of conditions and parameters governing the IMF and star formation histories (Frayer \\& Brown 1997 [FB97]).  At low metallicities, FB97 calculate lower C/O ratios and enhanced O$_2$/CO ratios (O$_2$/CO $\\sim 1$) within dark ($A_{V} > 5$) molecular clouds. At solar metallicities and above, the O$_2$/CO ratio is expected to decrease by several orders of magnitude.  Molecular oxygen has yet to be detected conclusively outside our solar system.  Due to atmospheric attenuation, the ground--based Galactic searches have been limited to observing the rarer $^{16}$O$^{18}$O isotope (Liszt \\& Vanden Bout 1985 [LV85]; Goldsmith et al. 1985; Combes et al.  1991 [C91]; Fuente et al. 1993; Mar\\'{e}chal et al. 1997 [M97]). The most sensitive Galactic studies have provided upper limits of approximately O$_2$/CO $<0.1$.  Extragalactic searches for the redshifted $^{16}$O$_2$ lines have provided more sensitive limits (O$_2$/CO $<0.01$) but have also been unsuccessful (Liszt 1985, Goldsmith \\& Young 1989 [GY89], C91, Liszt 1992 [L92]; Combes \\& Wiklind 1995).  The most sensitive limit to date is O$_2$/CO $<0.002$ ($1\\sigma$) from an absorption line study toward the radio source B0218+357 (Combes, Wiklind, \\& Nakai 1997).  All of these previous searches for O$_2$ have concentrated on chemically rich systems, or molecular clouds of unknown metallicity.  Since the theoretical models suggest the possible enhancement of O$_2$ in chemically young galaxies, we have carried out a search for O$_2$ in metal--deficient galaxies.  From the literature, we have compiled a list of metal--poor ($Z\\leq0.5Z_{\\sun}$) IRAS galaxies with redshifts of $z>0.02$.  There are only about 20 galaxies satisfying these constraints, which were primarily drawn from the samples of Salzer \\& MacAlpine (1988) and Dultzin--Hacyan, Masegosa, \\& Moles (1990).  Unfortunately, most of these galaxies are relatively weak IRAS sources $I(100\\micron) \\la 1$~Jy, and none have reported CO detections.  This is not terribly surprising since metal--deficient galaxies have lower amounts of dust. The metallicities for the galaxies in our sample were derived from their oxygen abundances calculated using standard {\\sc Hii} region analysis techniques whenever the 4363\\AA\\,[{\\sc Oiii}] line was observed (Osterbrock 1989, case B).  For galaxies with no 4363\\AA\\,line, the techniques of McGaugh (1991) were used to estimate the oxygen abundance. In this initial study, we observed the strongest IRAS sources satisfying the metallicity and redshift constraints. ",
        "conclusions": "In this paper, we report the nondetection of \\coa emission in four distant metal--deficient galaxies and the detection of CO emission in Mrk~109.  The \\coa observations of Mrk~109 imply the presence of $4 \\times 10^{9} \\msun$ of molecular gas within the central few kpc of Mrk~109.  The observational data favor a recent interaction of Mrk~109 with a nearby companion which has induced a starburst in the central regions of Mrk~109.  The low metallicity, high gas fraction, and the kinematics are consistent with a young starburst.  Based on the competing photodissociation and chemistry effects, we expect the largest global O$_2$/CO ratios for systems with metallicities in the range of $0.1 < Z/Z_{\\sun} < 0.5$ (FB97).  Although the metallicity of Mrk~109 falls within this optimal range, we fail to detect O$_2$ emission in Mrk~109.  The observed intensity upper limit is $I(\\oo)/I(\\co) < 0.15$.  These observations provide the first constraint on the abundance of O$_2$ in a metal--deficient galaxy.  We derive a column density abundance limit of $N(\\oo)/N(\\co)<0.31$ for Mrk~109. These results suggest that most of the molecular gas in Mrk~109 does not reside in dark clouds, and/or that the gas--phase steady--state chemistry models do not apply for Mrk~109.  It would be challenging to significantly improve the O$_2$/CO limit for chemically young environments derived from this work with current ground--based instrumentation.  Future observations with satellites, such as ODIN (Hjalmarson 1997), of the LMC and SMC should provide more useful limits."
    },
    "9803/astro-ph9803093_arXiv.txt": {
        "abstract": "Geomagnetic  effects  distort  the zenith angle  distribution of  sub--GeV and few--GeV atmospheric  neutrinos, breaking the up--down  symmetry that would be present in the absence of  neutrino oscillations and without a geomagnetic field. The  geomagnetic effects  also produce a  characteristic  azimuthal dependence of the $\\nu$--fluxes,  related to the well known east--west effect,  that should be detectable in neutrino experiments of sufficiently large mass.  We discuss these effects quantitatively. Because the azimuthal dependence is in first order independent of any oscillation effect, it is a useful diagnostic  tool for studying possible systematic effects in the search for neutrino  oscillations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803020_arXiv.txt": {
        "abstract": "We report the first detection of an inverse Compton X-ray emission, spatially correlated with a very steep spectrum radio source (VSSRS), 0038-096, without any detected optical counterpart, in cluster Abell~85. The ROSAT PSPC data and its multiscale wavelet analysis reveal a large scale (linear diameter of the order of 500 $h^{-1}_{50}$ kpc), diffuse X-ray component, in excess to the thermal bremsstrahlung, overlapping an equally large scale VSSRS. The primeval 3 K background photons, scattering off the relativistic electrons can produce the X-rays at the detected level. The inverse Compton flux is estimated to be $(6.5\\pm 0.5)\\times 10^{-13}$ erg~s$^{-1}$~cm$^{-2}$ in the 0.5--2.4~keV X-ray band. A new 327 MHz radio map is presented for the cluster field. The synchrotron emission flux is estimated to be $(6.6\\pm 0.90) \\times 10^{-14}$ erg s$^{-1}$ cm$^{-2}$ in the 10--100 MHz radio band. The positive detection of both radio and X-ray emission from a common ensemble of relativistic electrons leads to an estimate of $(0.95\\pm 0.10) \\times 10^{-6}$ G for the cluster-scale magnetic field strength. The estimated field is free of the `equipartition' conjecture, the distance, and the emission volume. Further, the radiative fluxes and the estimated magnetic field imply the presence of `relic' (radiative lifetime $\\ga 10^{9}$ yr) relativistic electrons with Lorentz factors $\\gamma \\approx$ 700--1700, that would be a significant source of radio emission in the hitherto unexplored frequency range $\\nu \\approx$ 2--10 MHz. ",
        "introduction": "Apart from dark-matter, the diffuse intracluster medium has the two main constituents: the bremsstrahlung emitting hot ($T \\sim 10^{7-8}$ K), tenuous ($n_0\\sim 10^{-(3-4)}$ cm$^{-3}$), and diffuse thermal gas, and the higher energy, relativistic particles (cosmic rays), with electrons emitting the magneto-bremsstrahlung (synchrotron) radiation. The presence of diffuse, large scale (0.2--1.0 Mpc) synchrotron sources with very steep spectra, is known in several clusters (e.g. Coma, Abell 2256, 2319, 85, etc.). The origin of these relativistic particles is currently not well understood but in view of the absence of any optical counterparts, these radio sources are believed to be the remnants (`relics' and `halos') of once active radio galaxies. These remnants are prevented from rapid fading from expansion by the thermal pressure of the surrounding intra cluster gas (Baldwin and Scott 1973). The main energy losses of their relativistic electrons come from synchrotron emission and the inverse Compton scattering of the 3K background radiation. Their presence imply that magnetic fields on similar large scales may exist in the intracluster space (see Feretti \\& Giovannini 1996, and Kronberg 1994 for reviews). That this magnetic field on large scales, may be a general property of clusters of galaxies, is suggested by the Faraday rotation data on radio sources observed in the direction of several clusters (Kim et al. 1991, Kronberg 1994). Estimates of the field strengths, based on Faraday rotation (Kim et al. 1991) or the `equipartition' hypothesis (e.g. Feretti \\& Giovannini 1996), both give a value $B \\sim$ 0.5 to 1.0~$\\mu$G. Currently, the origin and evolution of these fields are not well understood due to observational difficulties in estimating the magnetic fields with sufficient accuracy and to the small number of actual estimates that has been attempted so far for a few clusters only (Kronberg 1994).  Another method of considerable merit to estimate the cluster scale magnetic field is the detection of co-spatial inverse Compton (IC) X-ray emission with the synchrotron emission plasma. The ubiquitous 3K microwave background photons, scattering off the relativistic electrons (the IC/3K process), should produce a diffuse X-ray `glow' associated with the radio plasma (Feenberg \\& Primakoff 1948; Hoyle 1965; Baylis et al. 1967; Felten \\& Morrison 1966; Harris \\& Grindlay 1979; Rephaeli \\& Gruber 1988). The non-relativistic (thermal) analogue of this process is the well known Sunyaev-Zel'dovich effect (Rephaeli 1995). A possible detection of inverse Compton X-ray photons from the scattering of 3K microwave background would not only provide a magnetic field estimate based on a different physical process, but would also provide information on as yet unknown non-thermal X-ray component in galaxy clusters. The inverse Compton method has several advantages: both the magnetic field strength, and the electron energy spectrum are obtainable from the observed IC/3K and the synchrotron emission fluxes, and the field so obtained is independent of the `equipartition' conjecture (Pacholkzyck 1970), the distance, and the emission volume.  Despite a number of attempts over the last 20 years, the detection of IC/3K radiation has proved elusive, very often with only upper limits to the X-ray fluxes and lower limits to magnetic fields (e.g. Harris \\& Romanishin 1974, Rephaeli 1977, Rephaeli et al. 1987, Rephaeli \\& Gruber 1988, Bazzano et al. 1990, Rephaeli et al. 1994, Harris et al. 1995). The principal constraints arise from rarity of diffuse, steep spectrum cluster radio sources, limited sensitivity, and spatial/spectral resolution of low frequency (10--300 MHz) radio or X-ray telescopes, and confusion from cluster thermal emission. However, correlating the radio and the EINSTEIN satellite data, Bagchi (1992) suggested the possibility of IC/3K emission in Abell~85. Recently, using the ROSAT PSPC data, Laurent-Muehleisen et al. (1994) and Feigelson et al. (1995) claimed the detection of IC/3K X-ray emission associated with the radio lobes of the galaxy Fornax-A, further supported by Kaneda et al. (1995), employing the ASCA X-ray spectral data.  This work employs the good spatial resolution, high sensitivity and spectral capabilities of the ROSAT X-ray satellite, to search for co-spatial IC/3K emission from the diffuse radio source 0038-096 located in the rich cD cluster Abell~85. We present evidence for what might be the first detection of the IC/3K emission from the `relic' relativistic electrons residing in the intra-cluster medium, presumably the remnants of an once active radio galaxy that is presently unidentified. We also present a new 327 MHz radio map of the cluster field. Combining the radio and X-ray data, we estimate the magnetic field strength in the co-spatial emission volume. We then briefly discuss its implications. For a cluster redshift of 0.0555 (Pislar et al. 1997), 1 arcmin corresponds to $97 h^{-1}_{50}$ kpc, with  the Hubble constant expressed in units of 50~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "We have presented an evidence for correlated X-ray emission with the diffuse VSSRS in Abell~85. What are the possible emission processes (other than the IC/3K) that can give rise to the excess X-ray emission? The possibility that the emission is either the extension of the synchrotron radiation to the $\\sim$1 keV X-ray band, or that arises from the synchrotron self Compton process (Rees 1967, Harris et al. 1994), can be ruled out from the observed very steep spectral shape for $\\nu \\ge $ 1 GHz, signifying a dearth of necessary high energy photons. Further possibilities that the emission is generated by an active galaxy or that it arises from radio-jet cluster-medium interaction also appear remote due to the relaxed, `halo' type morphology of the VSSRS, and the absence of any optical counterpart.  Recent X-ray imaging and optical spectroscopy of clusters have shown the presence of thermal X-ray emission associated with secondary substructures possibly in the process of gravitational merger (e.g. Briel et al. 1991, Mohr et al. 1993, Henry and Briel 1993, Burns et al. 1994). Significantly, the presence of radio-halo sources appears to be correlated with the evidence for recent mergers, very high gas temperatures (7--14 keV), large velocity dispersions ($\\approx 1300$~km~s$^{-1}$), and the absence of both cooling flows and a single dominant central galaxy (BM type II or III; Edge et al. 1992, Tribble 1993, Feretti and Giovannini 1996). It has been proposed that these data possibly indicate the heating of intracluster gas and the reaccelaration of cosmic ray particles with stochastic amplification of magnetic fields, powered by the energy released in the ongoing mergers (Tribble 1993). Can the excess X-ray emission seen with the VSSRS in A85 come from such a process?  Although imaging data alone can not resolve this question, it is interesting to note that in terms of its physical properties, unlike the other radio-halo host clusters such as the Coma, Abell~2163, 2218, 2255,2256, and 2319, the cluster Abell~85 contains a central dominant cD galaxy, a central cooling flow of about 100 M$_{\\odot}$~yr$^{-1}$, and relatively `cool' gas at a temperature of T $\\sim$ 4 keV (Pislar et al. 1997; Lima Neto et al. 1997).  Durret et al. (1998) have detected a filamentary structure, visible both in optical and in X-rays, linking Abell 85 to the neighbouring cluster Abell 87. They present evidence for a possible merger of matter in this filament with the southern region of Abell 85, in the `south-blob' region. Markevitch et al. (1998) detect an enhancement in the gas temperature in the same region, possibly related to the merger. The position of the VSSRS, however, is not located there as it is shifted about 3 arcmin (300$h_{50}^{-1}$ kpc) to the northwest from the supposed shocked region and has a cooler temperature for the thermal matter in its vicinity.  The resolution of the question of exact physical process behind the excess X-ray emission would require  high spectral (and spatial) resolution imaging spectroscopy data. In the hard X-ray energy regime of $\\approx$ 10--20 keV, the contribution from thermal bremsstrahlung would drop sharply, whereas the non-thermal inverse Compton flux would be visible as an extra component with a steep power law spectrum of spectral index $\\alpha \\approx$ 1. Based on the data presented in this work, we predict the possible IC/3K flux of $\\approx 2.8 \\times 10^{-13}$ erg s$^{-1}$ cm$^{-2}$, and the thermal bremsstrahlung flux of $\\approx 3.0 \\times 10^{-13}$ erg s$^{-1}$ cm$^{-2}$ in the hard X-ray band of 5--10 keV. These fluxes are well within the reach of the currently operative BeppoSAX  telescope and the upcoming AXAF and the XMM missions. The nature of the other excess X-ray and the steep spectrum radio emission detected over the `south-blob' (cf. below) could possibly also be understood with the spectroscopic X-ray data.  With the evidence currently available to us, energetically the most feasible mechanism for the X-ray emission from the VSSRS appears to be the IC/3K process. From theory (e.g. Ginzburg 1989), the (total) energy loss ratio can be shown to be ${f_{IC} \\over f_{S}} \\approx U_{rad}/ {B^{2} \\over {8 \\pi}}$, where $U_{rad} \\approx 5 \\times 10^{-13}$ erg~cm$^{-3}$ is the energy density of the cosmological black-body radiation field, and ${B^{2} \\over {8 \\pi }} \\approx 4 \\times 10^{-14}$ erg~cm$^{-3}$ is the observed magnetic energy density. Therefore, if the observed radiation is from IC/3K process, we expect $f_{IC} / f_{S} \\approx 13$, which is comparable to the observed ratio $f_{IC} / f_{S} \\approx 10$ (the difference is mainly attributed to the finite bandwidths of our data). If the excess X-rays are produced by the thermal bremsstrahlung process, it is difficult to understand why the observed ratio is comparable to the ratio expected if IC/3K were the emission mechanism. This again suggests that the detected X-rays are indeed a product of inverse Compton scattering. This data on ratio of fluxes has enabled us to obtain a model independent magnetic field value, $B = 0.95 \\pm 0.10 \\mu$G, for the diffuse emission volume, located at $\\sim 700 h_{50}^{-1}$ kpc from the cluster centre.  It is apparent (Fig. 2) that another diffuse X-ray excess (the `south-blob') is located at $\\alpha = 00^{\\rm h}~41.8^{\\rm m}$, $\\delta = -09^{\\rm d}~27.5^{\\rm m}$. In the same region, two radio components, without any definite optical counterparts, are observed with the OSRT with $> 20 \\sigma$ signal (162 mJy: northern source, 148 mJy: southern source). A spectral index limit, $\\alpha \\ > 2.3$, is obtained for each of them, based on their non-detection in the new VLA 1.4 GHz sky-survey (Condon et al. in preparation) to the $\\approx$ 5 mJy flux density limit. The nature of this excess X-ray emission is currently not well understood (Lima Neto et al. 1997). Is it possible that a second IC/3K X-ray source is located here, co-spatial with yet another relic radio plasma? If true, this can provide another magnetic field estimate at $\\approx 1 h_{50}^{-1}$ Mpc from the cluster centre. Detailed, low-frequency ($<327$ MHz) radio and X-ray spectral data are necessary to explore this possibility.  Finally, the estimated magnetic field and the observed radiative fluxes from the VSSRS imply the presence of relativistic electrons with Lorentz factors $\\gamma \\approx$ 700--1700 (for 0.5--2.4 keV band, $\\nu_{x} \\propto \\nu_{bg} \\gamma^{2} $). The population of electrons in this energy range would in turn produce significant radio emission ($\\sim $ 1400 Jy at 2 MHz) in the presently unexplored frequency range $\\nu_{syn} \\approx $ 2--10 MHz ($\\nu_{syn} \\propto B\\;\\gamma^{2}$). The radiative lifetime of such `relic' electrons is $\\ga 10^{9}$ yr (Harris \\& Grindlay 1979), comparable to the time-scale for evolution of clusters. Although this frequency range is below the ionospheric cutoff, the future generation radio telescopes observing from above the earth's atmosphere or from the far side of the moon (Burns 1990), could make such measurements. Such low frequency radio data in association with sensitive X-ray data would prove invaluable in probing the physics of the intracluster media in large scale structures.  It is encouraging to note that the inverse Compton X-ray technique has the great potential, not only in measuring the elusive intra-cluster magnetic fields, but also in probing the hitherto unexplored population of very old relativistic electrons  residing in the diffuse intra-cluster space."
    },
    "9803/astro-ph9803216_arXiv.txt": {
        "abstract": "\\d\\/ in NGC 1536 is possibly the least luminous and energetic Type II supernova discovered to date. The entire light curve is subluminous, never reaching $M_V = -14.65$. The radioactive tail follows the \\co\\ decay slope. In the case of nearly complete trapping of the $\\gamma$-rays, the \\ni\\ mass derived from the tail brightness is extremely small, $\\sim 0.002$ \\M. At discovery the spectra showed a red continuum and line velocities of the order of 1000 \\kms. The luminosity and the photospheric expansion velocity suggest that the explosion occurred about 50 days before discovery, and that a plateau probably followed. Model light curves and spectra of the explosion of a 26 \\M\\ star successfully fit the observations. Low mass models are inconsistent with the observations. The radius of the progenitor, constrained by the prediscovery upper limits, is \\r0\\ \\ltsim 300 \\R. A low explosion energy of $\\sim 4 \\times 10^{50}$ ergs is then required in the modeling. The strong \\ion{Ba}{2} lines in the photospheric spectra are reproduced with a solar abundance and low $T_{\\rm eff}$. A scenario in which the low \\ni\\/ mass observed in \\d\\/ is due to fall--back of material onto the collapsed remnant of the explosion of a 25--40 \\M\\/ star appears to be favored over the case of the explosion of an 8--10 \\M\\ star with low \\ni\\/ production. ",
        "introduction": "\\d\\/ was serendipitously discovered on Jan. 14.15 U.T. (\\cite{demello}) in NGC~1536, a morphologically disturbed spiral galaxy belonging to a high density group (\\cite{maia}).  Although NGC~1536 ($v_{\\rm helio}=1296$ \\kms, RC3) is one of the galaxies patrolled visually by Rev. Evans for his SN search (private communication), he missed the SN because its brightness apparently never exceeded his detection limit.  The first spectrum showed the main features of type II SNe, but also revealed peculiarities which make this object unique. Particularly noteworthy were the extremely slow expansion velocities, the red color and the strong \\ion{Ba}{2} lines (\\cite{tur98}). A campaign of photometric and spectroscopic observations was therefore promptly started at ESO and CTIO. The complete data set will be presented elsewhere. ",
        "conclusions": "\\d\\/ is among the least luminous Type II supernovae known to date. The low luminosity persists also in the tail, where the decline rate is close to that of a \\co-powered tail.  \\d\\ is also characterized by a low temperature and a very low expansion velocity. These features are well reproduced by an explosion model with a relatively small progenitor radius (300 \\R), massive ejecta ($\\sim$ 24 \\M), low explosion energy ($4 \\times 10^{50}$ ergs), and small ejected \\ni\\/ mass ($\\sim$ 0.002 \\M).  A similarly small amount of \\ni\\ has recently been proposed for another SN~II, SN~1994W (\\cite{soll}), but this object was rather luminous and its tail declined faster than the \\co\\ slope.  The neutrino-heating mechanism of massive star explosions is not well understood, so the explosion energy and the amount of \\ni\\ ejected as a function of progenitor mass are difficult to predict from the hydrodynamical models. According to pre-supernova model calculations, stars more massive than $\\sim$ 25 \\M\\ form an Fe core in a greater gravitational potential because of the significantly smaller C/O ratio and the consequently weak carbon shell burning (\\cite{nom93}; \\cite{woowae}). Therefore, if the efficiency of neutrino heating does not change with exploding mass, more massive stars tend to produce lower explosion energies, and they suffer from more fall-back of material onto the collapsed remnant. The mass cut is then placed further out. However, the progenitor must be less massive than the Wolf-Rayet progenitors (i.e., \\ltsim 30 - 40 \\M) because of the presence of an envelope with a mass of at least 15 \\M\\ above the photosphere at day 50.  Stars in the range $\\sim$ 10 - 25 \\M\\ are likely to explode with typical explosion energies of $\\sim 1 \\times 10^{51}$ ergs, and to produce 0.07 - 0.15 \\M\\ of \\ni, as indicated by the brightness and light curve shapes of SN 1993J, SN 1987A and of Type Ibc supernovae (\\cite{nom95}; \\cite{young}). Stars of 8 - 10 \\M\\/ are predicted to produce little \\ni\\ and other heavy elements (\\cite{nom84}; \\cite{maywil}). However, these stars are at the top of the AGB in the pre-supernova stages (\\cite{hashi}), which is inconsistent with the small radius of \\d. Also, synthetic spectra obtained from the low mass models have shallow absorption lines compared to the observations.  In conclusion, the progenitor of \\d\\ was probably as massive as 25 - 40 \\M.  Our 26 \\M\\ model is in this mass range, thus being consistent with the low $E/M$ and the small \\ni\\ mass. The presupernova radius depends on many parameters, and the reason for its being rather small in \\d\\/ is an open question.  Spectral models appear to rule out the case of low metallicity and He enhancement. A possible scenario involves a close binary system in which the companion star spiraled in to make the envelope more massive (\\cite{pod}).  The ejection of a small mass of \\ni\\ provides a constraint on the mass of the collapsed remnant, $M_{\\rm rem}$ (\\cite{nom93}; \\cite{thiel}). In our 26 \\M\\ model, is $M_{\\rm rem} = 1.8$ \\M. If $M_{\\rm rem}$ were assumed to be smaller, more \\ni\\ would be ejected: e.g. if $M_{\\rm rem} = 1.7$ \\M, then M(\\ni) = 0.1 \\M. For more massive progenitor models, $M_{\\rm rem}$ would be larger than 1.8 \\M. If the equation of state of nuclear matter is relatively soft, the maximum (gravitational) mass of a cold neutron star is $\\sim$ 1.5 \\M\\ (\\cite{bb94}). If this is the case, the remnant of \\d\\ must be a small-mass black hole, and the low luminosity tail of \\d\\ might be a signature of the formation of such an object.  \\bigskip We would like to thank Piero Rosati for observing \\d\\/ and Rev. Robert Evans for providing his pre-discovery limits.  P.A. Mazzali acknowledges receipt of a Foreign Research Fellowship at N.A.O., and is grateful to T. Kajino and the Dept. of Astronomy at the University of Tokyo for the hospitality.  A. Piemonte aknowledges financial support from ESO during his stay in Chile.  This work has been supported in part by the Grant-in-Aid for Scientific Research (05242102, 06233101) and COE research (07CE2002) of the Ministry of Education, Science, and Culture in Japan, and by the National Science Foundation in US under Grant No. PHY94-07194.  Part of the observations have been obtained at the ESO 1.5m telescope operated under the agreement between ESO and Observatorio Nacional -- Brasil.   \\noindent"
    },
    "9803/astro-ph9803166_arXiv.txt": {
        "abstract": " ",
        "introduction": "The heavy elements present in the hot gas in galaxies are thought to be supplied both by stellar winds and by supernovae (SNe). Theoretical calculations of the nucleosynthesis of the SNe have been improved greatly in the past ten years (e.g., Thielemann, Nomoto, \\& Hashimoto 1990; Thielemann, Nomoto, \\& Hashimoto 1996, hereafter TNH).  In contrast, the observational evidences of metal rich gas have been obtained for only a limited number of young supernova remnants (SNRs).  Ku et al. (1984) first performed spatially-resolved spectroscopy on the Cygnus Loop with the Einstein observatory. They constructed an X-ray image with spatial resolution of $\\sim$1$^\\prime$ in the energy region of 0.1--4.0 keV. The X-ray image clearly showed limb-brightening structure.  They noticed that kTe at the limb was generally lower than that at the center.  Charles, Kahn, \\& McKee (1985) found that kTe gradually increased toward the center from the vicinity of the limb.  Tsunemi et al. (1988) observed the whole Cygnus Loop with the Gas Scintillation Proportional Counters on board the Tenma satellite, which possessed roughly twice better energy resolution than that of the proportional counter (Koyama et al. 1984). They detected emission lines from highly ionized Si and S. They fitted the X-ray spectrum in combination with the spectrum obtained with a previous rocket observation (Inoue et al. 1979, 1980). They found that a two-component NEI model could reproduce the X-ray spectra.  Together with Charles et al. (1985), Tsunemi et al. (1988) suggested that tenuous high-kTe plasma was sitting in the interior region which was surrounded by dense low-kTe plasma. Such high-kTe component was confirmed by Hatsukade \\& Tsunemi (1990) with the Ginga satellite.  Due to its high sensitivity, high energy resolution, and wide energy band, the ASCA observatory has opened a new window for the physics of hot plasmas like SNRs. We can perform the plasma diagnostics in detail by using the X-ray CCD cameras.  So far, the abundances of heavy elements have been well determined for many young SNRs with the ASCA observatory (see review by Tsunemi \\& Miyata 1997).  For evolved SNRs, it is difficult to determine the abundances due to low-kTe, low surface brightness, and large interstellar absorption. Furthermore, the swept up interstellar medium (ISM) dominates the ejecta mass from the progenitor star. This makes it difficult to detect the ejecta.  The Cygnus Loop is an evolved SNR and one of the best studied SNRs. Since the Cygnus Loop is roughly 8 degrees away from the Galactic plane, the neutral hydrogen column (${\\rm N_H}$) is only a few$\\times {10}^{20}\\ {\\rm cm}^{-2}$. This enables us to detect emission lines from O and to determine kTe with high accuracy. The mass of the swept up ISM is estimated to be roughly 100 ${\\rm M_{\\odot}}$ since the average density of the ISM is 0.2 ${\\rm cm}^{-3}$ and the mean radius is 18.8 pc (Ku et al.  1984). In comparison, the mass of the ejecta ranges from one to several tens ${\\rm M_{\\odot}}$. Thus, the ejecta are submerged under the sea of the ISM.  However, due to its high surface brightness and large apparent size, we can perform spatially-resolved analysis for the Cygnus Loop. This enables us to search ejecta in detail if ejecta have not yet mixed well with the ISM.  The first observation of the Cygnus Loop with the ASCA observatory was performed by Miyata et al. (1994; hereafter MTPK) on the NE limb. They found that kTe increased toward the center whereas log\\hspace{0.2mm}($\\tau$) decreased. They determined the abundances of heavy elements from O to Fe and found that metals were deficient at the NE limb.  In this {\\it paper}, we present the observational results of the center portion of the Cygnus Loop observed with the ASCA observatory. The X-ray spectrum obtained from the center portion was quite different from that obtained at the NE limb.  We studied the plasma diagnostics of the interior of the Cygnus Loop and nucleosynthesis at the SN, apart from dense shell regions. ",
        "conclusions": "We observed the center portion of the Cygnus Loop supernova remnant with the X-ray CCD cameras on board the ASCA observatory.  We confirmed a significant departure from a collisional ionization equilibrium condition at the center portion. Obtained abundances at the core region were larger than those of cosmic abundances for Si, S. Moreover, abundances of Si, S, and Fe were $\\sim$40 times larger than those obtained at the NE limb. This strongly supports the hypothesis that the X-ray emitting plasma in the core is ejecta in origin. Although previous X-ray observations could not detect any signature of ejecta, we could find it by using both the imaging capability and the high energy resolving power of the ASCA observatory.  Obtained abundances can be compared with theoretical calculations of nucleosynthesis by a SN explosion. Since we only observed the core region, we integrated the model calculations from the explosion center to some mass radius. For the explosion models for type Ia SN or for type II with the progenitor mass less than 20 ${\\rm M_{\\odot}}$, the abundance of Fe is much larger than that of Si or S even if we integrate to the outermost radius.  On the other hand, the abundances of Si, S, and Fe can be explained with the type II model with 25${\\rm M_{\\odot}}$ if we integrate these abundances from the mass cut radius 1.77${\\rm M_{\\odot}}$ to $2 \\sim 3.9 {\\rm M_{\\odot}}$.  Therefore, our results suggest the massive progenitor origin of the Cygnus Loop and the existence of the stellar remnant associated with the Cygnus Loop. \\par  \\vspace{1pc}\\par We would like to acknowledge Prof. Nomoto and Drs. T. Suzuki and K. Iwamoto for fruitful discussions about the theoretical aspects of the SNRs at conference of Thermonuclear Supernovae.  Dr. B. Aschenbach kindly gave us the whole X-ray image of the Cygnus Loop obtained with the ROSAT all-sky survey.  We thank Dr. Thielemann to use his data on nucleosynthesis of SNe.  We would like to thank the anonymous referee for her or his detailed comments and suggestions which greatly improved this paper.  We are greatful to all the members of ASCA team for their contributions to the fabrication of the apparatus, the operation of ASCA, and the data acquisition. We thank the members of {\\tt ASCA\\_ANL/Sim ASCA} software team."
    },
    "9803/astro-ph9803344_arXiv.txt": {
        "abstract": "We present X-ray and optical data of the new\\-ly discovered AM Her variable \\rxj18.  This X-ray source was observed in the \\ros\\ All-Sky-Survey in September 1990 and subsequently discovered as a highly variable and soft X-ray source.  Follow-up pointed observations confirmed the strong variability and revealed a periodic flux modulation including a sharp and nearly complete eclipse of the X-ray emission.  Based on the shape and duration of this eclipse as well as the lack of optical eclipses we favour an interpretation in terms of a self-eclipse by the accretion stream. The X-ray spectrum averaged over the full period is dominated by soft emission below 0.5 keV. The ratio of  soft (blackbody) to hard (bremsstrahlung) bolometric energy flux is 1670, distinguishing this object as another example of a polar with an extreme strong excess in soft X-rays. ",
        "introduction": "AM Her type variables are accreting binaries in which the strong magnetic field of the white dwarf controls the geometry of the material flow coming from the late-type companion star. The radial inflow of matter produces a shock front above the white dwarf surface which gives rise to hard X-rays and polarized cyclotron radiation in the IR to UV range. In the simplest physical model (Lamb 1985) half of this shock radiation intercepts the white dwarf surface and is reradiated as thermal emission from the heated accretion spot in the UV and soft X-ray band.  X-ray observations, most notably with the \\ros\\ satellite, had and have a great impact on the study of magnetic cataclysmic variables. The rather high intensity, strongly variable and soft X-ray emission allows to set up selection criteria (e.g. using the \\ros\\ data base) with a high detection/identification rate which reflects in the fact that about 80\\% of all magnetic cataclysmic variables have been discovered at X-ray wavelengths (Beuermann 1998).  During the search for supersoft X-ray sources in the \\ros\\ (Tr\\\"umper 1983) All-Sky-Survey (RASS), we have found a bright X-ray source showing strong scan-to-scan variability in the X-ray count rate as measured in the PSPC.  Subsequent spectroscopic observations re\\-vealed a 15th magnitude cataclysmic variable located near the X-ray position.  Additional pointed \\ros\\ PSPC observations strengthened the evidence of a strongly modulated X-ray light curve as well as of a very soft X-ray spectrum. These combined characteristics, and particularly the supersoft X-ray spectrum, are typical signatures of the AM Her subclass (Beuermann \\& Schwope 1994).  Here, we report on the results of our extensive \\ros\\ and optical observations performed to study the accretion geometry in \\rxj18\\ in more detail. The identification as an AM Her system (polar) and some first details have been reported in June 1994 at the Abano-Padova conference on Cataclysmic Variables (Greiner \\etal\\ 1995a) after which the object got the variable star name V884 Her. The results of quick follow-up observations of \\rxj18\\, (= WGA J1802.1+1804) have also been communicated in Singh \\etal\\, 1995 (note the flux error; see Greiner \\etal\\ 1995b) and Szkody \\etal\\ (1995). Recently, Shrader \\etal\\, (1997) reported the results of IUE observations. Due to its soft X-ray spectrum \\rxj18\\ has also been detected with the EUVE satellite (Lampton \\etal\\ 1997; = EUVE J1802+18.0). ",
        "conclusions": "\\subsection{AM Her type classification}  The coincidence of periods derived from the X-ray and optical data and the stability of the eclipse feature over more than 5 years suggest that the binary system is synchronized. The evidence of an X-ray eclipse without a corresponding optical eclipse together with the shape and variations with time of the X-ray eclipse imply that (1) the optical and X-ray radiation originate from different emission regions and (2) the X-ray emitting region is confined to a small region near the white dwarf.  The localized X-ray emitting region strongly indicates a magnetic CV subclass, since the X-ray emission from an inner accretion disk would only show eclipses when the disk is occulted by the companion star -- and such eclipses are evident in both X-ray and optical light.   The supersoft X-ray spectrum further suggests an AM Her / `polar' subtype (rather than `intermediate polar' type), since the association between AM Her objects and luminous emission in soft X-rays is well established.  We note, however, that the discovery of three intermediate polars with strong soft X-ray emission has changed this clear association (Haberl and Motch 1995).  However, the consistency of singular optical and X-ray periods in the case of \\rxj18 indicates the likelihood of spin / orbital synchronous rotation, another classic property of AM Her binaries. Finally, the large equivalent widths of the optical emission lines also suggest a polar nature.  An independent confirmation of the polar classification comes from polarimetric observations (Szkody \\etal\\ 1995). The circular polarization in the 5700--7700 \\AA\\ range reaches a maximum of 4\\% in two peaks separated by a minimum which occurs at the time of photometric minimum.  \\subsection{Geometric configuration}  The relatively short time scales observed for X-ray eclipse ingress, egress, and duration argue against the interpretation of eclipses as the disappearance of the X-ray emitting region behind the limb of the rotating white dwarf. We therefore interpret the eclipse as a self-eclipse of the X-ray emitting region by the accretion stream. X-ray spectral evidence for increased absorption column associated with the eclipse and the variable eclipse length provide further support for this conclusion.  Although one would expect the optical light curve to be more complex than the X-ray light curve due to several reasons (cyclotron radiation is beam-modulated and can be self-eclipsing; the recombination component is modulated by projection of the partially optically thick stream; additional light may be seen from the illuminated side of the secondary), the observed optical light curve is surprisingly smooth. Despite the stream-eclipse of the accretion spot in X-rays, the optical light is not affected due to its probable origin from a more extended region, including a portion of the accretion column upstream from the X-ray emitting area. The lack of optical eclipses suggests an inclination of $i < 78$\\degr\\ (assuming a typical M5/6 companion).   During mid eclipse, there is still detectable emission not only in the total flux, but also in the soft X-ray component. Thus, it seems possible that the accretion spot is not fully eclipsed or we see another, uneclipsed emission component.  The presence of X-ray emission at all phases implies that $i + \\beta <$ 90\\degr\\ (with $\\beta$ being the colatitude of the accreting magnetic pole). Combining this with the condition of stream eclipse ($i > \\beta$) we derive the following constraints on $i$ and $\\beta$:  $$ 45\\degr < i < 78\\degr\\ \\hspace{0.5cm} {\\rm and} \\hspace{0.5cm} 0\\degr < \\beta < 45\\degr $$  The X-ray light curve is very complex, and there are large differences seen when comparing observations well separated in time.  This might suggest that the density and/or the size of the accretion stream vary with time and thus cause changes in the size and/or vertical extent of the accretion spot. In this regard, we note that the eclipse duration is not constant.  Also, there are large changes in the characteristics of a secondary minimum, as noted in Section 2.1.2 and shown in Fig. \\ref{foldlc3}. Furthermore, the eclipse length seems to be correlated with the intensity of the X-ray emission before and after the eclipse. During the Oct. 1992 observation (cycle --4269 in Fig. \\ref{unfoldlc}) with its X-ray low-state, the eclipse length is only 0.07 phase units, whereas it is 0.1 during the high-state in Sep. 1993 and even 0.12 in the last eclipse observed.  A higher X-ray intensity may be caused by an increase in the rate of mass transfer, perhaps leading to a larger width of the accretion stream, which then increases the duration of the X-ray eclipse.  A significant larger ingress (5--7 min.) than egress (2--2.5 min.) is measured. Such asymmetry is possible if the impact area is not a circle but an arc due to the coupling of the stream to nonpolar field lines (see e.g. Beuermann \\etal\\ 1987 and their notion of ``X-ray auroral oval''. Alternatively, the accretion stream can be imagined to impact the white dwarf surface not from a perpendicular direction thus forming an elongated footprint on the white dwarf surface. The slow fall of X-ray intensity into the eclipse minimum then can be understood as being due to two effects: first, the projected area of the footprint decreases due to the rotation of the white dwarf and secondly, the stream starts to occult the footprint area. This is consistent with the observed X-ray spectral property that the absorbing column does not change during the start of the slow fall into eclipse minimum (phase interval 0.0--0.1, see Fig. \\ref{foldlc}), but only during the eclipse (phase interval 0.1--0.2). Such behaviour is also consistent with the interpretation of the variation of the polarization being due to varying aspects of an emission region extended in longitude (Szkody \\etal\\ 1995). A size estimate of this elongated emitting region beyond an axis ratio of about 2--3:1 is hard to evaluate since the relative sizes of the emitting region and the stream are important.  The X-ray eclipse occurs about 35\\degr\\ after the blue-to-red crossing of the radial velocity curve. In our interpretation of a stream-eclipsing geometry this is consistent with the picture that the emission lines are produced in the ballistic part of the accretion stream. We note, however, that the radial velocity amplitude is lower than one would expect as the maximum free fall velocity.  The best fit blackbody temperature and the normalization transform into a corresponding blackbody radius of the emission region of 750 (D/100 pc) km. Bearing in mind the above mentioned elongation this is understood as the radius of a circle which has the same area as the elongated emitting region. If true, this implies a stream diameter near the white dwarf surface of the order of 850--1050 km.  Besides the herewith presented \\rxj18\\ several other polars have been found where the accretion stream was identified as the cause of narrow absorption dips in the EUV/X-ray light curves: EF Eri (Patterson \\etal\\ 1981, Beuermann \\etal\\ 1991), UZ For (Warren, Vallerga and Sirk 1995), EK UMa (Clayton \\& Osborne 1994), QS Tel (Beuermann \\& Thomas 1993, Buckley \\etal\\ 1993, Schwope \\etal\\ 1995), HU Aqr (Schwope \\etal\\ 1993, Hakala \\etal\\ 1993, Schwope \\etal\\ 1997), QQ Vul (Beardmore \\etal\\ 1995) and V2301 Oph (= 1H1752+081) (Hessman \\etal\\ 1997). Another stream-eclipsing polar may be AX\\,J2315--592 which has recently been found from ASCA observations (Misaki \\etal\\ 1996, Thomas and Reinsch 1996). The EUVE light curves of UZ For show that both the phase and the amplitude of the dips vary over the 3-day observation suggesting variations in the density and position of the accretion stream similar to what we observed in the \\ros\\ light curves of \\rxj18. Also, the X-ray light curve of QQ Vul as observed with ROSAT is surprisingly similar to that of \\rxj18\\ (as given in Fig. \\ref{foldlc}).  \\subsection{Soft X-ray excess}  As in several other AM Her binaries, the soft X-ray luminosity of \\rxj18 is larger than the hard one, a fact known as `soft X-ray problem' (Rothschild \\etal\\ 1981). Kujpers \\& Pringle (1982) proposed that non-stationary accretion of dense blobs can heat the photosphere from below. This implies that high magnetic field systems should have a weak bremsstrahlung component, which is proved by a comparison of soft-to-hard flux ratios with measured magnetic field strengths in different AM Her systems (Beuermann \\& Schwope 1994). According to this correlation a magnetic field larger than 40 MG is suggested for \\rxj18\\ which has one of the highest ratios of soft/hard emission among polars. This would produce a very impressive Zeeman split in the H absorption lines coming from the white dwarf as well as strong cyclotron humps (depending on the viewing angle). However, since we observed the system ``only'' in its robust state of accretion, the optical spectrum is strongly dominated by the accretion stream making the determination of the magnetic field impossible."
    },
    "9803/gr-qc9803073_arXiv.txt": {
        "abstract": "In the hyperbolic slicing of de Sitter space appropriate for open universe models, a curvature scale is present and supercurvature fluctuations are possible. In some cases, the expansion of a scalar field in the Bunch-Davies vacuum includes supercurvature modes, as shown by Sasaki, Tanaka and Yamamoto. We express the normalizable vacuum supercurvature modes for a massless scalar field in terms of the basis modes for the spatially-flat slicing of de Sitter space. ",
        "introduction": "Scalar fields in de Sitter spacetime have long provided a testing ground for issues of quantum field theory in curved spacetime \\cite{birdavies,sfulling}.  Further motivation for their study stems from the central role they play in inflationary cosmology \\cite{KolbTurner}. Several different coordinate systems can be used to cover de Sitter space, and subtleties in the quantization of fields can arise in some of the less familiar coordinate systems.  These subtleties have been highlighted by recent models of open inflation \\cite{bgt,openi}, in which two periods of inflation are separated by nucleation of a bubble.  The bubble interior includes an open universe ($\\Omega_0 <1$), where we could be living today, described by hyperbolic, spatially-curved coordinates \\cite{colegott}.  A key difference between these hyperbolic coordinates and the more familiar spatially-flat slicing of de Sitter space is the presence of a curvature scale.  This in turn leads to the possibility of supercurvature \\cite{lythwos} fluctuations, fluctuations with wavelength longer than the curvature scale.  Unlike the continuum of modes familiar from the spatially-flat slicing of de Sitter space, a normalizable supercurvature mode may exist for an isolated, discrete eigenvalue of the spatial Laplacian, or not at all.  Although such modes have no analogue in the spatially-flat slicing of de Sitter space, it has been shown \\cite{sasaki95} (see also \\cite{Mosch}) that the supercurvature modes must be included in the vacuum spectra of low-mass scalar fields in order to produce a complete set of states, and hence the proper Wightman function.  For massless, minimally-coupled scalar fields in de Sitter space, there is in addition a well-known infrared divergence in the (coordinate-independent) Wightman function.  The infrared divergence is related to a dynamical zero mode in the spectrum of a massless field \\cite{desinfra,kg93}.  Kirsten and Garriga \\cite{kg93} covariantly quantized this zero mode in a spatially-closed slicing of de Sitter space. Extending their result from these closed coordinates to the coordinate system appropriate to open inflation has not yet been done. Here we identify the zero mode as one of the supercurvature modes when quantizing the minimally-coupled massless field in the open hyperbolic coordinates.  It has already been noted that supercurvature modes can make significant contributions to the fluctuations in the cosmic microwave background (CMB) radiation, and several of their effects have been calculated for models of open inflation \\cite{sccmb,garriga,sasaki96,gar-97}.    (The massless field zero-mode subtleties, except for a variant studied in \\cite{gar-97}, do not arise for these specific CMB calculations, which are sensitive to higher multipoles.) In addition to contributing to observable density fluctuations, such long wavelength, supercurvature modes might play a role \\cite{dkpre} at the end of inflation, when recent advances \\cite{newreh} in the theory of reheating are taken into consideration.   \\indent  In summary, increased understanding of these supercurvature modes is motivated both by general questions of quantizing fields in curved backgrounds, and by recent inflationary model-building.   We will focus here on the example of supercurvature modes for a massless, minimally-coupled scalar field, expressing them as a sum over the basis modes for a spatially-flat slicing of de Sitter space.  This overlap gives a measure of \\lq\\lq where all the supercurvature modes go\" in the familiar flat-space spectrum of such fields. The massless case is chosen for tractability, and questions about the zero mode are postponed for future work \\cite{jdcdk}.  Throughout this paper, we consider only an unperturbed de Sitter metric; a more complete treatment would include study of the backreaction of such fields on the background metric.  Because the supercurvature modes stretch beyond the horizon, any such study of the coupled metric fluctuations would need to pay special attention to the gauge subtleties which always accompany superhorizon fluctuations \\cite{rhb}, and such issues are not pursued here.  In section II, the two covers of de Sitter space (including the flat and hyperbolic slicings) are given, and the field quantization pertinent to open inflation in both systems is reviewed.  As these supercurvature modes have no analogue in the usual flat slicing of de Sitter space, this section gathers some previous work on supercurvature modes and provides notation and context for the rest of the paper. Section III specializes to the massless case.  The explicit calculation of the overlap for some of these supercurvature modes and the more familiar flat space modes is given, and indicates a general form for the overlap between all the massless supercurvature and spatially flat modes. We verify this general form by integrating over the spatially flat modes, weighted by the overlap, to obtain the original supercurvature modes. Concluding remarks follow in Section IV.  Three Appendices include the supercurvature mode normalization at fixed time in the flat coordinates, and details of the overlap calculation along different hypersurfaces within de Sitter space. ",
        "conclusions": "In conclusion, we have given the explicit form for the overlap between the flat basis functions (\\ref{flatF}) and the massless supercurvature modes.  As a result, the supercurvature modes within that patch of de Sitter space which would contain our open observable universe can be written \\eq \\uel (x_r(x_f)) = -i \\sqrt{2\\ell (\\ell + 1)} \\int_0^\\infty dk \\> k^{-1/2} \\> j_\\ell (k) \\> \\phi_{k\\ell m} (x_f) . \\en The long-wavelength supercurvature modes are distributed over the spatially-flat basis modes, oscillating over flat-space comoving wavenumber $k$ with decreasing amplitude. More quantitatively, the spherical bessel functions $j_\\ell(k)$ have their first and largest maximum (\\cite{Absteg} 10.1.59) near $k \\sim (\\ell + \\frac{1}{2}) [1 + O(\\ell^{- 2/3})]$ with the approximation improving as $\\ell$ increases, but the damping envelope going as $k^{-1/2}$ lowers this peak for higher $\\ell$. It may be possible to use the description \\cite{gar-97} of $M^2 \\ne 0$ supercurvature modes as small perturbations of the massless supercurvature modes to extend the above to small $M^2$.  This expression for the supercurvature mode on fixed $\\eta$ surfaces, extending into region $R$, may be useful for understanding the effects of supercurvature modes during reheating. Unlike the event of bubble nucleation, reheating occurs in the future of region $C$ and hence descriptions using fixed time $r_c$ surfaces are not appropriate.  Rather, it is important to understand the dynamics of these modes within region $R$, corresponding to our observable, open universe.  Having an expression for the normalizable supercurvature modes on slicings extending into region $R$ is a step in separating long wavelength properties of these modes from issues related to their non-normalizability at fixed time $t_r$ in region $R$.  We showed by comparing Cauchy and non-Cauchy surface calculations that in some cases non-Cauchy surface calculations of norms (in the appendix) and overlaps (in the text) agree for supercurvature modes, up to an identifiable boundary term. These non-Cauchy surfaces (fixed time in the flat coordinates)  extend into the open universe and thus could be used (with caution) to calculate other properties. The complementary calculations along different surfaces were required to verify that the non-Cauchy surface representation was indeed correct.   In addition, we identify one specific supercurvature mode as responsible for the well-known infrared divergence for massless scalar fields in de Sitter space.  Not only does this mode have divergent norm; it is demonstrated here to be a dynamical zero-mode.  This identification is a necessary first step toward its eventual replacement with an appropriate collective coordinate, similar to what has been done in closed coordinates\\cite{kg93}.  Something similar has been done in \\cite{gar-97} in the context of two field models and quasi-open inflation, here it is found more generally as a property of the massless supercurvature modes."
    },
    "9803/astro-ph9803034_arXiv.txt": {
        "abstract": "It is shown that the new observed redshift distributions of various flux-limited samples of radio sources in general are consistent with the predictions of two basic evolutionary models published by Condon (1984) and Dunlop \\& Peacock (1990), i.e. none of them can be rejected at the confidence level of about 95 per cent. However, the models allowing a free-form evolution and suggesting both density and luminosity evolution are more consistent with the observational data \\underline{at lower redshifts}, while the 'pure luminosity evolution' model fits better the data \\underline{at higher redshifts}. This leads to a suspicion that the {\\it same evolution} governing {\\it all} radio sources, suggested by Condon (1984), might not be the case. ",
        "introduction": " ",
        "conclusions": ""
    },
    "9803/astro-ph9803278_arXiv.txt": {
        "abstract": "We present {\\it Rossi X-ray Timing Explorer}\\ (\\RXTE) All-Sky Monitor observations of the X-ray binary Circinus~X-1 which illustrate the variety of intensity profiles associated with the 16.55~d flaring cycle of the source. We also present eight observations of Cir~X-1 made with the \\RXTE\\ Proportional Counter Array over the course of a cycle wherein the average intensity of the flaring state decreased gradually over $\\sim$12 days. Fourier power density spectra for these observations show a narrow quasi-periodic oscillation (QPO) peak which shifts in frequency between 6.8~Hz and 32~Hz, as well as a broad QPO peak that remains roughly stationary at $\\sim$4~Hz. We identify these as Z-source horizontal and normal branch oscillations (HBOs/NBOs) respectively. Color-color and hardness-intensity diagrams (CDs/HIDs) show curvilinear tracks for each of the observations. The properties of the QPOs and very low frequency noise allow us to identify segments of these tracks with Z-source horizontal, normal, and flaring branches which shift location in the CDs and HIDs over the course of the 16.55~d cycle. These results contradict a previous prediction, based on the hypothesis that Cir~X-1 is a high-$\\dot{M}$ atoll source, that HBOs should never occur in this source (\\cite{oosterbroek95}; \\cite{klis94}). ",
        "introduction": "The X-ray binary Circinus~X-1 is unique in its complex temporal and spectral variability. A 16.55~day cycle of flaring is observed in the X-ray (\\cite{kaluzienski76}) as well as optical (\\cite{moneti92}), IR (\\cite{glass78}), and radio bands (\\cite{whelan77}). The high degree of stability of the period of this cycle is evidence that it is the orbital period.  The onset of flaring has been suggested to be the result of enhanced mass transfer occurring near periastron of a highly eccentric binary orbit (\\cite{murdin80}; \\cite{oosterbroek95}).  The X-ray profile and average intensity of the 16.55~day cycle has varied considerably over timescales of years (see e.g, Dower, Bradt, \\& Morgan 1982\\nocite{dower82}; \\cite{stewart91}; \\cite{oosterbroek95}; Shirey et~al.\\ 1996, hereafter \\cite{shirey96}). Observations with the {\\it Rossi X-ray Timing Explorer} (\\RXTE) All-Sky Monitor (ASM, 2--12~keV) showed Cir~X-1 in a sustained bright state with a baseline intensity level of $\\sim$1.0~Crab (75 c/s; 1060 $\\mu$Jy at 5.2 keV) and strong flaring up to as high as 3.5~Crab (\\cite{shirey96}). The flaring state began during the day following phase zero (based on the radio ephemeris of \\cite{stewart91}) and typically lasted 2--5~days (see Figure~\\ref{fig:asm} and discussion below). The ratio of count rates in different ASM energy channels showed dramatic spectral softening at the onset of the flaring state and gradual hardening during the remainder of each cycle (\\cite{shirey96}). Similar behavior was seen in \\Ginga\\ ASM observations folded at the 16.55 day period (\\cite{tsunemi89}). Near phase zero, some of the cycles observed with the \\RXTE\\ ASM also showed brief dips below the 1~Crab level. Observations of dips near phase zero in Cir~X-1 with \\ASCA\\ (\\cite{brandt96}) and the \\RXTE\\ PCA (\\cite{bradt98}) indicate the presence of both a strongly absorbed spectral component and an unabsorbed component.  Observations of type~1 X-ray bursts demonstrate that Cir~X-1 is a low magnetic field neutron star (Tennant, Fabian, \\& Shafer 1986\\nocite{tennant86b}). Additional type~1 bursts have not been observed from Cir~X-1 since the \\EXOSAT\\ discovery, possibly because the source intensity has been higher during subsequent observations.  The rapid X-ray variability of Cir~X-1 at times resembles that of both ``atoll'' and ``Z'' low-mass X-ray binaries (LMXBs) as well as black-hole candidates (\\cite{oosterbroek95}). Quasi-periodic oscillations (QPOs) were reported at 1.4~Hz, 5--20~Hz, and 100--200~Hz in \\EXOSAT\\ data (\\cite{tennant87}, 1988\\nocite{tennant88}). Based on these data, it has been suggested that Cir~X-1 is an atoll source that can uniquely reach the Eddington accretion rate and exhibit normal/flaring branch QPOs at 5--20~Hz (\\cite{oosterbroek95}; \\cite{klis94}). Observations made during non-flaring phases with the \\RXTE\\ Proportional Counter Array (PCA) showed a QPO peak that varied from 1.3 to 12~Hz, flat-topped low-frequency noise (LFN), and a broad peak that varied from 20--100~Hz (\\cite{shirey96}). The two QPO frequencies and the cut-off frequency of the flat-topped noise were highly correlated. These QPOs are likely to be essentially the same phenomenon as those previously seen in the \\EXOSAT\\ observations at 5--20~Hz and 100--200~Hz.  In this paper we present additional \\RXTE\\ ASM observations of Cir~X-1 which further illustrate how its intensity profile varies from one 16.55~d cycle to another. We also present the results of \\RXTE\\ PCA observations made over the course of one cycle in which the intensity declined unusually gradually from the flaring state to the quiescent level. This slow transition allows us to demonstrate how the time-variability properties of the source are related to its spectral properties. ",
        "conclusions": "The combined temporal and spectral-branch properties of the observations presented here suggest Z-like behavior. We identify the 6.8--32~Hz QPOs as horizontal-branch oscillations (HBOs), the 4~Hz QPO as normal-branch oscillations (NBOs), and the strong VLFN as flaring-branch behavior (see discussion below).  These identifications of characteristic time-variability patterns then help to identify the tracks in the HID as horizontal, normal, and flaring branches (HB/NB/FB), where each 6 ks observation of Cir~X-1 appears to have captured a snapshot of portions of one or two of the branches.  The spectral branches appear to shift around as the flaring gradually subsides, rather than forming a stable Z pattern. It is likely that the shapes of the spectral branches become distorted somewhat during these large shifts.  We now describe the inferred properties of each of the spectral branches in more detail.  \\subsection{Horizontal Branch} HID regions VIII, VII, and VI-1 show a narrow QPO peak or knee at 6.8--7.6~Hz, 11.3--13.1~Hz, and 32~Hz respectively. This frequency range overlaps the 13--60~Hz range of typical horizontal branch QPOs (\\cite{klis95}). The associated low-frequency noise and harmonic peak are also typical of horizontal branch power spectra. The broad high frequency peak in Cir~X-1 may be related to the high frequency noise component often observed on the horizontal branch.  The HID track for observation VI shows the narrow QPO at 32~Hz on a roughly horizontal segment (region VI-1) and a knee at 37~Hz on the right end of this segment (region VI-2). The apex of region VI-2 brings a transition to the 4~Hz QPO, which is dominant on the downward branch of this track (region VI-3). This is very similar to the HB/NB transition in Z sources.  When Cir~X-1 is in ``quiescence'' in observations VII and VIII, the ``horizontal branch'' turns upward and becomes vertical in the HID. For comparison, \\RXTE\\ PCA observations of Cir~X-1 from 1996 March 10--19 which show a narrow QPO peak at 1.3--12~Hz (\\cite{shirey96}) are almost entirely confined to the 12.3--14.7 kcts/s (2--21 keV) intensity range. The HID tracks for those observations lie along a nearly vertical line, and probably represent sections of the ``horizontal'' branch.  Observation II may also be on part of the HB, since a weak narrow QPO appears to evolve into a knee and increase in frequency from 22 to 30~Hz as the intensity increases. However, the broad QPO is also weakly visible in PDSs for this observation. The fact that observations II, VI, VII, and VIII all show little variation of the hard color used in Figure~\\ref{fig:cchid_all}a suggests that observation~II may be associated with the other HB observations.  \\subsection{Normal Branch}  The 4~Hz QPO is observed when the source intensity rises above the ``quiescent'' 1-Crab level ($\\sim$13~kcts/s). It is roughly stationary in frequency (3.3--4.3~Hz when clearly peaked) and broader than the HBO. The feature is easily seen in observations III--VI; at these times the location in the HID moves along diagonal tracks. The $\\sim$4~Hz frequency and motion along diagonal tracks in the HID is consistent with the 4--7~Hz NBOs observed at nearly constant frequency on the NB of typical Z sources (\\cite{hk89}). We therefore identify the broad 4~Hz QPO as a normal branch oscillation, and the diagonal tracks for observations III--VI as shifted normal branches.  The broad QPO component may be also present in the highest intensity observations, as a weak feature in observation II and in the form of a break near the 4~Hz QPO frequency in observation~I.  We also note that at the top of the normal branch (regions I-1, III-1, IV-1, VI-2) a knee above 30~Hz is present in addition to the NBO component.  A similar broad 4~Hz QPO is present in observations from 1996 March 5--6 made immediately before phase zero of the cycle showing the 1.3--12~Hz narrow QPO.  \\subsection{Flaring Branch} Beyond the left apex of the normal branch a short upturned branch is observed in HID region~V-3 and possibly III-3. The PDS for these regions are dominated by very low frequency noise, which is typical for flaring branches, and no QPO peaks are obviously apparent. We note that in the well-established Z sources neither NBOs nor HBOs are present on the flaring branch, except for Sco~X-1 and GX~17+2, in which the NBO evolves into a 6--20~Hz QPO (\\cite{klis95} and references therein).  The left end of the spectral track for observation~V bends upward in the HID shown in Figure~\\ref{fig:cchid_all}b, but bends downward in the CD in Figure~\\ref{fig:cchid_all}a. This behavior is demonstrated more clearly in Figure~\\ref{fig:cchid_56medhard}, which shows CDs and HIDs for observations V and VI. When a broad color (\\broadcolor) is used as the ordinate of the diagrams (Figure~\\ref{fig:cchid_56medhard}a,b), the track for observation~V turns upward on the left end. When a harder color (\\hardcolor) is used as the ordinate (Figure~\\ref{fig:cchid_56medhard}c,d), this branch turns downward. The CD and particularly the HID version based on the harder color show the most clear similarity to canonical Z diagrams, with the temporal behavior of observations V and VI being generally consistent with horizontal, normal, and flaring branches. The broad-color HID (Figure~\\ref{fig:cchid_56medhard}b) shows evidence for a shift of the normal branch that does not show up in the other three diagrams of that figure.  \\subsection{Relation to Other Sources} Our observations reveal spectral branches which shift in the CD and HID as Cir~X-1 evolves from a soft, high-intensity state to a hard, lower-intensity state. The ASM light curves and hardness ratios (Figure~\\ref{fig:asm}) show that this evolution occurs periodically with the 16.55~day cycle, thus suggesting that the CD/HID shifts may also be periodic. Shifts of the ``Z'' pattern in CDs and HIDs have been observed in the so-called Cyg-like Z sources: Cyg~X-2 (Kuulkers, van~der~Klis, \\& Vaughan 1996\\nocite{kuulkers96:cygx-2}), GX~5-1 (\\cite{kuulkers94:gx5-1}), and GX~340+0 (\\cite{kuulkers96:gx340+0}). However, the shifts do not occur periodically in those sources, nor do they have the magnitude of the shifts observed in Cir~X-1.  The flaring branch of Cir X-1 turns upward when a soft or broad color is used on the vertical axis. When a harder color is used, this branch turns downward but then bends to the left. In the Cyg-like Z sources, the flaring branch sometimes turns upward or starts toward higher intensity and then loops back to lower intensity (\\cite{kuulkers96:cygx-2}; \\cite{kuulkers94:gx5-1}; \\cite{kuulkers96:gx340+0}; \\cite{penninx91:gx340+0}). In some cases, these sources are observed to ``dip'' while on the flaring branch (\\cite{kuulkers94:gx5-1}; \\cite{penninx91:gx340+0}; \\cite{wijnands97}), with tracks which turn down and then to the left, similar to that of Cir~X-1 in Figure~\\ref{fig:cchid_56medhard}c.  The left end of the horizontal branch in Cir~X-1 turns upward and becomes vertical at low intensity (Figure~\\ref{fig:hidqpo}). On this section of the branch, HBO frequencies are low: 6.8--13~Hz in observations VII and VIII and 1.3--12~Hz in the earlier 1996 March observations. A similar effect was reported in GX 5-1 (\\cite{lewin92:gx5-1}; \\cite{kuulkers94:gx5-1}), in which the HB turns upward at the low-intensity end while HBOs are observed at relatively low frequency (13--17~Hz). In fact, Lewin et~al.\\ (1992\\nocite{lewin92:gx5-1}) suggested that other Z sources might show such an upward turn of the HB if their intensities and QPO frequencies became sufficiently low.  The 5--20~Hz narrow QPO was detected with \\EXOSAT\\ at an intensity similar to the quiescent level observed by \\RXTE. We note that absorption dips are responsible for much of the structure seen in the CD shown for that observation; however, the HIDs show that the narrow QPO occurred on an upturned left end of a horizontally oriented track as in our data (see Figures~2--4, 8, \\& 10 in \\cite{oosterbroek95}). At higher intensity during the same observation, the narrow QPO was not present, and we note that some of the high-intensity PDSs show hints of a broad peak near 4~Hz. We thus conclude that the behavior observed by \\EXOSAT\\ during that observation is related to the Z-like behavior we observe with \\RXTE.  Most of the other \\EXOSAT\\ observations took place when Cir~X-1 was significantly lower in intensity than the ``quiescent'' level of the current observations. The CDs and HIDs for these \\EXOSAT\\ observations did not show tracks which could clearly be identified as Z or atoll. Their power spectra were generally dominated by VLFN, typical of atoll sources in the banana state, and sometimes also showed a broad red noise component resembling atoll high-frequency noise (\\cite{oosterbroek95}). However, these power-spectral shapes are not unique to atoll sources: power spectra for black hole candidates in the high state are dominated by VLFN, as are those of the current observations on the low-intensity end of the normal branch and on the flaring branch (i.e., regions III-3 and V-3).  Cir~X-1 was expected to never show HBOs since atoll-like behavior was taken as evidence that the magnetic field is not strong enough to allow the magnetospheric beat frequency mechanism (MBFM) to operate (\\cite{klis94}; \\cite{oosterbroek95}). (However, it is also possible that the HBOs are not produced by the MBFM\\@.) The results presented here demonstrate both HBOs and NBOs in Cir~X-1 and show no evidence for atoll behavior. Since the atoll-like behavior observed with \\EXOSAT\\ occurred at lower intensity than in the present observations, it is possible that they do represent a different state of the source. If Cir~X-1 actually can show atoll behavior as well as the Z-like behavior shown here, then we would have new clues to the differences between the two types of sources. Such observations would challenge the hypothesis that differences in both $\\dot{M}$ and magnetic field distinguish these two classes."
    },
    "9803/astro-ph9803108_arXiv.txt": {
        "abstract": "We report on mid-infrared (MIR) continuum and line emission mapping of the nucleus of NGC~253. The data, with a resolution of 1\\as.4, reveal a double-peaked arc-like [NeII] emission region. Comparison with published data shows that the [NeII] arc is centered on the nucleus of the galaxy. The brightest [NeII] source coincides with the infrared continuum peak. The interpretation of these results is complicated by the edge-on orientation of NGC~253, but a self-consistent explanation is starformation triggered by dynamical resonances in a barred potential. ",
        "introduction": "The prototypical starburst galaxy NGC~253 is a highly inclined (i=78\\deg , \\cite{pen81}) disk galaxy at a distance of about 3~Mpc (\\cite{tul88}). The view towards its inner region is heavily obscured, and the exact location of its nucleus has been the matter of some debate. The astrometry of the various infrared sources in the central 10\\as\\ has been discussed in detail by Kalas \\& Wynn-Williams (1994, hereafter KW94) in the near-infrared and \\cite{ket93} in the mid-infrared. From their $\\approx$1\\as\\ resolution H, K, L, and M observations, KW94 conclude that the brightest source in all the maps (their peak~1) lies at R.A.(1950)=$00^h45^m5^s.62$, Decl.(1950)=$-25^{\\circ}33^{\\prime}40^{\\prime\\prime}.2$. This position is in very good agreement with that of the emission peak at MIR wavelengths as determined by Keto \\ea\\ (1993) at 10~\\mm\\ from astrometric methods and Pi\\~na \\ea\\ (1992) from low-level contour fitting to published maps with lower resolution. It is safe to assume that the location of peak~1 is independent of wavelength up to 20~\\mm , enabling us to adopt the above coordinates for the continuum peak in our maps.  Peak~1 is not connected to any of the strong radio point sources found at 2~cm by Turner \\& Ho (1985, hereafter TH85) and 6~cm by \\cite{ulv91} and \\cite{ant88}. The radio sources are highly aligned along the major axis of NGC~253 at P.A. 51\\deg . All of them are most likely either radio supernovae or HII regions, based on their spectral index. The only exception is the source TH~2 in TH85, a powerful flat spectrum radio source that they proposed to be the nucleus of NGC~253. TH85 argue that because of its high radio luminosity ($10^4$~\\lsol ) and its unique location at the center of a synchroton disk this source probably is a compact synchroton source, similar to those observed in active galactic nuclei (AGNs). It lies close to a secondary MIR peak (peak~2) $\\approx$ 2\\as.2 northeast of peak~1 (\\cite{pin92}, \\cite{ket93}). However, as we will show, the radio nucleus is most likely not identical with any characteristic infrared source. This view was strongly supported from 0.5\\as\\ resolution NIR color maps by \\cite{sam94}. Their maps were referenced to other observations by cross-correlation of low level contours and show that the radio nucleus TH~2 coincides with an extinction maximum of $\\rm A_V\\geq 24$. In addition, KW94 show that the exact position of peak~2 is wavelength dependent in the sense that for J, H, K observations it lies $\\approx$ 1\\as\\ further east than in M-band.  The physical nature behind the various infrared sources is uncertain. For example, according to \\cite{for93}, the emission of the \\brg -line at 2.17~\\mm\\ is strongest at peak~1. On the other hand, KW94 show that peak~1 has a low PAH-feature/continuum emission ratio. Since PAH emission is usually indicative of ongoing star formation, KW94 argue against peak~1 being an intense starburst which would be the natural explanation for the \\brg\\ emission. To investigate this matter further, we have observed NGC~253 at MIR wavelengths, both in continuum and line emission. The forbidden $\\rm ^2P_{3/2}-^2P_{1/2}$ transition of singly ionized Neon at 12.81~\\mm\\ traces the photoionization regions of hot stars and thus carries similar information than e.g. \\brg , but at a much lower extinction: $\\rm A_{2.17\\mu m} \\approx 2.5\\cdot A_{12.8\\mu m}$ (\\cite{lut96,gen98}).  In the next section we will describe how the data were obtained. We also discuss the special observing techniques in the MIR and their implications on data reduction. We present the results of our observations in Sec. \\ref{results} and interpret them in Sec. \\ref{discussion}. ",
        "conclusions": "\\label{discussion} \\subsection{The bar in NGC~253} From 10\\as\\ resolution K-band imaging, \\cite{sco85} find a prominent NIR bar at P.A. 68\\deg . However, the NIR continuum of the central $\\approx$ 10\\as\\ in all high resolution NIR maps (e.g. Forbes \\ea\\ 1993, KW94, Sams \\ea\\ 1994) is oriented at P.A. 51\\deg , roughly parallel to the major axis of the optical disk. This change in orientation is also visible as an isophote twist in \\cite{sco85} and all MIR maps (\\cite{pin92,ket93}, Fig. \\ref{fig2}). What is the reason for this behavior? From kinematical studies of the CS(2-1) line at 98~GHz, \\cite{pen96} have tried to answer this question. They derived a model for the dynamical processes in the center of NGC~253. In short, they find five prominent CS emission spots. The two innermost spots are aligned along the axis of the radio knots, i.e. the major axis of NGC~253. They are separated by $\\approx$ 3\\as\\ and positioned symmetrically on either side of the radio nucleus. Two other knots lie again symmetrically to the nucleus, but are aligned with the large scale NIR bar. Their observations are well explained by a rotating gas torus around the dynamical center of NGC~253. The torus is the response of the viscous molecular gas to the potential of the NIR stellar bar. Gas clouds in the inner $\\approx$ 10\\as\\ move along elliptical $x_2$ orbits whose major axis is aligned perpendicular to the bar (Athanassoula 1992a,b). At the inclination of NGC~253, these orbits appear to be oriented parallel to the major axis of the optical disk as demonstrated in Fig.~5 of \\cite{pen96}. The gas clouds outside of $\\approx$ 10\\as\\ move on $x_1$ orbits, their major axis being aligned with the bar. The existence of $x_2$ orbits depends on the presence of at least one Inner Lindblad Resonance (ILR). For the case of NGC~253, \\cite{arn95} have shown that there are in fact two ILRs, one at a radius of 25\\as, the other close to the nucleus. It is at the inner ILR, where the orientation of the dominant orbits changes from $x_2$ to $x_1$ (e.g. \\cite{tel88}). Because of orbit crowding, massive star formation occurs predominantly at the apocenters of the $x_2$ orbits (\\cite{ath92b}), oriented symmetrically on either side of the nucleus. Therefore, \\cite{pen96} interpret their CS results as evidence for a rotating torus of dense gas around the radio nucleus TH~2. These results are in agreement with observations of other tracers of high density gas like HCN (\\cite{pag95}). \\subsection{The [NeII] arc - a starformation ring} Based on their model, \\cite{pen96} predict that higher resolution line mapping should prove that gas emission from the central 10\\as\\ (145~pc) of NGC~253 is aligned with the radio knots rather than with the optical disk. In that context, our results provide an independent confirmation of this scenario. Firstly, the [NeII] emission is well aligned with the radio knots (Figs. \\ref{fig1} (Plate 1) and \\ref{fig2}). Secondly, based on the structure of the [NeII] emission, we also confirm the distribution of the star forming material in a ring around the radio nucleus. The ring has a rotation velocity of $\\approx$ 60~km/s as seen in the CS data by \\cite{pen96}, the western side moving away from the observer. Unfortunately, our observations are unable to resolve the velocity structure of the [NeII], this has to await higher resolution line mapping, e.g. with a cryogenic Fabry-Perot interferometer.  Our data only show weak emission from the southeastern half of the ring. On the other hand, the [FeII] map of \\cite{for93} does show emission from this region (their source B). Thus, we favor the interpretation of the double-peaked morphology as a starburst ring with a diameter of $\\approx$ 4\\as\\ (60~pc). The total observed flux from a circular aperture with 1\\as.4 (20~pc) diameter centered on peak~1 is $1.7\\cdot 10^{-15}$~W/m$^2$. After dereddening with $\\rm A_V=24$ (Sams \\ea\\ 1994) in a mixed case scenario, this corresponds to $7.7\\cdot 10^5$\\lsol\\footnote{We have adopted $\\rm A_{12.8\\mu m}=0.04\\cdot A_V$ (\\cite{gen98})}.  With the conversion factor $\\rm \\frac{L_{Lyc}}{L_{[NeII]}} = 64$ (\\cite{gen98}), the intrinsic Lyman-continuum luminosity is $4.9\\cdot 10^7$\\lsol. This value should be used with caution given its uncertainties, but it compares well to other bright star forming regions like the core of 30~Doradus in the LMC (\\cite{bra96}) or the brightest SFR in the gaseous ring of IC~342 (\\cite{boe97b}). This supports the view that the [NeII] emission probably stems from individual giant molecular clouds (GMCs) that actively form stars and are located in a ring at the ILR. We will extend the comparison to a similar structure in IC~342 in the next section. \\subsection{NGC~253 and IC~342 - two of the same kind?} It is interesting to note the close similarities between NGC~253 and IC~342, another nearby late type spiral. B\\\"oker \\ea\\ (1997b) describe the dynamical processes in the central 100~pc of IC~342 as derived from NIR integral field spectroscopy. Table \\ref{tab1} compares the two galaxies. \\begin{table} \\dummytable \\label{tab1} \\end{table} The dynamical processes and the morphology of the molecular gas are almost identical. This was also noted by Paglione \\ea\\ (1995). One important difference, however, is that IC~342 does not house a central synchroton source. In fact, in the case of IC~342, there is no radio emission from the dynamical center (\\cite{tur83}). Rather, it is occupied by a cluster of young red supergiants that dominate the NIR continuum. For NGC~253, there is no obvious NIR component at the location of the radio nucleus, but the edge-on orientation and patchy extinction might complicate its detection. Both the diameter and rotation velocity of the molecular rings are very similar. This adds to the increasing evidence that stellar bars and small star formation rings on scales of 50-100~pc are a common feature in late type spirals. This points to an evolutionary connection between spirals of different Hubble type, as suggested by \\cite{pfe94}."
    }
}