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9803 | astro-ph9803011_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803011_arXiv.txt |
9803 | astro-ph9803227_arXiv.txt | 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. | 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. | 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} | 98 | 3 | astro-ph9803227_arXiv.txt |
9803 | astro-ph9803157_arXiv.txt | 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. | \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. | \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. | 98 | 3 | astro-ph9803157_arXiv.txt |
9803 | astro-ph9803196.txt | } \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. | %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% 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. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % | %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% 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. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 98 | 3 | 9803.196 |
9803 | astro-ph9803143_arXiv.txt | 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 | 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. | 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. | 98 | 3 | astro-ph9803143_arXiv.txt |
9803 | astro-ph9803233_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803233_arXiv.txt |
9803 | astro-ph9803005_arXiv.txt | 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. | \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. | \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}). | 98 | 3 | astro-ph9803005_arXiv.txt |
9803 | astro-ph9803280_arXiv.txt | 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. | \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. | 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. | 98 | 3 | astro-ph9803280_arXiv.txt |
9803 | astro-ph9803249_arXiv.txt | 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. | 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$^{-1}$. | 98 | 3 | astro-ph9803249_arXiv.txt |
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9803 | astro-ph9803139_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803139_arXiv.txt |
9803 | astro-ph9803162_arXiv.txt | 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. | \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}. | \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). | 98 | 3 | astro-ph9803162_arXiv.txt |
9803 | astro-ph9803024_arXiv.txt | \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. | } 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. | 98 | 3 | astro-ph9803024_arXiv.txt |
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9803 | astro-ph9803268_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803268_arXiv.txt |
9803 | astro-ph9803118_arXiv.txt | 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. | \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. | 98 | 3 | astro-ph9803118_arXiv.txt |
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9803 | astro-ph9803083_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803083_arXiv.txt |
9803 | astro-ph9803340_arXiv.txt | 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. | \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}. | \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. | 98 | 3 | astro-ph9803340_arXiv.txt |
9803 | astro-ph9803263_arXiv.txt | 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). | 98 | 3 | astro-ph9803263_arXiv.txt |
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9803 | astro-ph9803055_arXiv.txt | 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). | 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. | 98 | 3 | astro-ph9803055_arXiv.txt |
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9803 | astro-ph9803325_arXiv.txt | 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. | 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. | 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). | 98 | 3 | astro-ph9803325_arXiv.txt |
9803 | astro-ph9803113_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803113_arXiv.txt |
9803 | astro-ph9803169_arXiv.txt | 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. | 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).}. | 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 | 98 | 3 | astro-ph9803169_arXiv.txt |
9803 | gr-qc9803068_arXiv.txt | 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. | \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. | 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* | 98 | 3 | gr-qc9803068_arXiv.txt |
9803 | astro-ph9803107_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803107_arXiv.txt |
9803 | astro-ph9803277_arXiv.txt | 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. | 98 | 3 | astro-ph9803277_arXiv.txt |
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9803 | astro-ph9803195_arXiv.txt | 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. | \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}. | 98 | 3 | astro-ph9803195_arXiv.txt |
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9803 | astro-ph9803126_arXiv.txt | 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. | 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 | 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). | 98 | 3 | astro-ph9803126_arXiv.txt |
9803 | astro-ph9803256_arXiv.txt | 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. | 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}). | 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. | 98 | 3 | astro-ph9803256_arXiv.txt |
9803 | astro-ph9803074_arXiv.txt | \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. | 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. | 98 | 3 | astro-ph9803074_arXiv.txt |
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9803 | astro-ph9803242_arXiv.txt | 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. | 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. | 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* | 98 | 3 | astro-ph9803242_arXiv.txt |
9803 | astro-ph9803132_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803132_arXiv.txt |
9803 | astro-ph9803304_arXiv.txt | \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. | 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. | 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 | 98 | 3 | astro-ph9803304_arXiv.txt |
9803 | astro-ph9803148_arXiv.txt | 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$0.7. These data support a model in which most production of B and Be comes from the spallation of energetic C and O nuclei onto protons and He nuclei, probably in the vicinity of massive supernovae in star-forming regions, rather than the spallation of cosmic ray protons and alpha particles onto CNO nuclei in the general interstellar medium. | The light elements lithium, beryllium, and boron are of great interest out of proportion to their very low abundances, having implications in Big Bang Nucleosynthesis and stellar structure, as well as in constraints on models of galactic chemical evolution. The ``canonical'' theory of the origin of the elements Li, Be, and B was first presented by Reeves, Fowler, \& Hoyle (1970) and further developed by Meneguzzi, Audouze, \& Reeves (1971), and then Reeves, Audouze, Fowler, \& Schramm (1973). In this model, most light element formation can be accounted for by galactic cosmic rays (GCR) impinging on the interstellar medium (ISM), assuming a constant flux of GCRs through the life of the Galaxy and making reasonable assumptions about CR confinement by the Galactic magnetic field. Meneguzzi \etal\ (1971) also introduced the idea of a large (up to three orders of magnitude) increase in the low energy (5-40 MeV nucleon$^{-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. | 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. | 98 | 3 | astro-ph9803148_arXiv.txt |
9803 | astro-ph9803181_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803181_arXiv.txt |
9803 | astro-ph9803238_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803238_arXiv.txt |
9803 | astro-ph9803328_arXiv.txt | 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. | 98 | 3 | astro-ph9803328_arXiv.txt |
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9803 | astro-ph9803058_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803058_arXiv.txt |
9803 | astro-ph9803022_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803022_arXiv.txt |
9803 | astro-ph9803214_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803214_arXiv.txt |
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\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{${\vcenter{\vbox{ % \hrule height 0.4pt\hbox{\vrule width 0.4pt height 6pt \kern5pt\vrule width 0.4pt}\hrule height 0.4pt}}}$}} \renewcommand{\thefootnote}{\fnsymbol{footnote}} % \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{} % \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}} | \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. | \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 | 98 | 3 | astro-ph9803164_arXiv.txt |
9803 | astro-ph9803346_arXiv.txt | 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) | 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. | 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. | 98 | 3 | astro-ph9803346_arXiv.txt |
9803 | astro-ph9803170_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803170_arXiv.txt |
9803 | astro-ph9803200_arXiv.txt | 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 | 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. | 98 | 3 | astro-ph9803200_arXiv.txt |
|
9803 | astro-ph9803085_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803085_arXiv.txt |
9803 | astro-ph9803198_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803198_arXiv.txt |
9803 | astro-ph9803017_arXiv.txt | 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. | 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$^{-1}$. | \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. | 98 | 3 | astro-ph9803017_arXiv.txt |
9803 | astro-ph9803221_arXiv.txt | 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. | 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. | 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}). | 98 | 3 | astro-ph9803221_arXiv.txt |
9803 | hep-ph9803378_arXiv.txt | 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. | 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. | 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. | 98 | 3 | hep-ph9803378_arXiv.txt |
9803 | hep-ph9803244_arXiv.txt | { \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 | \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} | \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. | 98 | 3 | hep-ph9803244_arXiv.txt |
9803 | astro-ph9803309_arXiv.txt | 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. | 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. | 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 | 98 | 3 | astro-ph9803309_arXiv.txt |
9803 | astro-ph9803286_arXiv.txt | \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. | 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$4.0, we discuss these results and their implications for the formation and evolution of dwarf galaxies, in particular with respect to various interaction scenarios. The 21 cm and optical observations together provide a unique source of information on the nature of BCDs and their companions in clusters. | \subsection{The \ion{H}{1} Cloud} A position-velocity cut through the center of both FCC 35 and the \ion{H}{1} cloud (Fig.~\ref{fig:f13}) shows that despite the spatial overlap , the two components are not connected in velocity space. The cloud may either be a unique source within the Fornax cluster or a foreground or background \ion{H}{1} object. The former seems to be the most likely considering the velocity of the cloud (V$_{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. | 98 | 3 | astro-ph9803286_arXiv.txt |
9803 | astro-ph9803235_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803235_arXiv.txt |
9803 | astro-ph9803003_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803003_arXiv.txt |
9803 | astro-ph9803145_arXiv.txt | 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. | 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} | 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} | 98 | 3 | astro-ph9803145_arXiv.txt |
9803 | astro-ph9803102_arXiv.txt | 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$\beta$). | 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. | \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. | 98 | 3 | astro-ph9803102_arXiv.txt |
9803 | astro-ph9803334_arXiv.txt | 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. | 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. | 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 | 98 | 3 | astro-ph9803334_arXiv.txt |
9803 | astro-ph9803044_arXiv.txt | 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. | 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. | 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 | 98 | 3 | astro-ph9803044_arXiv.txt |
9803 | astro-ph9803272_arXiv.txt | 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): $1.06^{+0.22}_{-0.26}$. | 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. | 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. | 98 | 3 | astro-ph9803272_arXiv.txt |
9803 | astro-ph9803208_arXiv.txt | 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}. | \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. | \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. | 98 | 3 | astro-ph9803208_arXiv.txt |
9803 | astro-ph9803178_arXiv.txt | 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. | 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} | 98 | 3 | astro-ph9803178_arXiv.txt |
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9803 | astro-ph9803099_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803099_arXiv.txt |
9803 | astro-ph9803050_arXiv.txt | 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}. | 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. | 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). | 98 | 3 | astro-ph9803050_arXiv.txt |
9803 | astro-ph9803320_arXiv.txt | 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. | \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}. | \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. | 98 | 3 | astro-ph9803320_arXiv.txt |
9803 | astro-ph9803116_arXiv.txt | 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. | 98 | 3 | astro-ph9803116_arXiv.txt |
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9803 | astro-ph9803184_arXiv.txt | 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. | \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. | \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. | 98 | 3 | astro-ph9803184_arXiv.txt |
9803 | astro-ph9803071_arXiv.txt | 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}. | 98 | 3 | astro-ph9803071_arXiv.txt |
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9803 | astro-ph9803247_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803247_arXiv.txt |
9803 | astro-ph9803137_arXiv.txt | 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 $z>5$. | 98 | 3 | astro-ph9803137_arXiv.txt |
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9803 | astro-ph9803301_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803301_arXiv.txt |
9803 | astro-ph9803315_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803315_arXiv.txt |
9803 | astro-ph9803123_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803123_arXiv.txt |
9803 | astro-ph9803065_arXiv.txt | 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. | 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*} | \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. | 98 | 3 | astro-ph9803065_arXiv.txt |
9803 | astro-ph9803229_arXiv.txt | 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. | 98 | 3 | astro-ph9803229_arXiv.txt |
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9803 | physics9803022_arXiv.txt | 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. | 98 | 3 | physics9803022_arXiv.txt |
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9803 | astro-ph9803159_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803159_arXiv.txt |
9803 | astro-ph9803190_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803190_arXiv.txt |
9803 | astro-ph9803049_arXiv.txt | 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.} | 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}? | 98 | 3 | astro-ph9803049_arXiv.txt |
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9803 | astro-ph9803080_arXiv.txt | \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. | \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. | \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. | 98 | 3 | astro-ph9803080_arXiv.txt |
9803 | astro-ph9803339_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803339_arXiv.txt |
9803 | astro-ph9803175_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803175_arXiv.txt |
9803 | astro-ph9803205_arXiv.txt | 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 | 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. | 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 The 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 The 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 The 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 The 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 | 98 | 3 | astro-ph9803205_arXiv.txt |
9803 | astro-ph9803027_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803027_arXiv.txt |
9803 | astro-ph9803094_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803094_arXiv.txt |
9803 | astro-ph9803230_arXiv.txt | 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 ($>$1) likely indicate optically thin gas created by the high temperatures caused by star forming regions in the nucleus of this starburst galaxy. | The circumnuclear regions of galaxies are very often the setting for starbursts and other extraordinary events. Observations of the gas kinematics and distribution indicate that bars, resonances, gas inflow, and tidal shear play important roles in the formation and evolution of nuclear starbursts (\eg\ Handa \etal\ 1990; Kenney \etal\ 1992). Previous observations show that the molecular gas in the central regions of barred spiral galaxies often does not extend all the way into the centre of the nucleus. It accumulates some distance away from the centre, giving the emission a double peaked appearance with each peak occurring where the bar meets the nucleus (Kenney \etal\ 1992; Ishizuki \etal\ 1990). Dynamical models indicate that in the presence of a barred potential, gas will flow inward along the bar and slow its descent temporarily near inner Lindblad resonances (ILR, \eg\ Combes 1988; Shlosman, Frank \& Begelman 1989). At these locations the gas may accumulate into larger complexes of molecular clouds. In addition to complex dynamics, there is most likely a profusion of complicated photo-chemistry occurring within the nuclei of starburst galaxies. Interstellar clouds are believed to consist of smaller high density dark cores interspersed throughout a larger region of lower density. These high density regions are self-shielded from ultraviolet (UV) radiation that tends to dissociate molecular gas. The result is the population of the diffuse region by hydrogen atoms (\ion{H}{1}), atomic and ionized carbon (C, C$^+$), and many other atomic and ionized species (\eg\ Morton \etal\ 1973), while the dark cores can contain molecular species such as H$_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. | } 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} | 98 | 3 | astro-ph9803230_arXiv.txt |
9803 | astro-ph9803006_arXiv.txt | 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. | 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$1.2\arcsec) and sensitivity requirements. Contrary to the synthesis aperture maps so far published, we recover entirely the zero-spacing flux of the {\bf ND}. The comparison between the different gaseous and stellar tracers of the {\bf ND} is reexamined in this work. Particular attention is paid to the bias introduced by extinction in optical and even near-infrared images of the nucleus, and what might be the implications on the interpretation of the data. Observations are confronted to the result of new numerical simulations of the clouds hydrodynamics, based on a mass distribution directly derived from the infrared luminosity image of M100. We will focus on the feasibility of two independent patterns in the disk and how this scenario accounts better for the observations. | Simulations of the H$_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. | 98 | 3 | astro-ph9803006_arXiv.txt |
9803 | astro-ph9803140_arXiv.txt | 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. | 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*} | 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. | 98 | 3 | astro-ph9803140_arXiv.txt |
9803 | astro-ph9803283_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803283_arXiv.txt |
9803 | astro-ph9803297_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803297_arXiv.txt |
9803 | astro-ph9803318_arXiv.txt | 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. | 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. | \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} | 98 | 3 | astro-ph9803318_arXiv.txt |
9803 | astro-ph9803154_arXiv.txt | 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. | 98 | 3 | astro-ph9803154_arXiv.txt |
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9803 | astro-ph9803224_arXiv.txt | 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. | 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. | 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 | 98 | 3 | astro-ph9803224_arXiv.txt |
9803 | astro-ph9803131_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803131_arXiv.txt |
9803 | astro-ph9803307_arXiv.txt | 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. | 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). | 98 | 3 | astro-ph9803307_arXiv.txt |
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9803 | astro-ph9803288_arXiv.txt | 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. | 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. | 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. | 98 | 3 | astro-ph9803288_arXiv.txt |
9803 | astro-ph9803182_arXiv.txt | 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$-$63 on a measurable level. | The pulsar PSR B1259$-$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. | 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. | 98 | 3 | astro-ph9803182_arXiv.txt |
9803 | astro-ph9803019_arXiv.txt | 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. | 98 | 3 | astro-ph9803019_arXiv.txt |
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9803 | astro-ph9803255_arXiv.txt | 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. | \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. | \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.} | 98 | 3 | astro-ph9803255_arXiv.txt |
9803 | astro-ph9803313_arXiv.txt | 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. | 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. | 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$), and are separated by a large luminosity gap (0.5~dex) from stars on the populous EHB. As discussed by Landsman et al.\ (\cite{land96}), these two characteristics are consistent with the non-canonical HB models of Sweigart (\cite{sweig97}), which include helium mixing on the RGB. However, a more definitive test of EHB evolutionary tracks will require a larger sample of post-EHB stars. | 98 | 3 | astro-ph9803313_arXiv.txt |
9803 | astro-ph9803125_arXiv.txt | We describe an automated method for detecting clusters of galaxies in imaging and redshift galaxy surveys. The Adaptive Matched Filter (AMF) method utilizes galaxy positions, magnitudes, and---when available---photometric or spectroscopic redshifts to find clusters and determine their redshift and richness. The AMF can be applied to most types of galaxy surveys: from two-dimensional (2D) imaging surveys, to multi-band imaging surveys with photometric redshifts of any accuracy (2$\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. | 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. | 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. | 98 | 3 | astro-ph9803125_arXiv.txt |
9803 | astro-ph9803038_arXiv.txt | 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. | 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. | 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). | 98 | 3 | astro-ph9803038_arXiv.txt |
9803 | astro-ph9803206.txt | 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. | 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. | 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. | 98 | 3 | 9803.206 |
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