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9803 | nucl-th9803026_arXiv.txt | The properties of hot matter are studied in the frame of the relativistic Brueckner-Hartree-Fock theory. The equations are solved self-consistently in the full Dirac space. For the interaction we used the potentials given by Brockmann and Machleidt. The obtained critical temperatures are smaller than in most of the nonrelativistic investigations. We also calculated the thermodynamic properties of hot matter in the relativistic Hartree--Fock approximation, where the force parameters were adjusted to the outcome of the relativistic Brueckner--Hartree--Fock calculations at zero temperature. Here, one obtains higher critical temperatures, which are comparable with other relativistic calculations in the Hartree scheme. | The properties of hot and dense nuclear matter play an essential role in the understanding of high-energy heavy-ion collisions, supernova explosions and proto-neutron stars. For that reason the problem of hot nuclear matter has been studied over the last decades in several investigations, which however were predominantly performed within the nonrelativistic scheme \cite{1a}, using either effective density dependent interactions \cite{1a,1,2,23} or the Brueckner approach \cite{1a,3,4,5,5a}. In the relativistic approach investigations of the equations of state for $T\neq 0$ are relatively scarce. The majority of such calculations were performed in the relativistic Hartree approximation (RH), where the extension to finite temperatures is straightforward. Details of this scheme are given, for instance, in Refs.\,\cite{23}, \cite{6}-\cite{11}. More complicated are the relativistic Hartree--Fock-- \cite{13,11} and the Brueckner--Hartree--Fock approximation \cite{12}, and the application to finite nuclei \cite{14}. In this contribution we will concentrate on the relativistic Brueckner--Hartree--Fock treatment (RBHF) of symmetric and asymmetric nuclear matter generalizing the formalism as described in Refs.\,\cite{15,15a} to $T\neq 0$. The RBHF--approach seems to be of special interest, since it is known for $T = 0$ that the resulting EOSs are much stiffer than their nonrelativistic counterparts \cite{15b}. To our knowledge such an investigation has been only performed so far by the Groningen group for symmetric nuclear matter \cite{12}. As described in more details in Refs.\,\cite{12,15}, their method uses a nonunique ansatz for the T--matrix in terms of five Fermi invariants, which can lead to ambiguous results for the self-energies (see, e.g., Refs.\,\cite{17,18}). In order to avoid this problem we solve, according to the original scheme of the Brooklyn group \cite{19}, the RBHF--approximation in the full Dirac space, which is more tedious (for a more detailed discussion, see Ref.\,\cite{15}). Since the formal structure of the problem is the same as for $T=0$, where one has to solve three coupled equations, namely the Dyson equation for the one-body Green's function $G$, the (reduced) Bethe--Salpeter equation for the effective scattering matrix $T$ in matter and the equation for the self--energy $\Sigma$, we will not repeat here the equations. As in Refs.\,\cite{12,15,15a} we will restrict ourselves to the incorporation of intermediate positive--energy nucleon states only, where now the Fermi step functions are replaced by the Fermi distribution functions $n_{\vec p}\,(T)$ at finite temperature $T$ (for details, see Ref.\,\cite{21}). The Green's function obeys for $T\neq 0$ the spectral representation \cite{11,20} \begin{equation} \label{I.1} G(p) = \int d\omega\,A(\vec{p},\omega) \left\{\frac{f(\omega)}{p_\rho -\omega - i\eta} + \frac{f(-\omega)}{p_\rho -\omega+i\eta}\right\} ~, \end{equation} with \begin{equation} \label{I.2} f(\omega) = (e^{\beta\omega}+1)^{-1} ~ {\stackrel{T=0}{_{\mbox{$\longrightarrow$}}}} ~ \Theta(-\omega) ~. \end{equation} The formal structure of the spectral function $A(\vec{p},\omega)$ \cite{15} remains unaltered to the case for $T=0$. A further difference in comparison with Ref.\,\cite{12} is that we take the momentum dependence of the self--energies into account. Since the pole of the quasi--particle propagators occurs for $T\neq 0$ in the integration domain of the intermediate states, one obtains, in principle, complex effective scattering matrix elements and self--energies. It was checked in Ref.\,\cite{12} that the imaginary part of $\Sigma$ turned out to be small. Therefore we neglect also Im~$\Sigma$ in the calculations. For the one--boson--exchange interaction we used the modern potentials constructed by Brockmann and Machleidt \cite{22}. We select for the presentation the so-called potential $B$, which gives the best results for the EOS ($E/A = -15.73$~MeV; $\rho_0 = 0.172$~fm$^{-3}; K = 249$~MeV; $J = 32.8$~MeV) at zero temperature in RBHF--calculations (see Refs.\,\cite{15,15a}, for the potential $A$ the outcome is similar, see Ref.\,\cite{21}). For the sake of comparison we also treated the RHF--approximation, where we adjusted the force parameters to the outcome for the EOS for symmetric and asymmetric nuclear matter at $T=0$ \cite{15a,21}. The RHF--approximation has in comparison with the RH--approximation the advantage to resemble in its formal structure more to the RBHF--approximation with the benefit of a much simpler numerical treatment than in the RBHF case. For finite temperatures one needs for the determination of the pressure the free energy per baryon, defined as \begin{equation} \label{I.3} f = u - T s ~, \end{equation} with the entropy per baryon: \begin{equation} \label{I.4} s = - \frac{2}{\rho h^3} \sum_\tau \int d^3p\,[n(\vec{p})\; {\rm ln} \,n(\vec{p}) + \left(1 - n(\vec{p}\right)\; {\rm ln}\,\left(1 - n(\vec{p}\right)]~. \end{equation} The pressure is given by: \begin{equation} \label{I.5} P = \rho \sum_\tau \rho_\tau \left(\frac{\partial f} {\partial\rho_\tau}\right)_{T,\rho_{-\tau}} ~. \end{equation} | In conclusion, we have performed a calculation of hot symmetric and asymmetric nuclear matter within the relativistic Brueckner--Hartree--Fock scheme using modern OBE--interactions constructed by Brockmann and Machleidt. It turned out that the critical temperatures are smaller than it is the case for the majority of nonrelativistic treatments. We have additionally treated the relativistic Hartree--Fock approximation at $T\neq 0$, where the Lagrangian parameters were adjusted to the outcome of the RBHF--treatment for $T=0$. Here the critical temperature is in the range of other relativistic treatments performed in the Hartree scheme. | 98 | 3 | nucl-th9803026_arXiv.txt |
9803 | astro-ph9803042_arXiv.txt | We report on the discovery of a new kind of thermal pulse in intermediate mass AGB stars. Deep dredge-up during normal thermal pulses on the AGB leads to the formation of a long, unburnt tail to the helium profile. Eventually the tail ignites under partially degenerate conditions producing a strong shell flash with very deep subsequent dredge-up. The carbon content of the intershell convective region (X$_{C}$ $\sim$ 0.6) is substantially higher than in a normal thermal pulse (X$_{C}$ $\sim$ 0.25) and about 4 times more carbon is dredged-up than in a normal pulse. | It is now just over 30 years since thermal pulses were discovered in AGB stars by \cite{Sch65} and over 20 years since third dredge-up was discovered by \cite{Ibe75}, and we are still learning about the consequences of these events for stellar evolution and nucleosynthesis. Although much is known there are still many uncertainties, especially concerning third dredge-up (\cite{Fro96}) and mass-loss. We are presently studying the effects of hot bottom burning (HBB) on intermediate mass AGB stars during their thermally pulsing evolution. During these calculations we found a new kind of thermal pulse which we report on in this paper. | We believe that the development of degenerate thermal pulses is inevitable provided two criteria are met:~1)~dredge-up is very deep, providing cooling of the helium shell soon after ignition;~and~2)~the star lives long enough on the AGB for the helium shell to reach (partially) degenerate conditions. Whether these conditions are realised in real stars remains an open question. The two aspects of AGB evolution which are the most uncertain are the extent of dredge-up and the rate of mass-loss (hence AGB lifetime), precisely the two phenomena which govern the occurrence or not of degenerate thermal pulses. It seems prudent to look for some observational consequence of degenerate thermal pulses. The nucleosynthetic consequences are currently being investigated, and will be reported elsewhere. \clearpage | 98 | 3 | astro-ph9803042_arXiv.txt |
9803 | astro-ph9803274_arXiv.txt | Deep $J, H,$ and $K$ images are used to probe the evolved stellar contents in the central regions of the Sculptor group galaxies NGC55, NGC300, and NGC7793. The brightest stars are massive red supergiants (RSGs) with $K \sim 15 - 15.5$. The peak RSG brightness is constant to within $\sim 0.5$ mag in $K$, suggesting that NGC55, NGC300, and NGC7793 are at comparable distances. Comparisons with bright RSGs in the Magellanic Clouds indicate that the difference in distance modulus with respect to the LMC is $\Delta \mu = 7.5$. A rich population of asymptotic giant branch (AGB) stars, which isochrones indicate have ages between 0.1 and 10 Gyr, dominates the $(K, J-K)$ color-magnitude diagram (CMD) of each galaxy. The detection of significant numbers of AGB stars with ages near 10 Gyr indicates that the disks of these galaxies contain an underlying old population. The CMDs and luminosity functions reveal significant galaxy-to-galaxy variations in stellar content. Star-forming activity in the central arcmin of NGC300 has been suppressed for the past Gyr with respect to disk fields at larger radii. Nevertheless, comparisons between fields within each galaxy indicate that star-forming activity during intermediate epochs was coherent on spatial scales of a kpc or more. A large cluster of stars, which isochrones suggest has an age near 100 Myr, is seen in one of the NGC55 fields. The luminosity function of the brightest stars in this cluster is flat, as expected if a linear luminosity-core mass relation is present. | Recent studies of the structural characteristics of spiral galaxies suggest that the disk and spheroid do not evolve in isolation, but interact throughout the lifetime of a galaxy (e.g. Andredakis, Peletier, \& Balcells 1995). The observational signatures of these interactions are readily apparent in late-type spiral galaxies. For example, Courteau, de Jong, \& Broeils (1996) find a scale-free Hubble sequence among late-type spirals, a result which could be explained if the central light concentrations \footnote{ The presence of a traditional bulge in late-type spirals galaxies has been challenged by Bothun (1992) and Regan \& Vogt (1994), who concluded that the `bulge' in M33 is actually the central extension of the halo. However, Minniti, Olszewski, \& Rieke (1993) and Mighell \& Rich (1995) resolved the innermost regions of M33 into stars, and detected a significant intermediate age population. Minniti {\it et al.} (1993) argue that, on the basis of stellar content alone, M33 contains a central component that is distinct from the halo. Given this debate, in the current paper the term `central light concentration' is used to refer to what has traditionally been called the `bulge' in late-type spirals.} formed after disks. WFPC1 images discussed by Phillips {\it et al.} (1996) reveal that the central light concentrations of late-type spirals are structurally distinct from the bulges of early-type spirals, suggesting differences in evolutionary histories. Surveys of the bright stellar content in nearby galaxies provide a direct means of studying the evolution of their central regions. Photometric studies of first ascent and asymptotic giant branch (AGB) populations are of particular interest, as these stars probe evolution during early and intermediate epochs, when the basic properties of the disk and spheroid were being imprinted. While efforts to resolve the central light concentrations of these systems into stars require near-diffraction limited image quality to overcome crowding, the inner disks of many nearby systems can easily be resolved into stars from the ground. Although the Local Group contains the closest, most obvious objects for stellar content surveys, the number of targets is limited to three morphologically diverse spiral galaxies: the Milky-Way, M31, and M33. This limited sample makes it necessary to study more distant systems and, as the nearest collection of galaxies outside the Local Group, the Sculptor Group offers a number of lucrative targets that can be resolved from the ground. In the current paper, deep $J, H$, and $K$ images are used to investigate the photometric properties of cool stars in the inner disks of the Sculptor galaxies NGC55, NGC300, and NGC7793. The morphological types and integrated brightnesses of these galaxies, as assigned by Sandage \& Tammann (1987), are summarized in Table 1. Throughout this study it is assumed that these galaxies are equidistant with $\mu_0 = 26.0$, as derived for NGC300 by van den Bergh (1992) from a number of different standard candles. The Galactic reddening towards these galaxies is negligible (Burstein \& Heiles 1984). There are a number of advantages to conducting photometric surveys of luminous, cool evolved stars at wavelengths longward of $1\mu$m. Not only is the contrast between bright cool stars and fainter unresolved objects in the disk enhanced at infrared wavelengths, but it is also possible to overcome the effects of line blanketing, which can affect the spectral energy distributions of moderately metal-rich giants at optical wavelengths (e.g. Bica, Barbuy, \& Ortolani 1991), and complicate efforts to derive bolometric corrections. In addition, near-infrared two-color diagrams can also be used to identify foreground stars, contamination from which may be significant at faint visible magnitudes. NGC55 and NGC300 have been the targets of earlier photometric investigations. Deep broad- and narrow-band surveys of NGC55 (Pritchet {\it et al.} 1987) and NGC300 (Richer, Pritchet, \& Crabtree 1985; Zijlstra, Minniti, \& Brewer 1996) have revealed that the outer disks of these galaxies contain rich AGB populations. A comparison of the AGB luminosity functions suggests that the star-forming histories of NGC55 and NGC300 during intermediate epochs were similar, but not identical (Pritchet {\it et al.} 1987). Freedman (1984) and Pierre \& Azzopardi (1988) used $B$ and $V$ photometry to survey the bright young stellar content of NGC300, while Kiszhurno-Kozrey (1988) used CCD observations to construct $(V, B-V)$ CMDs of two fields in NGC55. There is no published photometric study of the stellar content of NGC7793, although Catanzarite {\it et al.} (1995) report the discovery of Cepheids in this galaxy. The only published infrared survey of these galaxies was carried out by Humphreys \& Graham (1986), who obtained $JHK$ aperture measurements of red supergiant (RSG) candidates in NGC300. Spectroscopy revealed that almost half of the candidate objects were cool Galactic main sequence stars. The paper is structured as follows. The observations, reduction techniques, and methods used to measure stellar brightnesses are discussed in \S 2. The luminosity functions (LFs), two-color diagrams (TCDs), and color-magnitude diagrams (CMDs) derived from these data are presented and compared in \S 3. In \S 4 the data are used to search for radial population gradients in NGC300 and NGC7793. A summary of the results follows in \S 5. | 98 | 3 | astro-ph9803274_arXiv.txt |
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9803 | astro-ph9803104_arXiv.txt | We report the results of a 5-GHz southern-hemisphere snapshot VLBI observation of a sample of blazars. The observations were performed with the Southern Hemisphere VLBI Network plus the Shanghai station in 1993 May. Twenty-three flat-spectrum, radio-loud sources were imaged. These are the first VLBI images for 15 of the sources. Eight of the sources are EGRET ($>$~100~MeV) $\gamma$-ray sources. The milliarcsecond morphology shows a core-jet structure for 12 sources, and a single compact core for the remaining 11. No compact doubles were seen. Compared with other radio images at different epochs and/or different frequencies, 3 core-jet blazars show evidence of bent jets, and there is some evidence for superluminal motion in the cases of 2 blazars. The detailed descriptions for individual blazars are given. This is the second part of a survey: the first part was reported by Shen \etal (1997). | Blazar is the collective name for BL Lac objects, optically violent variables and highly polarized quasars, all of which share extreme observational properties that distinguish them from other active galactic nuclei. These properties include strong and rapid variability, high optical polarization, weak emission lines, and compact radio structure (cf. Impey 1992). About 200 blazars have been identified (cf. Burbidge \& Hewitt 1992). A possible explanation for the blazar phenomenon within a unified scheme for active galactic nuclei is that their emission is beamed by the relativistic motion of the jets traveling in a direction close to the observer's line of sight. This beaming argument is strengthened by the recent CGRO (Compton Gamma Ray Observatory) discovery that most of the detected high-latitude $\gamma$-ray sources are blazars (e.g. Dondi \& Ghisellini 1995). A comprehensive theoretical review of these sources has been made by Urry \& Padovani (1995). Blazars are an important class of active galactic nuclei because they are thought to be sources with relativistic jets seen nearly end-on. Such sources generally have very compact, flat-spectrum radio cores, which are appropriate for VLBI study. Pearson \& Readhead (1988) have undertaken a survey of a complete sample consisting of 65 strong northern-hemisphere radio sources. They provided the first well-defined morphological classification scheme, based primarily on the large-scale radio structure and radio spectra of the sources. Most surveys to date, however, including the recent Caltech--Jodrell Bank VLBI Surveys (Polatidis \etal 1995; Thakkar \etal 1995; Xu \etal 1995; Taylor \etal 1994; Henstock \etal 1995), have been restricted to northern-hemisphere sources. For example, all the confirmed superluminal radio sources, except the well-known equatorial source 3C~279 (1253$-$055), are in the northern sky (Vermeulen \& Cohen 1994). This reflects the paucity of southern VLBI observations, the notable exceptions being the systematic one-baseline surveys by Preston \etal (1985) and Morabito \etal (1986), and the more extensive SHEVE survey (Preston \etal 1989 and references therein). Since 1992 we have been carrying out a program to address this deficiency, using VLBI at 5 GHz to study southern radio sources. In an earlier paper (Shen \etal 1997, hereafter Paper I), we reported the results from the first observing session in 1992 November, and presented images of 20 strong sources selected on the basis of their correlated fluxes on intercontinental baselines. In 1993 May we observed a second sample of southern sources, which is the subject of this paper. Section~2 introduces this blazar sample; Section~3 briefly describes the observations and data reduction procedures; Section~4 presents the results; the summary and conclusions are presented in Section~5. Throughout the paper, we define the spectral index, $\alpha$, by the convention {S$_{\small \nu}\,\propto\,{\nu}^{\alpha}$}, and assume {H$_0=100$~km~s$^{-1}$~Mpc$^{-1}$} and q$_0=0.5$. | In this paper we have defined a sample of southern-hemisphere core-dominated blazars. Of the 24 blazars in the sample, 3 were observed earlier with the same array. The other 21 in the sample and 2 other sources were observed in 1993 May with the Southern VLBI Network plus the Shanghai radio telescope. This is part of the Southern Hemisphere 5-GHz VLBI Survey project, the aim of which is to improve the study of southern extragalactic radio sources (see Paper I). Our study also adds significantly to the number of sources whose structures can be compared on arcsecond (kpc) and milliarcsecond (pc) scales (Table~5). The misalignment of jet-like structures on these scales is an important unsolved problem for the understanding of compact sources. The main conclusions presented in this paper can be summarized as follows: \begin{enumerate} \item We have detected and imaged all 23 radio sources, of which 15 are first-epoch VLBI images. These are PKS~0118$-$272, PKS~0332$-$403, PKS~0426$-$380, PKS~0454$-$234, PKS~0823$-$223, PKS~1034$-$293, PKS~1244$-$255, PKS~1514$-$241, PKS~1936$-$155, PKS~1954$-$388, PKS~2005$-$489, PKS~2155$-$152, PKS~2240$-$260, PKS~2243$-$123 and PKS~2355$-$534. \item Most of the blazars are resolved and display simple morphology, with 12 having core-jet structures and 11 having single-core structures. Observations with increased sensitivity will probably reveal many more core-jet structures (e.g. 2243--123). We have compared our VLBI images with other radio images. Only 3 (PKS~0438$-$436, PKS~0537$-$441 and PKS~1226+023) of the 12 core-jet blazars were found to have curved jets. Superluminal motion was inferred from two-epoch observations for 2 sources (PKS~0208$-$512, PKS~2243$-$123). \item Eight of these blazars (PKS~0208$-$512, PKS~0454$-$234, PKS~0521$-$365, PKS~0537$-$441, PKS~1127$-$145, PKS~1226+023, PKS~1424$-$418 and PKS~2005$-$489) have been detected at $>100$~MeV $\gamma$-ray energies. Together with the other 5 EGRET sources observed in 1992 November (Paper I), a total of 13 southern $\gamma$-ray-loud blazars have now been imaged by our survey project. A systematic study of the VLBI properties of these $\gamma$-ray blazars and comparison with other non-$\gamma$-ray sources will improve our understanding of the beaming characteristics in blazars and the properties of EGRET sources. \end{enumerate} | 98 | 3 | astro-ph9803104_arXiv.txt |
9803 | astro-ph9803332_arXiv.txt | A quantitative method is presented to compare observed and synthetic colour-magnitude diagrams (CMDs). The method is based on a $\chi^2$ merit function for a point $(c_i,m_i)$ in the observed CMD, which has a corresponding point in the simulated CMD within $n\sigma(c_i,m_i)$ of the error ellipse. The $\chi^2$ merit function is then combined with the Poisson merit function of the points for which no corresponding point was found within the $n\sigma(c_i,m_i)$ error ellipse boundary.\hfill\break Monte-Carlo simulations are presented to demonstrate the diagnostics obtained from the combined ($\chi^2$, Poisson) merit function through variation of different parameters in the stellar population synthesis tool. The simulations indicate that the merit function can potentially be used to reveal information about the initial mass function. Information about the star formation history of single stellar aggregates, such as open or globular clusters and possibly dwarf galaxies with a dominating stellar population, might not be reliable if one is dealing with a relatively small age range. | In the last decade the simulation of synthetic Hertzsprung-Russell and colour-magnitude diagrams (hereafter respectively referred to as HRDs and CMDs) has advanced at a rapid pace. It has been applied successfully in various studies of young clusters in the Magellanic Clouds (Chiosi \etal\ 1989; Bertelli \etal\ 1992; Han \etal\ 1996; Mould \etal\ 1997; Vallenari \etal\ 1990, 1992, and 1994ab), dwarf galaxies in the Local Group (Aparicio{\muspc\&\muspc}Gallart 1995, 1996; Aparicio \etal\ 1996, 1997{\mmuspc}a,b; Ferraro \etal\ 1989; Gallart \etal\ 1996a{\to}c; Han \etal~1997; Tolstoy~1995, 1996; Tosi \etal~1991), open clusters in our Galaxy (Aparicio \etal~1990; Carraro \etal~1993, 1994; Gozzoli \etal~1996), and the structure of our Galaxy (Bertelli \etal~1994, 1995, 1996; Ng 1994, 1997{\mmuspc}a,b; Ng{\muspc\&\muspc}Bertelli 1996{\mmuspc}a,b; Ng \etal~1995, 1996{\mmuspc}a,b, 1997). \par Generally all studies focus in the first place on matching the morphological structures at different regions in the CMDs (Gallart 1998). In the recent years a good similarity is obtained between the observed and the simulated CMDs. Unfortunately, the best fit is in some cases distinguished by eye. The morphological differences are large enough to do this and the eye is actually guided by a detailed knowledge of stellar evolutionary tracks. However, the stellar population technique has improved considerably and an objective evaluation tool is needed, to distinguish quantitatively one model from another. \par Bertelli \etal~(1992,\muspc1995) and Gallart \etal~(1996c) defined ratios to distinguish the contribution from different groups of stars. The ratios are defined so that they are sensitive to the age, the strength of the star formation burst and/or the slope of the initial mass function. Vallenari \etal~(1996{\mmuspc}a,b) demonstrated that this is a good method to map the spatial progression of the star formation in the Large Magellanic Cloud. Robin \etal~(1996), Han \etal~(1997) and Mould \etal~(1997) use a maximum likelihood method to find the best model parameter(s), while Chen (1996{\mmuspc}a,b) adopted a multivariate analysis technique. Different models can be quantitatively sampled through Bayesian inference (Tolstoy 1995; Tolstoy{\muspc\&\muspc}Saha 1996) or a chi-squared test (Dolphin 1997). \par In principle one aims with stellar population synthesis to generate a CMD which is identical to the observed one. The input parameters reveal the evolutionary status of the stellar aggregate under study. To obtain a good similarity between the observed and the simulated CMD one needs to implement in the model in the first place adequately extinction, photometric errors and crowding. It goes without saying that the synthetic population ought to be comparable with the age and metallicity (spread) of the stellar aggregate. Only with a proper choice of these parameters, one can start to study in more detail the stellar initial mass function and the star formation history for an aggregate. \par This paper describes a quantitative evaluation method based on the combination of a chi-squared and Poisson merit function% \footnote{The method introduced in Sect.~2 actually mimics closely the procedure used in the `fit by eye' method.}. It allows one to select the best model from a series of models. In the next section this method is explained and Monte-Carlo simulations are made, to display how the diagnostic diagrams of the residual points can be employed. It is demonstrated that the non-fitting, residual points provide a hint about the parameter that needs adjustment in order to improve the model. This paper ends with a discussion about the method and the diagnostic diagrams together with their limitations. \hfill\break It is emphasized that this paper deals with a description of a quantitative evaluation method for CMDs. Aspects related with the implementation of an automated CMD fitting program and comparison with simulated or real data sets will not be considered, because they are not relevant for the general validity of the method described in the following section. | A method based on the $\chi^2$ merit function is presented to compare `observed' with synthetic single stellar populations. Monte-Carlo simulations have been performed to display the diagnostic power from a CMD containing the points for which no corresponding synthetic point was found within a reasonable error ellipse. The simulations indicate that the CMDs of residual points might provide hints about model parameters to be improved. The simulations further indicate that one ought to be cautious with the analysis of stellar luminosity functions and that strong hints can be obtained about the shape of the initial mass function. | 98 | 3 | astro-ph9803332_arXiv.txt |
9803 | astro-ph9803326_arXiv.txt | We present WFPC2--HST photometry of the resolved stellar population in the post-starburst galaxy NGC~1569. The color-magnitude diagram (CMD) derived in the F439W and F555W photometric bands contains $\sim$2800 stars with photometric error $\leq$ 0.2 mag down to \mb, \mv~$\simeq$~26, and is complete for \mv~$\lsim$~23. Adopting the literature distance modulus and reddening, our CMD samples stars more massive than $\sim$4~\MSUN, allowing us to study the star formation (SF) history over the last $\sim$0.15~Gyr. The data are interpreted using theoretical simulations based on stellar evolutionary models. The synthetic diagrams include photometric errors and incompleteness factors. Testing various sets of tracks, we find that the ability of the models to reproduce the observed features in the CMD is strictly related to the shape of the blue loops of the sequences with masses around 5~\MSUN. The field of NGC~1569 has experienced a global burst of star formation of duration $\gsim$0.1~Gyr, ending $\sim$ 5$-$10~Myr ago. During the burst, the SF rate was approximately constant, and, if quiescent periods occurred, they lasted less than $\sim$10~Myr. The level of the SF rate was very high: for a single-slope initial mass function (IMF) ranging from 0.1 to 120~\MSUN\ we find values of 3, 1, and 0.5~\Myr\ for $\alpha = 3$, 2.6, and 2.35 (Salpeter), respectively. When scaled for the surveyed area, these rates are approximately 100 times larger than found in the most active dwarf irregulars in the Local Group. The data are consistent with a Salpeter IMF, though our best models indicate slightly steeper exponents. We discuss the implications of our results in the general context of the evolution of dwarf galaxies. | Understanding the status and evolution of dwarf galaxies is of crucial importance in current astrophysics and cosmology. They are important ingredients in common scenarios of galaxy formation, either as building blocks or as left-overs of the formation process (Silk 1987). In addition, dwarf galaxies forming at redshift around and below 1 have been suggested to be responsible for the excess of faint blue galaxies seen in deep photometric surveys (Babul $\&$ Rees 1992; Babul $\&$ Ferguson 1996). Dwarf galaxies come in three main classes: (i) dwarf spheroidals, which contain population~II stars and have low surface brightness (e.g., Wirth \& Gallagher 1984); (ii) dwarf irregulars, with various levels of star-formation activity and a high gas mass (Hunter \& Gallagher 1986); (iii) starbursting dwarfs, including blue compact dwarfs (Thuan 1991), H~II galaxies (Terlevich \etal\ 1991), and blue amorphous galaxies (Sandage \& Brucato 1979). The boundary between dwarf and giant galaxies is usually taken to be around $M_{\rm{B}} = -16$ mag (Tammann 1994). An intriguing property of dwarf galaxies is the pronounced dichotomy between the two main classes of dwarf spheroidals and dwarf irregulars: most dwarfs are either gas-poor spheroids or gas-rich disk-like systems, with very few transition objects (van den Bergh 1977). The relationship between the two classes is still poorly understood: attempts to unify the complex zoo of dwarf galaxies have not been fully successful (Binggeli 1994). Some effort has been made to understand the interrelationship between individual classes in terms of an evolutionary sequence (e.g., Gallagher, Hunter, \& Tutukov 1984). Starbursting dwarfs can play a key role: a powerful starburst might strip a gas-rich dwarf irregular of its gas and transform it into a gas-poor dwarf-elliptical (Kormendy 1985; Marlowe 1997). Understanding the star-formation histories in starbursting dwarf galaxies is a prerequisite for understanding evolutionary connections between different object classes (e.g., Gallagher 1996) since during cosmologically brief periods of time their star formation increases dramatically. This has profound implications for the interstellar medium. Stellar winds and supernovae inject energy and may cause a `blow-out', initiating a galactic superwind (Heckman 1995). Numerous examples of galactic superwinds are known from observations (e.g., Heckman, Armus, \& Miley 1990) and their cosmological implications have been discussed (e.g., De Young \& Heckman 1994). However, attempts to relate the superwind properties and the star-formation histories in individual galaxies are quite rare and do not go beyond a qualitative level (Heckman \etal\ 1990; Leitherer, Robert, \& Drissen 1992). Yet establishing such a relation is crucial for theoretical models of the superwind hydrodynamics (Suchkov \etal\ 1994; Tenorio-Tagle \& Mu\~noz-Tu\~non 1997). NGC~1569 (= UGC~3056 = Arp~210 = VII~Zw~16 = IRAS~4260+6444) has been indicated by Gallagher, Hunter, $\&$ Tutukov (1984) as an outstanding object in their sample of active star forming galaxies. Indeed, it is a prime candidate for a detailed study of the star-formation history in a starburst galaxy since at $D = 2.2 \pm 0.6$~Mpc it is the closest starburst galaxy known (Israel 1988a). Adopting the literature reddening $E(B-V)$ = 0.56 (Israel 1988a), photometry of individual stars is feasible with HST down to $M_{V,0}\simeq -1.5$. This corresponds to stars with mass of about 3~\Ms\ in their core-helium burning phase, and consequently to a look-back time of $\sim$0.4~Gyr (e.g., Schaller \etal\ 1992). The properties of NGC~1569 are typical of a dwarf galaxy. With a distance modulus of $(m-M)_0$ = 26.71, its total absolute B magnitude is $M_{B.0} \sim -17$ (Israel 1988a), which is intermediate between the two Magellanic Clouds. Its total mass and hydrogen content are estimated $M \simeq 3.3 \times 10^8$~\Ms\ and $M_H \simeq 1.3 \times 10^8$~\Ms, respectively (Israel 1988a). Numerous studies of the chemical composition exist. The published range of oxygen abundances is very narrow: $12 + \log (O/H) = 8.25$ (Hunter, Gallagher, \& Rautenkranz 1982), 8.37 (Calzetti, Kinney, \& Storchi-Bergmann 1994), 8.26 (Devost, Roy, \& Drissen 1997), 8.29 (Gonz\'alez-Delgado et al. 1997), and 8.19 (Kobulnicky \& Skillman 1997; Martin 1997). Searches for chemical composition gradients were done by Devost et al. and Kobulnicky \& Skillman. No significant evidence for chemical inhomogeneities was found. If [O/Fe]~=~0.0 is assumed, the average metallicity of NGC~1569 is $Z \simeq 0.25$~\Zsun, with about 0.2~dex uncertainty. Thus, NGC~1569 is a gas-rich system with SMC-like composition in a relatively early stage of its chemical evolution. The most detailed study of the warm ($\sim$10$^4$~K) gas of NGC~1569 was done by Waller (1991), who found evidence for a link between the star-formation history and the morphology and kinematics of extended gas structures. The same connection is suggested by hot ($\sim$10$^7$~K) gas observed with the ROSAT and ASCA satellites (Heckman \etal\ 1995; Della Ceca \etal\ 1996). The extended X-ray emission is consistent with a starburst driven galactic superwind which could in principle lead to a large-scale disruption of the interstellar medium. Evidence for large numbers of supernovae was also found from the non-thermal radio spectrum (Israel 1988b). One of the most spectacular properties of NGC~1569 are its `super star clusters'. These high-density star clusters were first detected and discussed by Arp \& Sandage (1985) and Melnick \& Moles, \& Terlevich (1985) and later studied by O'Connell, Gallagher, \& Hunter (1994), Ho \& Filippenko (1996), De~Marchi \etal\ (1997), and Gonz\'alez-Delgado \etal\ (1997). The super star clusters reflect a recent starburst in which at least $10^5$~\Ms\ of gas was transformed into stars within about 1~pc. Super star clusters are similar to young Galactic globular clusters observed shortly after their formation (Meurer 1995). Pre-Costar HST studies of the field population were done by O'Connell \etal\ (1994) and Vallenari \& Bomans (1996) (hereinafter VB). According to the latter authors, NGC~1569 experienced a strong starburst until a few Myr ago, when its activity subsided. The repaired HST offers the possibility of significant improvement over these previous studies. We have therefore embarked on an extensive spectroscopic and photometric HST study in an attempt to confront the observational data with the predictions of theoretical models, and derive detailed informations on the star formation (SF) history in NGC~1569. In a first paper (De Marchi \etal\ 1997) we discussed the photometry of the SSCs, while in this paper we present the results for the resolved stars in the same field, which are interpreted in terms of the recent SF history via theoretical simulations. The data and their analysis are described in Section~2. The derived color-magnitude diagram and luminosity function are in Section~3. Our theoretical models are described in Section~4. They are compared with the observations in Section~5. In Section~6 we discuss our results and their implications for the evolution of dwarf galaxies, and we present in Section~7 our general conclusions. | In this paper we have presented HST WFPC2 F555W and F439W photometry of resolved stars in the dwarf irregular galaxy NGC~1569. About 2800 stars were measured with photometric error smaller than \sigmada = 0.2\,mag out of a total of $\sim$7000 objects detected. The corresponding CMD extends down to \mb, \mv~$\simeq$~26, and is complete for \mv~$\lsim$~23. It samples stars more massive than $\sim$4~\Ms\, allowing us to derive the SF history over the last 0.15~Gyr. The interpretation of the data has been performed via theoretical simulations, based on stellar evolutionary tracks. We have considered several sets of stellar tracks from the Geneva and Padova groups, and were able to reach a satisfactory representation of the data adopting either the Geneva tracks with $Z$~=~0.001, or the Padova tracks with $Z$~=~0.004. Since the applicability of these tracks is related to the shape of the blue loops, which is sensitive to details in the input parameters used in the computations, we consider it premature to dismiss alternative sets of evolutionary tracks. The two sets provide a similar picture of the recent SF history in the field of NGC~1569. The galaxy has experienced a general burst of SF, of duration $\Delta t \gsim 0.1$~Gyr, at an approximately constant rate. The burst seems to have stopped $\sim$5~--~10~Myr ago but SF continues in the HII regions and in the SSCs. If this general burst consists of successive episodes, these must have occurred at a similar rate, and be separated by short quiescent periods. Qualitatively, this behavior looks similar to that inferred for dwarf irregulars in the Local Group. Quantitatively, though, the level of the SFR in this recent burst in NGC~1569 is approximately 2 orders of magnitude higher. We find that the Salpeter IMF is consistent with the observed CMD and LF, but our best simulations are characterized by slightly steeper slopes (2.6 -- 3). This seems to disagree with observations of other starbursts which generally indicate a Salpeter slope (Leitherer 1997). The present study, however, refers to the field population of stars less massive than $\sim$30~\Ms, whereas most starburst IMFs are derived for the nuclear burst and for dense OB clusters in the mass range above $\sim$20~\Ms. Since most of the SF activity in the NGC~1569 field has subsided about 7~--~10~Myr ago, we have no direct information on the IMF of those stars with lifetimes less than this value. The corresponding minimum masses are around 30~\Ms. The field-star IMF in our Galaxy and in the Magellanic Clouds is steeper (biased against massive stars) than the cluster IMF (Massey, Johnson, $\&$ DeGioia-Eastwood 1995a; Massey \etal\ 1995b), with the cluster value being close to Salpeter. This could be due to a richness effect, where the most massive stars are not observed simply because of their small expected numbers. NGC~1569 may be another case where the most massive stars with masses above $\sim$30~\Ms\ are preferentially formed in clusters, but we emphasize the different mass range sampled by our and that of LMC/SMC studies. Our preferred IMF is quite similar to that derived for field stars of comparable mass in our Galaxy. This is in agreement with the notion of a virtually constant IMF, at least as far as the shape is concerned, and does not support the expectations of flatter slopes in starbursting regions (e.g., Padoan, Nordlund, $\&$ Jones 1997), or in low metallicity environments, as sometimes invoked to better reproduce the properties of elliptical galaxies (e.g., Vazdekis \etal\ 1996). In the past, the SF in NGC~1569 is likely to have proceeded at a substantially lower average rate than in the recent burst. This follows from the estimate of the gas exhaustion timescale, and from the relatively low metallicity of the ISM in this galaxy. This leaves the possibilities of either an approximately constant SFR, or short strong bursts, with long interburst periods. Detailed chemical evolution models are needed to explore quantitatively the possible evolutionary paths for NGC~1569, and this will be the subject of a forthcoming paper. The recent burst of SF is certainly not the first episode in NGC~1569 (see also VB) but may be the last. The velocity field of the ionized gas (Tomita, Ohta, $\&$ Saito 1994) and energetics arguments support this possibility (Heckman \etal\ 1995). If this was the case, in the future NGC~1569 may turn into a dwarf spheroidal galaxy. In addition, its luminosity evolution would be characterized by passive fading, which seems required by the Babul $\&$ Ferguson (1996) bursting dwarfs model to explain the excess faint galaxy counts. However, it is also possible that the outflowing gas falls back onto the galaxy, triggering a successive burst of SF, an option which has to be explored with the computation of detailed hydrodynamic modelling for bursting dwarf galaxies (D'Ercole $\&$ Brighenti 1998). In spite of these uncertainties, our results shows that dwarf galaxies are capable of sustaining bursts of SF at large enough rates to be relevant for the interpretation of the faint galaxy counts. | 98 | 3 | astro-ph9803326_arXiv.txt |
9803 | astro-ph9803110_arXiv.txt | The velocity distribution $f(\bvel)$ of nearby stars is estimated, via a maximum-likelihood algorithm, from the positions and tangential velocities of a kinematically unbiased sample of 14\,369 stars observed by the \hip\ satellite. $f$ shows rich structure in the radial and azimuthal motions, $v_R$ and $v_\varphi$, but not in the vertical velocity, $v_z$: there are four prominent and many smaller maxima, many of which correspond to well known moving groups. While samples of early-type stars are dominated by these maxima, also up to about a quarter of red main-sequence stars are associated with them. These moving groups are responsible for the vertex deviation measured even for samples of late-type stars; they appear more frequently for ever redder samples; and as a whole they follow an asymmetric-drift relation, in the sense that those only present in red samples predominantly have large $|v_R|$ and lag in $v_\varphi$ w.r.t.\ the local standard of rest (LSR). The question arise, how these old moving groups got on their eccentric orbits? A plausible mechanism known from the solar system dynamics which is able to manage a shift in orbit space is sketched. This mechanism involves locking into an orbital resonance; in this respect is intriguing that Oort's constants, as derived from \hip\ data, imply a frequency ratio between azimuthal and radial motion of exactly $\Omega:\kappa=3:4$. Apart from these moving groups, there is a smooth background distribution, akin to Schwarzschild's ellipsoidal model, with axis ratios $\sigma_R: \sigma_\varphi:\sigma_z\approx1:0.6:0.35$. The contours are aligned with the $v_r$ direction, but not w.r.t.\ the $v_\varphi$ and $v_z$ axes: the mean $v_z$ increases for stars rotating faster than the LSR. This effect can be explained by the stellar warp of the Galactic disk. If this explanation is correct, the warp's inner edge must not be within the solar circle, while its pattern rotates with frequency $\gtrsim13\kmskpc$ retrograde w.r.t.\ the stellar orbits. | \else | \else | 98 | 3 | astro-ph9803110_arXiv.txt |
9803 | astro-ph9803260_arXiv.txt | There are now about fifty known radio pulsars in binary systems, including at least five in double neutron star binaries. In some cases, the stellar masses can be directly determined from measurements of relativistic orbital effects. In others, only an indirect or statistical estimate of the masses is possible. We review the general problem of mass measurement in radio pulsar binaries, and critically discuss all current estimates of the masses of radio pulsars and their companions. We find that significant constraints exist on the masses of twenty-one radio pulsars, and on five neutron star companions of radio pulsars. All the measurements are consistent with a remarkably narrow underlying gaussian mass distribution, $m=1.35\pm0.04M_\odot$. There is no evidence that extensive mass accretion ($\Delta m\gsim0.1M_\odot$) has occurred in these systems. We also show that the observed inclinations of millisecond pulsar binaries are consistent with a random distribution, and thus find no evidence for either alignment or counteralignment of millisecond pulsar magnetic fields. | Neutron stars have been the subject of considerable theoretical investigation since long before they were discovered as astronomical sources of radio and X-ray emission (\cite{bz34b,ov39,whe66}). Their properties are determined by the interplay of all four known fundamental forces---electromagnetism, gravitation, and the strong and weak nuclear forces---but neutron stars remain sufficiently simple in their internal structure that realistic stellar modeling can be done. Measurements of their masses and radii (as well as detailed study of their cooling histories and rotational instabilities) provide a unique window on the behavior of matter at densities well above that found in atomic nuclei ($\rho_{\rm nuc}\approx 2.8\times10^{14}\mbox{g~cm}^{-3}$). Observations of neutron stars also provide our only current probe of general relativity (GR) in the ``strong-field'' regime, where gravitational self-energy contributes significantly to the stellar mass. The most precisely measured physical parameter of any pulsar is its spin frequency. The frequencies of the fastest observed pulsars (PSR~B1937+21 at 641.9~Hz and B1957+20 at 622.1~Hz) have already been used to set constraints on the nuclear equation of state at high densities (e.g., \cite{fidp88}) under the assumption that these pulsars are near their maximum (breakup) spin frequency. However, the fastest observed spin frequencies may be limited by complex accretion physics rather than fundamental nuclear and gravitational physics. A quantity more directly useful for comparison with physical theories is the neutron star mass. The basis of most neutron star mass estimates is the analysis of binary motion. Soon after the discovery of the first binary radio pulsar (\cite{ht75a}), it became clear that the measurement of relativistic orbital effects allowed extremely precise mass estimates. Indeed, the measurement uncertainties in several cases now exceed in precision our knowledge of Newton's constant $G$, requiring masses to be quoted in solar units $GM_\odot$ rather than kilograms if full accuracy is to be retained. After several recent pulsar surveys, there are now about fifty known binary radio pulsar systems, of which five or six are thought to contain two neutron stars. It is thus possible for the first time to consider compiling a statistically significant sample of neutron star masses. It is our purpose here to provide a general, critical review of all current estimates of stellar masses in radio pulsar binaries. The resulting catalog, with a careful, uniform approach to measurement and systematic uncertainties, should be of value both to those who wish to apply mass measurements to studies of nuclear physics, GR, and stellar evolution, and as a guide to the critical observations for observational pulsar astronomers. We begin with a discussion of known methods for pulsar mass determination (\S\ref{sec:methods}), including a new statistical technique for estimating the masses of millisecond pulsars in non-relativistic systems. In \S\ref{sec:estimates} we review all known mass estimates, including new data and analysis where possible. Statistical analysis of the available pulsar mass measurements is presented in \S\ref{sec:stat}. We summarize in \S\ref{sec:summ}. A second paper will consider mass estimates for neutron stars in X-ray binary systems (\cite[Paper~II]{ct98}). A detailed discussion of the implications of the combined results of this work and Paper~II for studies of supernovae and neutron star formation, mass transfer in binary evolution, the nuclear equation of state, and GR will occur elsewhere (Paper~III). | \label{sec:stat} For a dozen neutron stars, useful mass constraints are available with no assumptions beyond the applicability of the general relativistic equations of orbital motion to binary pulsar systems. Ten of these stars are members of double neutron star binaries. With the possible exception of PSR~B2127+11C, in the globular cluster M15, the pulsar in each system is believed to have undergone a short period of mass accretion during a high-mass X-ray binary phase ($\Delta m\sim10^{-3}M_\odot$, Taam and van den Heuvel 1986). The companion stars have not undergone accretion; their masses most directly preserve information about the initial mass function of neutron stars. Only two ``millisecond'' pulsars, the end products of extended mass transfer in low-mass X-ray binaries, have interesting mass estimates based on GR alone: PSRs~B1802$-$07 and B1855+09. Because such pulsars must accrete $\sim0.1M_\odot$ to reach millisecond periods (\cite{tv86}), and much more ($\sim0.7M_\odot$) in some field decay models (e.g., \cite{vb95a}), obtaining additional mass measurements of millisecond pulsars is of particular interest in testing evolutionary models and in locating the maximum neutron star mass. As noted in \S\ref{sec:cmrr}, the $P_b$--$m_2$ relation can be used to estimate the companion mass in recycled binary systems with circular orbits and orbital periods $P_b\gtrsim 3$~d. There are now thirteen such millisecond ($P_{\rm spin}<10$~ms) pulsars known, excluding those in globular clusters (where gravitational interactions may have significantly perturbed the orbital parameters since spin-up). In each case, the measured mass function and the inferred companion mass, together with the requirement that $\sin i<1$, then yields an upper limit on the mass of the pulsar itself. A number of systems in which this upper limit is particularly constraining have been mentioned in \S\ref{sec:nswd}. Additional constraints on the neutron star mass in these systems can be derived using statistical arguments, given a prior assumption about the distribution of binary inclinations. The simplest such assumption is that the binaries are randomly oriented on the sky, though biases toward high or low inclinations are possible in some models (\S\ref{sec:uniform}). However, as discussed below, we believe there is currently no evidence for such a bias, so for the remainder of this discussion we assume random orbital orientations. For an individual system, we are interested in the probability distribution\footnote{We adopt the notation that $p(x ; A)$ is the (marginal) probability density for the random variable $x$, where $x$ depends upon the parameter $A$. Also, $p(x|y; A)$ is the conditional probability density for the random variable $x$ for a given value of the random variable $y$ and parameter $A$.} $p(m_1; f, P_b)$ for the neutron star mass $m_1$ given the measured mass function $f$ and binary period $P_b$. We can neglect the measurement uncertainty in $f$ and $P_b$. Then, the probability distribution for $m_1$ can be written schematically as \begin{equation} p(m_1; f, P_b) = \int_0^1 d(\cos i) \int_{m_{2,{\rm min}}(P_b)}^{m_{2,{\rm max}}(P_b)} dm_2\, p(m_2; P_b)\, p(\cos i)\, p(m_1|m_2,\cos i; f) , \end{equation} where $m_1=f^{-1/2}(m_2 \sin i)^{3/2}-m_2$ is restricted to positive values. We have evaluated this numerically for each system, assuming that $p(\cos i)$ is uniform between zero and unity and that $p(m_2; P_b)$ is uniformly distributed within the appropriate factor (see \S\ref{sec:cmrr}) of the $m_2$ implied by equations (9)--(11). Not surprisingly, the width of the distribution $p(m_1; f, P_b)$ is dominated by the range of allowed $\cos i$ rather than the uncertainty in $m_2$ for a given $P_b$. For each of the 13 binaries, we have plotted the cumulative distribution CDF$(m_1)=\int_0^{m_1} p(m_1')\, dm_1'$ in Figure~\ref{fig:cumprob}.\placefigure{fig:cumprob} The median and 68\% and 95\% confidence regions for each pulsar mass is given in Table~\ref{tab:pbm2}. \placetable{tab:pbm2} Although several of the pulsars have, under the assumptions made, most likely masses well above $2M_\odot$, some such results are expected even if all the masses are quite low. In fact, in only one case of the 13 pulsars does $1.35M_\odot$ lie outside the 95\% central confidence region (J1045$-$4509), and in 6 cases of 13 is $1.35M_\odot$ excluded at 68\% confidence, consistent with chance. It is interesting to ask whether a single, simple distribution of neutron star masses is consistent with all of our observational constraints. We considered two models for this question: a Gaussian distribution of masses with mean $\hat{m}$ and standard deviation $\sigma$, and a uniform distribution of masses between $m_l$ and $m_u$ ({\it cf.} Finn 1994). A maximum likelihood analysis was used to estimate the parameters $\hat{m}$, $\sigma$, $m_l$, and $m_u$ (assuming a uniform prior distribution for all four parameters). The resulting 68\% and 95\% joint confidence limits on $\hat{m}$ and $\sigma$ are shown in Figure~\ref{fig:all26}, and on $m_l$ and $m_h$ in Figure~\ref{fig:all26uni}. \placefigure{fig:all26} \placefigure{fig:all26uni} In each model, the distribution of neutron star masses is remarkably narrow: the maximum likelihood solutions are $\hat{m}=1.35M_\odot$ and $\sigma=0.04M_\odot$, and $m_l=1.26M_\odot$ and $m_u=1.45M_\odot$. Of course, {\em any} model (even a poor one) will yield maximum likelihood parameters for a given data set. However, it is obvious by inspection that both the Gaussian and uniform distributions for the neutron star mass are good fits to the extremely narrow observed range of neutron star masses in the double neutron star binaries. While it is difficult to quantify the goodness-of-fit for the entire data set, because of the diverse assumptions made in the various mass estimates and the sometimes highly non-gaussian error estimates, we can easily test some neutron star subsamples against the maximum likelihood gaussian model $m_1=1.35\pm0.04M_\odot$. For the thirteen neutron-star--white-dwarf binaries, we used a Monte Carlo technique to evaluate the fit quality. For each binary (with its measured $P_b$ and $f$), we simulated a large number of Monte Carlo trials where the neutron star mass $m_1$ was drawn from the maximum likelihood model, $m_2$ was drawn from the appropriate uniform distribution implied by $P_b$, and $\cos i$ was drawn from a uniform distribution. The Monte Carlo trials were then used to construct the probability distribution for the mass function, and this distribution was used to compute the cumulative probability for the measured mass function, $p(f^\prime<f)$. If the model and the associated assumptions are correct, then the cumulative probabilities for the 13 measured mass functions should be consistent with a uniform distribution betwen zero and unity ({\it cf.} a classical $V/V_{\rm max}$ test, Schmidt 1968). A KS test of the distribution shows consistency with a uniform distribution at the improbably good 99\% level (Figure~\ref{fig:ks}). \placefigure{fig:ks} For the ten stars for which gaussian error estimates $\sigma_e$ are available (both stars in the relativistic binaries B1534+12, B1913+16, and B2127+11C, as well as J1012+5307, J1713+0747, B1855+09, and J0045$-$7319), we can calculate a $\chi^2$ statistic, $\sum(m-\hat{m})^2/(\sigma^2+\sigma_e^2)=7.5$, consistent with expectations for a chi-square distribution with $10-2$ degrees of freedom. We conclude, therefore, that at least in the radio pulsar systems, there is no evidence for neutron star masses above about $1.45M_\odot$. Indeed, the data appear very well modeled by very narrow distributions centered around $1.35M_\odot$. | 98 | 3 | astro-ph9803260_arXiv.txt |
9803 | astro-ph9803056_arXiv.txt | Soft gamma repeaters are astrophysical sources which exhibit long periods of quiescence, often spanning years, punctuated by periods of intense bursting activity during which many brief (durations $<1$ s) and intense (luminosities L$\sim1-10^{3}$ L$_{Edd}$) bursts are emitted by the source (Norris et al. 1991). Besides the burst emission, SGRs are also characterized by steady - also referred to as ``quiescent'' - emission. Believed to be neutron stars, the mechanism(s) for both the steady and bursting X--ray emission is still not well understood (Thompson \& Duncan 1995). SGR 1806-20 is the most prolific SGR, and it has been studied in the X--ray (Sonobe et al. 1994), optical (van Kerkwijk et al. 1995), infrared (Kulkarni et al. 1995), and radio (Kulkarni et al. 1994) bands. The source became active again during the Fall of 1996, emitting many powerful bursts that were first detected with BATSE (Kouveliotou et al. 1996). A target of opportunity observation by the {\it Rossi X--ray Timing Explorer} (RXTE) was initiated on November 5, 1996. The data analyzed here were taken during that $50$ ks observation, which spanned the time interval starting at 10:53:20 UT (5/11/96) and ending at 10:52:00 UT (6/11/96). In addition, followup {\it RXTE} data were taken during the time interval 06:45:20 UT (15/7/97) to 09:06:24 (15/7/97), in which the source was not bursting. | Through analysis of {\it RXTE} TOO data, we have determined that the persistent emission from SGR 1806--20 is consistent with a constant spectral shape and intensity both during and away from the active bursting periods. The spectrum is best-fit by a nonthermal power law shape, with thermal bremsstrahlung and Raymond-Smith functional forms producing much worse fits to the {\it RXTE} data. The mean power law spectral index obtained by {\it RXTE} is $2.30\pm0.02$, which is consistent with the {\it ASCA} value of $\alpha=2.2\pm0.2$ (Sonobe et al. 1994). The nonthermal nature of the SGR 1806--20 quiescent spectrum supports the idea that the X--ray emission is due to a compact synchrotron nebula, or plerion, that derives its power from either a rapidly spinning-down pulsar (Kulkarni et al. 1994) or energetic particles ejected in the SGR bursts (Tavani 1994). The lack of an iron line in the SGR spectrum (Sonobe 1994) and short term time variability argue against the source being a low luminosity X--ray binary system (White, Nagase, \& Parmar 1995), although a more extensive monitoring campaign is necessary to rule out this hypothesis. \paragraph*{Acknowledgements.} We thank NASA for support under grants NAS5-30720 (D.M. and R.E.R.), NAG5-2560 (S.D. and C.K.), and NAG5-4878 (JvP) | 98 | 3 | astro-ph9803056_arXiv.txt |
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9803 | nucl-th9803032_arXiv.txt | The nucleus $^{54}$Mn has been observed in cosmic rays. In astrophysical environments it is fully stripped of its atomic electrons and its decay is dominated by the $\beta^-$ branch to the $^{54}$Fe ground state. Application of $^{54}$Mn based chronometer to study the confinement of the iron group cosmic rays requires knowledge of the corresponding halflife, but its measurement is impossible at the present time. However, the branching ratio for the related $\beta^+$ decay of $^{54}$Mn was determined recently. We use the shell model with only a minimal truncation and calculate both $\beta^+$ and $\beta^-$ decay rates of $^{54}$Mn. Good agreement for the $\beta^+$ branch suggests that the calculated partial halflife of the $\beta^-$ decay, $(4.94\pm 0.06)\times 10^5$ years, should be reliable. However, this halflife is noticeably shorter than the range 1-$2\times 10^6$~y indicated by the fit based on the $^{54}$Mn abundance in cosmic rays. We also evaluate other known unique second forbidden $\beta$ decays from the nuclear $p$ and $sd$ shells ($^{10}$Be, $^{22}$Na, and two decay branches of $^{26}$Al) and show that the shell model can describe them with reasonable accuracy as well. | The nucleus $^{54}$Mn decays in the laboratory dominantly by electron capture to the $2^+$ state in $^{54}$Cr with the halflife of 312 days. However, as a component of cosmic rays, $^{54}$Mn will be fully stripped of its atomic electrons, and this mode of decay is therefore impossible. The $^{54}$Mn nuclei were in fact detected in cosmic rays using the Ulysses spacecraft~\cite{Ulysses1,Ulysses2}. They offer an attractive possibility to use their measured abundance as a chronometer for the iron group nuclei (Sc - Ni) in cosmic rays in analogy to the chronometers based on the abundances of other long lived isotopes ($^{10}$Be, $^{26}$Al, and $^{36}$Cl). With them one can, in turn, determine the mean density of interstellar matter, a quantity of considerable interest. The use of the long lived nuclei as cosmic ray chronometers is reviewed in Ref.~\cite{Simpson}. The importance of $^{54}$Mn for the understanding of propagation of the iron group nuclei that are products of explosive nuclear burning have been stressed in Refs.~\cite{Grove,Leske}. For this program to succeed, however, one must know the halflife of the stripped $^{54}$Mn ($I^{\pi} = 3^+$). The decay scheme of $^{54}$Mn is shown in Fig.~\ref{fig:decay}; the dashed lines indicate the decay paths of the stripped $^{54}$Mn. In two recent difficult and elegant experiments the very small branching ratio for the $\beta^+$ decay to the ground state of $^{54}$Cr has been measured: $(2.2 \pm 0.9) \times 10^{-9}$~\cite{b+1} and $(1.20 \pm 0.26) \times 10^{-9}$~\cite{b+2}. By taking the weighed mean of these values we extract the averaged branching ratio of $(1.28 \pm 0.25) \times 10^{-9}$. Combining it with the known halflife for $^{54}$Mn of 312.3(4)~d~\cite{nds54}, it corresponds to an experimental partial $\beta^+$ halflife of $(6.7 \pm 1.3) \times 10^8$ years. As explained in~\cite{Ulysses1,b+1,b+2} one expects, however, that the decay of the fully stripped $^{54}$Mn will be dominated by the at present unobservable $\beta^-$ decay to the $^{54}$Fe ground state. Previously, the partial $\beta^-$ halflife was estimated assuming that the $\beta^-$ and $\beta^+$ form factors are identical. Very recently, in Ref. \cite{b+2}, the ratio of the $\beta^-$ and $\beta^+$ form factors was calculated using a very truncated shell-model and extending it by comparison with similar calculations in the $sd$-shell. The estimated $\beta^-$ halflife is $(6.3 \pm 1.3) \times 10^5$~y \cite{b+2}. In this work we will use the state of the art shell model and evaluate not only the $EC$ decay rate of the normal $^{54}$Mn, but also both decay branches of the unique second forbidden transitions $^{54}$Mn$(3^+) \rightarrow {}^{54}$Cr$(0^+)$ and $^{54}$Mn$(3^+) \rightarrow {}^{54}$Fe$(0^+)$. By comparing the calculated $\beta^+$ decay halflife (or branching ratio) to the measured one we hope to judge the reliability of the calculation. We then proceed to calculate the halflife of the unknown $\beta^-$ decay. The decays of stripped $^{54}$Mn are unique second forbidden transitions which depend on a single nuclear form factor (matrix element). Halflives of several such decays in the $p$ shell ($^{10}$Be) and $sd$ shell ($^{22}$Na, and two decay branches of $^{26}$Al) are known and have been compared to the nuclear shell model predictions in Ref.~\cite{War}. For the $sd$ shell nuclei, however, only calculations in a severely truncated space were performed in~\cite{War}. Since that time computation techniques and programming skills have improved considerably. Thus, in order to further test our ability to describe this kind of weak decays, we repeat the analysis~\cite{War}, using the exact shell model calculations without truncation. At the same time the availability of new (and different) experimental data for the $^{10}$Be~\cite{be10exp} and $^{26}$Al~\cite{al26exp} decays make necessary a new comparison between experiment and calculations. In order to evaluate the decay rate we use the formulation of~\cite{tables}. The number of particles with momentum $p$ emitted per unit time is: \begin{equation} N(p_e) dp_e = \frac{g^2}{2 \pi^3} p_e^2 p_{\nu}^2 F(Z,W_e) C(W_e) dp_e, \end{equation} where $g$ is the weak coupling constant, $p_e$ and $W_e$ are electron (or positron) momentum and energy, $p_{\nu}$ is the neutrino (or antineutrino) momentum, and $Z$ is the atomic number of the daughter nucleus. All momenta and energies are in units where the electron mass is unity. For the Fermi function $F(Z,W_e)$ we use the tabulated values, and the shape factor $C(W_e)$ for the case of the unique second forbidden transitions is of the form \begin{equation} C(W_e) = \frac{R^4}{15^2} \left| ^AF_{321}^{(0)} \right|^2 \left[ p_\nu^4 + \frac{10}{3} \lambda_2 p_\nu^2 p_e^2 + \lambda_3 p_e^4 \right]. \end{equation} The nuclear form factor, in turn, is defined as \begin{equation} ^AF_{321}^{(0)} = g_A \sqrt{\frac{4\pi}{2J_i + 1}} \frac{\langle f || r^2 [\bbox{Y}_2 \times \bbox{\sigma}]^{[3]}\bbox{t}_\pm || i \rangle}{R^2}, \label{eq:formfactor} \end{equation} where $i$ denotes the initial state and $f$ the final one; the matrix element is reduced with respect to the spin space only (Racah convention~\cite{edmons}); $\pm$ refers to $\beta^\pm$ decay; $\bbox{t}_\pm = (\bbox{\tau}_x \pm i \bbox{\tau}_y)/2$, with $\bbox{t}_+ p = n$; $g_A = -1.2599 \pm 0.0025$~\cite{towhar}; and $R$ is the nuclear radius (the final expression for $C(W_e)$ is obviously independent of $R$). The functions $\lambda_2$ and $\lambda_3$ are tabulated in~\cite{tables}. Integrating the rate formula up to the spectrum endpoint we obtain the expression for $1/\tau$ and, respectively, for the halflife ($T_{1/2} = \ln(2) \tau$) in terms of the nuclear form factor squared. (For the stripped atoms we correct the endpoint energy accordingly.). For the quantity $2 \pi^3(\ln 2)/g^2$ we use the value $6146\pm 6$~s~\cite{towhar}. Note the usual $f t$ value, commonly used to characterize a decay, uses the integrated phase space factor $f$ of Eq. (1), however without the constant $g^2/(2 \pi^3 15^2)$) and the radius factor $R^4$. | Our shell model calculations reproduce the experimental halflives of the unique second forbidden beta decays within a factor of less than two. No clear evidence for the quenching of the corresponding form factors emerges. For the stripped $^{54}$Mn decays, the shell model describes the $\beta^+$ branch within errors. It predicts that the form factor for the $\beta^-$ decay is larger than the one for the $\beta^+$ decay. The calculated $\beta^-$ halflife (and therefore also the total halflife) is noticeably shorter than the range based on the observation of $^{54}$Mn in cosmic rays. This conflict, albeit relatively mild, makes attempts to determine the branching ratio for the $\beta^-$ decay experimentally even more compelling. | 98 | 3 | nucl-th9803032_arXiv.txt |
9803 | astro-ph9803285_arXiv.txt | Analysis of several recent ROSAT HRI observations of the gravitationally lensed system Q0957+561 has led to the detection at the 3$\sigma$ level of the cluster lens containing the primary galaxy G1. The total mass was estimated by applying the equation of hydrostatic equilibrium to the detected hot intracluster gas for a range of cluster core radii, cluster sizes and for different values of the Hubble constant. X-ray estimates of the lensing cluster mass provide a means to determine the cluster contribution to the deflection of rays originating from the quasar Q0957+561. The present mass estimates were used to evaluate the convergence parameter $\kappa$, the ratio of the local surface mass density of the cluster to the critical surface mass density for lensing. The convergence parameter, $\kappa$, calculated in the vicinity of the lensed images, was found to range between 0.07 and 0.21, depending on the assumed cluster core radius and cluster extent. This range of uncertainty in $\kappa$ does not include possible systematic errors arising from the estimation of the cluster temperature through the use of the cluster luminosity-temperature relation and the assumption of spherical symmetry of the cluster gas. Applying this range of values of $\kappa$ to the lensing model of Grogin \& Narayan (1996) for Q0957+561 but not accounting for uncertainties in that model yields a range of values for the Hubble constant: $ 67 <$ H$_{0}$ $<$ 82 $\,$ km $\,$ s$^{-1}$ $\,$ Mpc$^{-1}$, for a time delay of 1.1 years. | Refsdal (1964a 1964b 1966) outlined a method for determining a global value for the Hubble constant by measuring the light travel delay between different images of a gravitationally lensed system. Together with a model that describes the gravitational potential of the lens, the time delay defines the geometrical dimensions of the lens system and therefore may lead to the determination of the Hubble constant. The criteria that make a gravitational lens system (GLS) suitable for measuring the Hubble constant can be summarized as follows: (a) The sky positions and redshifts of the lens and of the images of the lensed object should be measurable accurately; (b) The mass distribution of the lens should be subject to accurate and reliable estimation. When the lens has multiple components, as for Q0957+561, this estimation becomes especially difficult and benefits from many different types of observations and also from many and extended images, such as rings. For example, if the types of observations include radio and optical emission lines, which are less sensitive to microlensing, flux ratio measurements are more useful in constraining the lens model; and (c) The differences in propagation time from the lensed object to us through (at least two) different images (``time delay'') should be measurable with high accuracy. The source, i.e., the lensed object, must therefore be variable on time scales far shorter than the time delays, to allow accurate estimates to be made of the latter. Several potential candidates for determining $H_{0}$ are the quasar Q0957+561 with two images and a measured time delay of 1.1 years (Pelt et al. 1996; Haarsma et al. 1997; Kundi\'{c} et al. 1997), the quadruple image system PG 1115+080 with a measured time delay between components B and C of about 25 days (Schechter et al. 1997; Bar-Kana 1997), and B0218+357 with a time delay of 12 $\pm$ 3 days (Geiger \& Schneider 1996; Corbett et al. 1996). Combining measured time delays with a detailed lensing model, Refsdal's suggestion has been applied to Q0957+561 (Falco, Gorenstein, \& Shapiro 1991; Grogin \& Narayan 1996, hearafter GN) and PG1115+080 (Keeton \& Kochanek 1997). In this paper we focus on a study of the properties of Q0957+561 (Walsh, Carswell \& Weyman 1979), one of the most extensively studied lensing systems which has been monitored since its discovery in both the optical and radio (see, for example, Schild \& Thomson 1995; Vanderriest et al. 1989; Haarsma et al. 1997). The primary lensing components are an elliptical cD galaxy at a redshift of 0.355, usually referred to as the G1 galaxy, a cluster of galaxies containing the G1 galaxy, and a group of galaxies at a redshift of 0.5. The most thorough optical spectroscopic and photometric study to date of the field around Q0957+561 has been presented by Angonin-Willaime, Soucail, \& Vanderriest (1994). One of the major concerns in modeling Q0957+561 is a mass sheet degeneracy which arises from the presence of the cluster of galaxies at z=0.36. In particular if one were to modify the assumed radial mass density profile $\kappa(\theta)$ of the lens to ${\lambda}{\kappa(\theta)}$ + (1 - $\lambda$), where 1 - $\lambda$ represents a constant surface mass density sheet term, then the new mass distribution would also satisfy the observational constraints. However the resulting value for $H_{0}$ would be scaled by the constant 1 - $\lambda$. Since $\lambda$ must always be positive one may always determine an upper limit on the Hubble constant for $\lambda$ = 0 (Falco et al. 1991). As also noted by Falco et al. (1991) and later GN, this degeneracy may be broken if one were to measure the velocity dispersion or total mass distribution of either the cluster of galaxies or the principal lensing galaxy G1. Recently Falco et al. (1997) have obtained a velocity dispersion for the central lens galaxy G1 of $\sigma_{v} = 266 \pm 12$ km s$^{-1}$ or $\sigma_{v} = 279 \pm 12$ km s$^{-1}$ depending on the interpretation of the spectroscopic slit data.\\ In this paper we focus on improving the cluster model and we present an estimate of the total mass of this lensing cluster of galaxies through X-ray measurements of thermal emission from the intracluster gas. In section 2 we describe the spatial analysis of the X-ray data of Q0957+561 obtained from several ROSAT HRI observations and we present an estimate for the convergence parameter $\kappa(\theta)$ of the cluster. Section 3 describes the spectroscopic analysis of ROSAT position sensitive proportional counter (PSPC) and ASCA Gas Imaging Spectrometer (GIS) X-ray observations at different epochs in order to investigate the variability of different spectral components, and presents various scenarios to explain the observed variability in the X-ray flux. Finally section 4 presents our conclusions from the analysis of the X-ray observations of the Q0957+561 GL system. In particular, we discuss how our results allow bounds to be placed on the Hubble constant, albeit with these bounds still subject to systematic errors that we cannot yet evaluate accurately. | The need for a detailed measurement of the mass distribution of the cluster lens has been emphasized by Falco et al.(1991) and by GN (1996) in analyses of lens models for Q0957+561 that incorporate the cluster contribution through a convergence parameter $\kappa(r)$. The most important result of the present analysis of the ROSAT HRI data of Q0957+561, that we hope will eventually lead to an accurate determination of the mass distribution of the cluster lens, is the 3$\sigma$ detection of X-rays from the cluster of galaxies that contains the principal lens galaxy G1. \\ As indicated in Figure 3, the uncertainty in our estimated values for the mass distribution of the cluster spans a large range as do the consequent estimates of the convergence parameter of the lens (Figure 6). The main contributors to this uncertainty are the lack of knowledge of the shape of the cluster, the core radius of the cluster, the temperature profile of the hot gas, the cluster extent, and the location of the cluster center as well as the sensitivity to the value of $H{_0}$ used for deriving the cluster hot gas temperature from the X-ray luminosity. The present analysis yields a total cluster mass in the range 1.5 - 3.2 $\times$ 10$^{14}$M${\odot}$, and does not include allowances for errors in several of the above characteristics, but does consider the effects of cluster core radii between 5$''$ and 45$''$, and cluster extents between 120$''$ and 280$''$. The calculations of the X-ray luminosity were performed for Hubble constants of 50 and 75 km s$^{-1}$ Mpc$^{-1}$. The average convergence parameter $\overline{\kappa} = ({\kappa}_{A} + {\kappa_B})/2$ of the cluster at the image locations A and B was found to lie between 0.07 and 0.21. \\ One of the other assumptions in the present calculation is that the temperature of the cluster follows the luminosity temperature relation of David et al. (1993) for a large sample of clusters of galaxies. The derived cluster temperature depends weakly on the cluster X-ray luminosity. An error of a factor of 2 in the X-ray luminosity derived from the HRI data will propagate as an error of about 30$\%$ in the cluster temperature and about 7$\%$ (for $\kappa$ = 0.2) in the derived value of the Hubble constant. The residual rms scatter in the $L_{X}$ - T relation along the temperature axis is $\sigma_{logT}$ = 0.104/2.1 according to Markevitch (1998) and propagates as an error of about 3$\%$ in the derived value of the Hubble constant. Our calculations for the deflection angles produced from considering the influence of the cluster alone (not including G1) found the deflection for images A and B to lie in the range of 2$''$ to 16$''$ as shown in Figure 6; the differences in these deflections for images A and B are much smaller. The resulting percent contribution, ${\Delta}S_{AB}$, to the observed 6.$''$1 separation of images A and B due to the cluster alone, lies in the range of 4\% to 22\%. This range includes only the effects of varying the cluster core radius between 5$''$ and 45$''$ and the cluster limit between 120$''$ and 280$''$. This result is in contrast with the commonly accepted notion that the large separation of the two images of Q0957+561 is an indicator that the lens contains a cluster of galaxies. To test the sensitivity of this result to the adopted cluster center position, we computed the cluster contribution to the image separation for cluster center locations within a distance of 5$''$ of the location derived from the optical observations by Angonin-Willaime et al. (1994). As expected, when the cluster center is moved closer to the locations of images A and B, the derived cluster contribution to the image separation becomes more significant reaching a value as high as ${\Delta}S_{AB}$ = 31\% for a cluster center at a distance of about 16$''$ from image B. \\ Recently the mass distribution of the lens of Q0957+561 was determined by Fischer et al. (1997) based on the distortion of background galaxies produced by the lens. Their estimate of 3.7 $\pm$ 1.2 $\times$ 10$^{14}$M${\odot}$ within a radius of 1Mpc is consistent with the X-ray estimates. As pointed out by Schneider \& Seitz (1995) however, the Kaiser \& Squires (1993) mass reconstruction technique also is insensitive to constant density sheets of matter. An interesting conclusion from the mass reconstruction analysis by Fischer et al. (1997) is the relatively low value of 5$''$ derived for the cluster core radius. The cluster core radius, r$_{cF}$, as defined in Fischer et al. (1997), however, differs from the cluster core radius, r$_{c}$, used in this paper and defined in equation 2. The relation between r$_{c}$ and r$_{cF}$ is given approximately by r$_{cF}$ = -0.987 + 0.685r$_{c}$ for 5$''$ $<$ r$_{c}$ $<$ 50$''$. We will have to await future X-ray observations of Q0957+561 to compare more accurately X-ray and weak-lens results for the cluster core radius. \\ The present ROSAT detection of cluster emission suggests that future X-ray observations will greatly add to our understanding of Q0957+561. AXAF observations should provide direct determination of the temperature profile, the core radius and the shape of the cluster lens, reveal possible clumpiness in the total mass profile, and ultimately provide a tighter constraint on the Hubble constant.\\ | 98 | 3 | astro-ph9803285_arXiv.txt |
9803 | astro-ph9803146_arXiv.txt | Many disk galaxies are lopsided: their brightest inner parts are displaced from the center of the outer isophotes, or the outer contours of the H{\sc i} disk. This asymmetry is particularly common in small, low-luminosity galaxies. We argue here that long-lived lopsidedness is a consequence of the disk lying off-center in the potential of the galaxy's extended dark halo, and spinning in a sense retrograde to its orbit about the halo center. The stellar velocity field predicted by our gravitational $N$-body simulations is clearly asymmetric. | Many galaxies display lopsided, rather than bisymmetric, structure, on both large and small scales. Hubble Space Telescope observations show that the nuclei of M31 and NGC 4486B do not lie at the center of the bulge isophotes (Lauer \etal\ 1993, 1996; Davidge \etal\ 1997). In our own Milky Way, the distribution of molecular gas as measured in CO is offset about $1^{\circ}$ in longitude and +15\kms\ in velocity from the center (Binney \etal\ 1991); other mass concentrations in that region also appear to be off-center (Blitz 1995). On kiloparsec scales, many galaxies are asymmetric in both their optical appearance and their gas distribution (\eg\ Baldwin, Lynden-Bell \& Sancisi 1980). Working from a sample of 1700 spectra, Richter \& Sancisi (1994) estimate that about half of all late-type spiral galaxies show clearly asymmetric H{\sc i} profiles, indicating a lopsided gas disk; about half of this sample are nearby field galaxies, and none lie in rich clusters where interactions between galaxies are frequent. Zaritsky \& Rix (1997) estimate that about 50\% of spiral galaxies in the field are significantly lopsided at optical and near-infrared wavelengths; the asymmetry involves both young and old stellar populations, and most of their lopsided galaxies lack any obvious companions. Asymmetric bars are particularly common in late type or small disk galaxies, less luminous than $M_B \sim 18$ (\eg\ Feitzinger 1980; Odewahn 1996; Matthews \& Gallagher 1997). While some lopsidedness in the outer parts of galaxies is undoubtedly due to recent accretion, the fact that lopsided asymmetries are so common suggests that they can be long-lived. The origin and persistence of these disk asymmetries remain a mystery. Baldwin, \etal\ (1980) suggest that the lopsided distortions seen in the outer parts of galaxies are kinematic features; in that case they should wind themselves into a leading spiral, since the slow pattern speed $\Omega - \kappa$ is negative. In fact the one-armed spiral often has the same sense of winding as the two-armed components, which presumably trail (Colin \& Athanassoula 1989; Phookun \etal\ 1992, 1993). The only analytic disk models which are strongly unstable to lopsided distortions are those with many retrograde-streaming particles (Zang \& Hohl 1978; Sawamura 1988; Sellwood \& Merritt 1994); but observed counter-rotation is rare in disk galaxies (\eg\ Kuijken, Fisher \& Merrifield 1996). Adams, Ruden \& Shu (1989) suggested that a protostellar disk could become unstable to a lopsided distortion; in the galactic context, this would correspond to the stellar disk being off-center in the dark halo. However, Heemskerk, Papaloizou \& Savonije (1992) showed that the $m=1$ distortion is generally evanescent in the disk, so that resonant growth will not occur. Matthias (1993), investigating orbits in a weakly `egg-shaped' tumbling potential, found that prograde orbits were distorted so as to oppose the `eggness'. Miller \& Smith (1992) describe an $N$-body simulation of an elliptical galaxy in which a central tilted disk developed of an off-center motion, tracing an oscillation in the underlying particle distribution, and Taga \& Iye (1998) have done simulations in which a massive object wanders off-center in a dense stellar system; but no similar results have been reported in the literature. By contrast, kinematic models for lopsided systems have been quite successful; orbits in the potential of an off-center bar rotating rigidly in an axisymmetric disk can account for the rotation curves of Magellanic barred systems (de Vaucouleurs \& Freeman 1973; Christiansen \& Jeffreys 1976), while the gas response has a trailing one-armed form (Colin \& Athanassoula 1989). We propose that the key to the puzzling ubiquity of lopsided galactic disks is that they are not dynamically isolated inside the galaxy. Late-type spirals and dwarf galaxies, in which larger-scale lopsidedness is most frequently observed, are most likely to be dominated by dark halo mass: \eg\ Casertano \& van Gorkom (1991); C\^ot\'e, Carignan \& Sancisi (1991); Broeils (1992a, b). (Athanassoula, Bosma \& Papaoiannou 1987 suggested that in later-type galaxies the halo core is larger in relation to the radius of the stellar disk, although this was not confirmed by Broeils 1992a, b.) If a disk found itself off-center in a dominant dark halo with a core of nearly constant density that was large compared to the disk dimensions, most of the disk mass would lie in a region where the angular speed of an orbit about the halo center did not vary strongly. Self-gravity might then act to maintain the disk's coherence as it orbits the center of the halo potential well. This paper presents our investigation of just such a model using a self-consistent, self-gravitating disk embedded in a static dark halo. | Our experiments imply that a gravitating galactic disk can remain off-center in a halo potential as long as it orbits in a region where the halo has nearly constant density; the effect is more pronounced if the spin of the disk is in the opposite sense to its orbit around the halo center. This behavior can be understood in the limit where the disk dimensions are small, and its self-gravity is weak: a particle in near-circular orbit in the halo then follows an epicycle in a sense retrograde to the orbit. A collection of particles with the same angular momentum, and hence with a common guiding-center radius, would stay together as they orbited the halo center. In the halo core, where the orbital frequency does not vary strongly with radius, the disk's self-gravity can fairly easily maintain coherence in approximately this configuration. The relatively large fraction of lopsided systems found at redshifts $z \sim 1$, which was apparently a period of rapid galaxy formation (\eg\ Driver, Windhorst \& Griffiths 1995), suggests that many disks can slowly lose their off-center character as they settle down. In dwarf galaxies, the halo is relatively more massive (Broeils 1992a, b), so lopsidedness is more easily maintained than in more luminous systems. The relation between a lopsided disk and the presence of a central bar has not been studied systematically, but both these features are characteristic of Magellanic irregular systems. The lopsided Sc galaxy NGC~1637 (Block \etal\ 1994) appears prominently barred in an infrared image, while at optical wavelengths the bar is hidden by dust and confused by star formation. The spatial scale of the predicted asymmetry is approximately that of the constant-density core; consistent with this idea, the off-center nuclei of M31 and the elliptical NGC~4486B lie within resolved central cores in those galaxies. Our results are consistent with those of Bontekoe (1988), who investigated the sinking of a satellite galaxy into a larger system. In his unpublished fifth chapter, he reports that the satellite never sank all the way to the center, but the orbit shrank until it enclosed about 10\% of the main galaxy's mass. The larger system was represented by a series of polytropes; the more concentrated the polytrope, the closer the satellite sank to its center. Our study also has aspects in common with work on galaxy mergers by Miller \& Smith (1995), who found that if two polytropic `cores' orbited each other in isolation, dynamical friction caused them to merge within a couple of orbits. But in the presence of a `halo' of constant density, whether it was represented by a fixed potential, or consisted of `live' simulation particles, the cores survived for $15-30$ orbits without coming noticeably closer, although the interpretation of the `live halo' model was complicated by an overstable oscillation in the cores' orbital radii. Persistent off-center disks have not been reported in published $N$-body simulations following a galactic disk in an imposed external bulge potential (see \eg\ Sellwood \& Wilkinson 1993); but in these computations, the mass distribution of the bulge was in general more concentrated than the disk. We must be concerned about the effect of dynamical friction on the disk, from the particles which would make up a `live' galactic halo. In the context of M31's nucleus, King, Stanford \& Crane (1995) point out that this drag could well be small, if the bulge stars stream rapidly in the same direction as the orbiting cluster's motion. It is unclear even how far the usual estimates of dynamical friction are to be trusted. Bontekoe's (1988) satellite stopped sinking at some distance from the center of the system into which it was accreting. The physical situation of an off-center disk has much in common with that of a rotating bar, where estimates of the drag from the particles of a spherical halo suggest that the bar should rapidly spin down (Weinberg 1985), and gravitational $N$-body experiments also indicate strong braking. But observed galactic bars seem to be fast-rotating: Sellwood (1996) summarizes this confusing situation. The long-lived pairs of {\it counter-rotating} bars which develop in the gravitational $N$-body simulations of Sellwood \& Merritt (1994) and Friedli (1996) are clearly not discouraged by dynamical friction. It may well be more useful, instead of calculating only the back-reaction of the halo on the orbiting disk, to look for long-lived modes of oscillation in the combined system. Weinberg (1991, 1994) has developed an analytic method for doing this; examining spherical King models with an isotropic velocity dispersion, he found $m=1$ modes which, once excited, require many orbital times to decay. The distortions lasted longest in the models in which the core was largest in relation to the total extent of the system. He then used gravitational $N$-body simulations to confirm the existence of these modes. Similar lopsided modes may well exist for a disk within a galactic halo. We would then expect a relatively long-lived lopsided structure to result when material is accreted such as to push it off-center in the underlying potential. In the central regions of a galaxy, where the gravitational force is provided largely by a `hot' stellar system such as a bulge, an inner disk or nuclear star cluster could remain off-center over many orbital periods. Gas might be captured by a galactic disk such as to push it off-center in the halo; alternatively, accretion of dark matter onto the halo may result in the disk lying off-center. If the orbit of the disk is retrograde with respect to its spin, the lopsidedness is likely to be stronger and more persistent. To maintain appreciable asymmetry, the near-constant-density halo core should be large enough to encompass a substantial fraction of the disk mass. | 98 | 3 | astro-ph9803146_arXiv.txt |
9803 | astro-ph9803236_arXiv.txt | By redshift of $10$, star formation in the first objects should have produced considerable amounts of Carbon, Nitrogen and Oxygen. The submillimeter lines of C, N and O redshift into the millimeter and centimeter bands ($0.5\,{\rm mm}$--$1.2\,{\rm cm}$), where they may be detectable. High spectral resolution observations could potentially detect inhomogeneities in C, N and O emission, and see the first objects forming at high redshift. We calculate expected intensity fluctuations and discuss frequency and angular resolution required to detect them. For CII emission, we estimate the intensity using two independent methods: the line emission coefficient argument and the luminosity density argument. We find they are in good agreement. At $1+z \sim 10$, the typical protogalaxy has a velocity dispersion of $30 \,{\rm km\,s^{-1}}$ and angular size of $1 \,{\rm arcsecond}$. If CII is the dominant coolant, then we estimate a characteristic line strength of $\sim 0.1 \,{\rm K\, km\, s}^{-1}$. We also discuss other atomic lines and estimate their signal. Observations with angular resolution of $10^{-3}$ can detect moderately nonlinear fluctuations of amplitude $2 \cdot 10^{-5}$ times the microwave background. If the intensity fluctuations are detected, they will probe matter density inhomogeneity, chemical evolution and ionization history at high redshifts. | When did the first stars and galaxies form? How can we detect them? Quasars were once the most distant objects that have ever been observed (\cite{sch91}). Now galaxies are beginning to take their place. Discovery of a large number of galaxies at $z>3$ (\cite{ste96}), and even up to $z\sim 5$ (\cite{fra97,dey98}), suggests that galaxies existed well before quasars did. In hierarchical models, star formation starts relatively early, by $z\sim30$ (\cite{cou86,fuk94,gne96,teg97}), so even these $z \sim 5$ galaxies are not representative of the first generation of objects. The purpose of this paper is to investigate the possibility of detecting primordial objects through the atomic lines of carbon, nitrogen and oxygen, which should have been produced as a result of star formation. At low redshift, these submillimeter lines are important coolants (e.~g., \cite{ben94,isr95,stu97,you97,barv97,mal97}). If the source is at redshift $1+z \sim 10$--$20$, the lines are redshifted to radio frequencies, where the atmosphere is more transparent and ground--based detectors with very high sensitivity and resolution exist. If the source was uniformly distributed through space, then its emission would merely distort the spectrum. Since heavy elements produced by stars are distributed inhomogeneously, the observed emission should have intensity fluctuations as a function of both position and frequency. With high enough resolution one can in principle distinguish them from background. If we take $1 h^{-1}\,{\rm Mpc}$ as a characteristic comoving scale of the inhomogeneity, this corresponds to $\sim 1$ arcmin in angular scale, and $\sim 10^{-3}$ in frequency resolution. On smaller scales, velocity dispersion of matter limits the scale that can be resolved in frequency space. In this paper we formulate the intensity fluctuations and estimate their magnitude. This signal should give us information on number density fluctuations of specific source elements at $1+z>10$. Detailed determination of the peak of this signal will also tell us at which epoch elements were ionized. The rest of this paper is organized as follows. In the next section, we present our formalism for estimating the fluctuating signal. Then we estimate its magnitude in the case that the matter density fluctuations are calculated from the linear power spectrum in section \ref{sect:linear}. We also estimate optimal frequency and angular resolution for an object search. We focus on atomic line redshifted into present radio band, investigating carbon, nitrogen and oxygen emission lines at rest frame wavelength of hundreds of micrometers. In section \ref{sect:nonlin}, we discuss the size and the number density of nonlinear clumps at the reionization epoch, and estimate intensity from nonlinear clumps and the interval between enhance of intensity fluctuations. Section \ref{sect:paremis} gives discussion of merits and demerits as the signal of specific emission. After alternative estimation of CII luminosity in section \ref{sect:altcii}, the final section gives our conclusion. | We have proposed a new and promising method to detect the first objects: search for carbon, nitrogen and oxygen line emission at $1+z=10$--$20$. We conclude that the intensity fluctuations originated from the emission are detectable in radio band ($0.5\,{\rm mm}$--$1.2\,{\rm cm}$). If observation is carried out with angular resolution $\theta_{\rm res} = 1 \,{\rm arcmin}$ and spectral resolution $\Delta \nu / \nu = 10^{-3}$, the magnitude of the smoothed intensity fluctuations corresponding with linear matter density fluctuations is predicted as $\Delta \bar{I}_{\rm line} / I_{\rm tot}= (1$ and $3) \times 10^{-6}$ (in CI $609$ and $370 \,\mu {\rm m}$ cases) from sources at $1+z=10$, in the standard CDM model. The predictions for other lines are $\Delta \bar{I}_{\rm line} / I_{\rm tot} = (20, 70, 3, 20) \times 10^{-6}$, respectively, in [$158$(CII), $146$(OI), $204$(NII), $122$(NII)] $\,\mu {\rm m}$ cases. Observation with $\theta_{\rm res} = 1 \,{\rm arcsec}$ and $\Delta \nu / \nu = 10^{-4}$ resolve individual protogalaxies with velocity dispersion of $30\,{\rm km\,s^{-1}}$. We predict line intensities relative to the microwave background of $\Delta I_{\rm line} / I_{\rm tot} = (1, 4, 30, 80, 4, 30) \times 10^{-4}$ or equivalently $\Delta T_{\rm line} v_{\rm typ} = (0.8, 2, 7, 20, 1, 6) \times 10^{-2} \,{\rm K\, km\, s}^{-1}$ for [$609$(CI), $370$(CI), $158$(CII), $146$(OI), $204$(NII), $122$(NII)] $\,\mu {\rm m}$ lines, respectively. If detected, these features will provide useful information on matter inhomogeneities at $1+z=10$--$20$. They will also probe chemical evolution and reionization time $z_{\rm re}$ itself. | 98 | 3 | astro-ph9803236_arXiv.txt |
9803 | astro-ph9803222_arXiv.txt | We define a complete sample of thirty-three GHz-Peaked-Spectrum (GPS) radio sources based on their spectral properties. We present measurements of the radio spectra and polarization of the complete sample and a list of additional GPS sources which fail one or more criteria to be included in the complete sample. The majority of the data have been obtained from quasi-simultaneous multi-frequency observations at the Very Large Array (VLA) during 3 observing sessions. Low frequency data from the Westerbork Synthesis Radio Telescope (WSRT) and from the literature have been combined with the VLA data in order to better define the spectral shape. The objects presented here show a rather wide range of spectral indices at high and low frequencies, including a few cases where the spectral index below the turnover is close to the theoretical value of 2.5 typical of self-absorbed incoherent synchrotron emission. Faint and diffuse extended emission is found in about 10\% of the sources. In the majority of the GPS sources, the fractional polarization is found to be very low, consistent with the residual instrumental polarization of 0.3 $\%$. \footnote{Tables 4 and 5 are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/Abstract.html} | The GHz-peaked-spectrum (GPS) radio sources are characterized by a simple convex spectrum which peaks in a range of about a decade around 1 GHz. General discussions on the general properties of these objects are given by O'Dea \etal\ (1991) O'Dea and Baum (1997), and O'Dea (1998) where an exhaustive bibliography can also be found. Common characteristics of the bright sample of GPS radio sources are: small size ($\lae 1$ kpc), high radio luminosity, low fractional polarization, and apparently low variability. They are a mixed class of quasars and galaxies. Galaxies tend to be $L_*$ or brighter and at redshifts $0.1 \lae z \lae 1$ (O'Dea \etal\ 1996) while quasars are often found at very large redshift $1 \lae z \lae 4$, (O'Dea 1990). Currently, there are two main competing hypotheses to explain the origin of GPS radio sources and their possible evolution. In the ``young source" scenario, first suggested by Phillips and Mutel (1982), GPS radio sources with compact double (CD) morphology or compact symmetric morphology (Compact Symmetric Objects: CSO) are classical double radio sources at the very first stage of their lives. In the ``frustration'' scenario, GPS radio sources will never become as large as the classical doubles since they are confined to the sub-kpc scale by a dense and turbulent ambient medium (e.g., O'Dea \etal\ 1991, Carvalho 1994, 1998). However, it is unclear whether the gas density and its distribution in the nuclear region is sufficient to confine the radio source for times of the order of 10$^7$ years (see e.g., O'Dea 1998). Recent results support the young source hypothesis. Fanti \etal\ (1995) presented a model for the evolution of the galactic-scale Compact Steep Spectrum (CSS) sources into large scale classical doubles. They argued that the objects showing symmetric morphology are probably not confined by a dense and clumpy medium (see also De Young 1993). They also suggest that the typical CSS source age is of the order of 10$^6$ years and that the radio source luminosity decreases by an order of magnitude as the size of the radio source grows from a few kpc to hundreds of kpc (see also Begelman 1996). The same conclusions are reached by Readhead et al. (1996a,b) in their study of a few CSO's. O'Dea and Baum (1997), combining the complete sample of GPS presented here with the CSS sources by Fanti \etal\ (1990) reach similar conclusions. They further note note that the luminosity evolution required makes it likely that GPS and CSS radio sources will evolve into objects less powerful than the most powerful classical doubles. The detection of arc-second scale faint extended emission around $\sim 10\%$ of GPS sources (Baum \etal\ 1990, Stanghellini \etal\ 1990) motivated Baum \etal\ (1990) to suggest that nuclear activity is recurrent in these sources. In this hypothesis, we see the relic of a previous epoch of activity as faint diffuse emission surrounding the current young nuclear source. In order to achieve a deeper understanding of the GPS radio sources we pursued the following project. (1) We created a complete sample of GPS radio sources by means of a preliminary bibliographic research, checked with subsequent new multi-frequency data from the Very Large Array (VLA) and the Westerbork Synthesis Radio Telescope (WSRT). (2) We determined the properties of the radio spectrum and the polarization. (3) We obtained optical imaging to determine the host galaxy properties. (4) We obtained VLBI observations to study the milliarcsecond radio morphology. In this paper we present the selected sample and the results from the VLA and WSRT radio observations. The optical properties of GPS radio sources are discussed by O'Dea, Baum and Morris (1990), Stanghellini \etal\ (1993), and O'Dea \etal\ (1996). The milliarcsecond morphology is presented in Stanghellini \etal\ (1997). Constraints on radio source evolution based on our observations of the complete GPS sample are discussed by O'Dea \& Baum (1997). H$_0$=100 km sec$^{-1}$ Mpc$^{-1}$, and q$_0$=0.5 have been used in this paper. | In this paper we present the observational results and the primary qualitative conclusions. In a future paper we will present a quantitative discussion of the results. \subsection{The radio spectra} The flux densities are presented in Tables 4 and 5 and plotted in Figures 1 to 5. In order to extract information from the radio spectra, we fitted them with a hyperbola (which is a curve tending asymptotically to a straight line at the extrema). We first fit the spectral indices in the thick and thin part of the spectrum independently, and then fixed these two values (which correspond to the angular coefficients of the asymptotes) and we performed a least square fit, solving for the other parameters of the hyperbola. We calculated the frequency of the spectral peak (Tables 1 and 2) using the fitted curves. We show the distributions of the observed and rest frame turnover frequency for the 33 sources of the complete sample in Figures 6c and 6d. In Fig. 6a and 6b we show the high frequency (above the peak) and low frequency (below the peak) spectral index distribution, respectively. \subsection{The spectral indices and the turnover frequency} The GPS radio sources are a mixed group of galaxies and quasars, with some remarkable differences between the two classes. The histogram in Fig. 16a shows that the redshift distribution is very different for galaxies and quasars. The galaxies have a typical redshift of $\sim$ 0.5, and none has a redshift higher than 1. For the galaxies without redshift information, we note that only 0316+161 has an optical magnitude slightly fainter than the galaxy 2128+048 at redshift 0.99 (Table. 1). Since these galaxies follow the Hubble diagram (O'Dea \etal\ 1996; Snellen \etal\ 1996) it is unlikely they will be found at a redshift much higher than the others in the sample. The quasars are instead found at any redshift (we included the galaxy 1404+286 (OQ208), which has a Seyfert 1 nucleus, in the quasar class) with the majority at very high z (see also O'Dea 1990). \begin{figure} \vbox{\psfig{figure=figure7.ps,width=8cm,height=9cm}}\par \caption[]{$0248+430$ at 1.35 GHz. The restoring beam is 1.41 $\times$ 1.32 arcsec in P.A. $-78^\circ$. The r.m.s. noise on the image is 0.2 mJy. The peak flux is 810 mJy/beam. The contour levels for all the images are -3, 3, 6, 12, 25, 50, 100, 200, 500, 1000 $\times$ the r.m.s noise} \end{figure} \begin{figure} \vbox{\psfig{figure=figure8.ps,width=8cm,height=9cm}}\par \caption[]{$0528+134$ at 1.66 GHz. The restoring beam is 1.14 $\times$ 1.01 arcsec in P.A. $+2^\circ$. The r.m.s. noise on the image is 0.25 mJy. The peak flux is 2088 mJy/beam. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure9.ps,width=8cm,height=9cm}}\par \caption[]{$0738+313$ at 1.33 GHz. The restoring beam is 2 $\times$ 2 arcsec. The r.m.s. noise on the image is 0.3 mJy. The peak flux is 1943 mJy. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure10.ps,width=8cm,height=6cm}}\par \vskip -1cm \caption[]{$0941-080$ at 1.33 GHz. The restoring beam is 2.44 $\times$ 1.54 arcsec in P.A. $+29^\circ$. The r.m.s. noise on the image is 0.5 mJy. The peak flux is 2088 mJy/beam. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure11.ps,width=8cm,height=9cm}}\par \caption[]{$2134+004$ at 1.33 GHz. The restoring beam is 1.91 $\times$ 1.55 arcsec in P.A. $-5^\circ$. The r.m.s. noise on the image is 0.4 mJy. The peak flux is 3249 mJy/beam. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure12.ps,width=8cm,height=9cm}}\par \caption[]{$2223+210$ at 1.33 GHz. The restoring beam is 1.37 $\times$ 1.32 arcsec in P.A. $-7^\circ$. The r.m.s. noise on the image is 0.3 mJy. The peak flux is 1663 mJy/beam. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure13.ps,width=8cm,height=9cm}}\par \caption[]{$2223+210$ at 5 GHz. The restoring beam is 0.37 $\times$ 0.35 arcsec in P.A. $+5^\circ$. The r.m.s. noise on the image is 0.3 mJy. The peak flux is 1034 mJy/beam. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure14.ps,width=8cm,height=9cm}}\par \caption[]{$2230+114$ at 1.33 GHz. The restoring beam is 1.46 $\times$ 1.36 arcsec in P.A. $+1^\circ$. The r.m.s. noise on the image is 0.8 mJy. The peak flux is 6444 mJy/beam. } \end{figure} \begin{figure} \vbox{\psfig{figure=figure15.ps,width=8cm,height=9cm}}\par \caption[]{$2230+114$ at 5 GHz. The restoring beam is 0.38 $\times$ 0.36 arcsec in P.A. $-11^\circ$. The r.m.s. noise on the image is 0.5 mJy. The peak flux is 4137 mJy/beam. } \end{figure} \begin{figure*} \vbox{\psfig{figure=figure16.ps,width=16.5cm,height=22cm}}\par \vskip -6cm \caption[]{Histograms for the complete sample: a) redshift distribution; b) fractional polarization at 1.3 GHz; c) fractional polarization at 4.9 GHz; d) fractional polarization at 8.5 GHz } \end{figure*} The high frequency spectral index ranges from 0.5 (the limit set in the selection criteria) to 1.3 with the galaxies having perhaps slightly steeper values (but the 2 objects with the steepest spectral indices are quasars). The low frequency spectral index ranges from -0.2 to -2.1 without any clear difference between galaxies and quasars, but the higher values are biased since in several objects the low frequency part of the spectrum is under-sampled and the spectral index is calculated close to the turnover frequency where it is likely to be flatter. Similar results were found by De Vries \etal\ (1997) though their poorer frequency coverage resulted in a somewhat smaller range in spectral index. The quasars tend to peak at higher frequencies than the galaxies in both the rest frame and observed frame and some quasars have a turnover frequency in their rest frame exceeding 10 GHz. This suggests that, on the assumption that the turnover is caused by synchrotron self-absorption, the GPS quasars are more compact than the galaxies. This effect has been also found by De Vries \etal\ (1997) in a bright heterogeneous sample and by Snellen (1997) in a fainter sample. In addition, VLBI images of several sources belonging to the complete sample show that quasars are more compact than galaxies, and in general exhibit different morphologies (Stanghellini \etal\ 1997). These results suggest that either GPS galaxies and GPS quasars are different types of objects, or that beaming of compact components plays a role in the quasars (see also O'Dea 1998). \begin{table*} \caption{Polarization for the complete sample } \begin{flushleft} \begin{tabular}{rrrrrrrrrrrrrrr} \hline \noalign{\smallskip} \hline \noalign{\smallskip} & 1380 && 1640 && 4835 && 4885&& && &&&\\ name & $\%$ & PA & $\%$ & PA& $\%$ & PA & $\%$ & PA & & & &&& \\ \noalign{\smallskip} \hline\noalign{\smallskip} 0019$-$000& 0.5& +0 & 0.3 &-2 & $<$0.1& & $<$0.1& & & & &&&\\ 0237$-$233& 2.0& +7 & 1.2 &+4 & 4.0& -35 & 4.0 &-35 & & & &&&\\ 0500+019& $<$0.2& & $<$0.2 & & $<$0.1& & $<$0.1& & & & &&&\\ \noalign{\smallskip} \hline \noalign{\smallskip} & 1380&& 1630&& 4815&& 4865 && 8435&& 8485&&&\\ name & $\%$ & PA & $\%$ & PA& $\%$ & PA & $\%$ & PA & $\%$ & PA & $\%$ & PA&&\\ \noalign{\smallskip} \hline \noalign{\smallskip} 0108+388& $<$0.2& & 0.5 & & $<$0.2&& $<$0.1&& $<$0.1& & $<$0.1& & &\\ 0316+161& 0.5& -63 & 0.7 & -60 & $<$0.1&& $<$0.1&& $<$0.1& & $<$0.1& &&\\ 0428+205& 0.4& -59 & 0.6 & -58 & $<$0.1&& $<$0.1&& $<$0.1& & $<$0.1& &&\\ 0710+439& $<$0.1& & $<$0.3 & & $<$0.1&& $<$0.1&& $<$0.1& & $<$0.1& &&\\ 1031+567& $<$0.2& & $<$0.2 & & $<$0.3&& $<$0.2&& $<$0.2& & $<$0.2& && \\ 1323+321& 0.3& -35 & 0.4 & -01 & $<$0.1&& $<$0.1&& 0.9& +14 & 0.8& +15&&\\ 1404+286& 0.4& -50 & $<$0.3 & -23 & $<$0.2&& $<$0.1&& $<$0.2& & $<$0.2&&&\\ 1518+047& $<$0.2& & $<$0.2 & & $<$0.3&& $<$0.1&& $<$0.2& & $<$0.2& & & \\ 1607+268& $<$0.3& & $<$0.2 & & $<$0.1&& $<$0.1&& $<$0.1& & $<$0.1& &&\\ 2008$-$068& 1.7& -19 & 2.5 & -21 & $<$0.3&& $<$0.3&& $<$0.3& & $<$0.3& & \\ 2352+495& $<$0.2& & $<$0.3 & -34 & $<$0.1&& $<$0.1&& $<$0.1& & $<$0.1& &&\\ \noalign{\smallskip} \hline \noalign{\smallskip} & 1335&& 1665&& 4535 && 4985&& 8085 && 8465&&RM(rad/m$^2$)&\\ name & $\%$ & PA & $\%$ & PA & $\%$ & PA & $\%$ & PA & $\%$ & PA& $\%$ & PA&obs&rest\\ \noalign{\smallskip} \hline \noalign{\smallskip} 0248+430& 1.7& +52? & 1.5 & +23 & 2.0& -18 & 2.0 & -27 & 0.5& -40 & 0.4& -44&131&703\\ 0457+024& $<$0.1& & $<$0.1 & &$<$0.3& & 0.3 & +47 & 1.0& +14 & 1.2& +12&258&2954\\ 0738+313& $<$0.1& & $<$0.3 & & 1.7& +41 & 1.6 & +67 & 3.0& -2 & 3.2& +2&-813&-2160\\ 0742+103& $<$0.2& & $<$0.2 & &$<$0.1& & $<$0.2 & & $<$0.1& & $<$0.1& & & \\ 0743$-$006& 0.4& -76 & 0.5 & -7 &$<$0.2& & $<$0.2 & & 0.8& +1 & 1.0& -18& &\\ 0941$-$080& $<$0.2& & $<$0.1 & &$<$0.1& & $<$0.1 & & $<$0.2& & $<$0.2& & &\\ 1117+146& $<$0.2& & $<$0.1 & &$<$0.2& & $<$0.3 & & $<$0.2& & $<$0.2& && \\ 1127$-$145& 4.3& -84 & 3.9 & +61 & 2.6& -26 & 2.8 & -24 & 3.3& -17 & 3.3& -18 &-49&-234\\ 1143$-$245& 1.9& +72 & 2.1 & +80 & 0.7& -18 & 0.7 & -27 & 1.3& -43 & 1.4 & -46&146&1270\\ 1245$-$197& $<$0.2& & $<$0.1 & &$<$0.1& & $<$0.1 & & $<$0.1& & $<$0.1 & && \\ 1345+125& $<$0.2& & $<$0.1 & &$<$0.1& & $<$0.2 & & $<$0.1& & $<$0.1 & & & \\ 1358+624& $<$0.2& & $<$0.1 & &$<$0.1& & $<$0.3 & & $<$0.2& & $<$0.2 & &&\\ 1442+101& 0.8& -74 & 0.8 & -78 & 2.0& +60 & 1.7 & +64 & 1.4& +82 & 1.6 & +78&-116&-2395\\ 1600+335& $<$0.2& & $<$0.1 & &$<$0.1& & $<$0.1 & & $<$0.1& & $<$0.1 & & & \\ 2126$-$158& $<$0.2& & $<$0.2 & &$<$0.1& & $<$0.2 & & $<$0.1& & $<$0.2 & & &\\ 2128+048& $<$0.1& & $<$0.1 & &$<$0.1& & $<$0.1 & & $<$0.2& & $<$0.2 & & &\\ 2134+004& $<$0.1& & $<$0.1 & & 0.9& +60 & 0.9 & +43 & 0.5& -4 & 0.5 & +0&349&3008\\ 2210+016& $<$0.1& & $<$0.1 & &$<$0.1& & $<$0.1 & & $<$0.2& & $<$0.2 & && \\ 2342+821& $<$0.2& & $<$0.1 & &$<$0.1& & 0.5 & -63 & &&&&&\\ \noalign{\smallskip} \hline \noalign{\smallskip} \end{tabular} \end{flushleft} \end{table*} \begin{table*} \caption{Polarization of additional objects } \begin{flushleft} \begin{tabular}{rrrrrrrrrrrrrrr} \hline \noalign{\smallskip} \hline \noalign{\smallskip} & 1335&& 1665&& 4535&& 4985&& 8085&& 8465&&RM(rad/m$^2$)&\\ name & $\%$ & PA & $\%$ & PA & $\%$ &PA & $\%$ &PA &$\%$ &PA &$\%$ &PA&obs&rest\\ \noalign{\smallskip} \hline \noalign{\smallskip} 0404+768& $<$0.2& &$<$0.1 & &$<$0.1& & $<$0.3& & & & & &&\\ 0440$-$003& 0.9& +49 & 0.9 & -27 & 2.3& +61& 2.1& +58& 2.6& +55& 2.6 & +56&28&\\ 0528+134& 0.6& -56 &$<$0.3 & & 2.2& -31& 2.6& -33& 3.9& -39 & 4.0 & -41&52&\\ 2149+056& $<$0.2& &$<$0.2 & &$<$0.1& & $<$0.1& & $<$0.2 & & $<$0.2 & && \\ 2223+210& 13.3& -28 &10.5 & -82 &10.0& -64& 9.1& -58& & & &-123&-1077\\ 2230+114& 2.5& -87 & 2.6 & -33 & 1.2& +61& 1.5& +75& & & &-59&-244\\ \noalign{\smallskip} \hline \noalign{\smallskip} & 1380& & 1630& & 4815&& 4865 && 8435&& 8485&&RM(rad/m$^2$)&\\ name & $\%$ & PA & $\%$ & PA & $\%$ &PA & $\%$ &PA &$\%$ &PA &$\%$ &PA&obs&rest\\ \noalign{\smallskip} \hline \noalign{\smallskip} 0026+346& 0.4& -45 & 0.6 & -43 &$<$0.2& & $<$0.2& & $<$0.1& & $<$0.1&&&\\ 0201+113& 0.7& -48 & 0.9 & -38 & 1.4& +33& 1.4& +33& 0.5& +4 & 0.5& +5&131&2724\\ 0552+398& 0.4& -52 & 1.0 & +78 & 0.8& +43& 0.6& +53& 0.8& +66& 0.8& +66&-1344&-15218\\ 0703+468& $<$0.1& &$<$0.2 & &$<$0.2& & $<$0.2& & $<$0.3& & $<$0.3&&&\\ 0711+356& 1.3& +5 & 1.0 & -33 & 1.8& +83& 1.8& +80& 1.9& +76& 2.0& +79&40&275\\ 0904+039& $<$0.2& &$<$0.2 & &$<$0.5& & $<$0.5& & $<$0.9& & $<$0.9&&&\\ 0914+114& $<$0.2& &$<$0.2 & &$<$0.8& & $<$0.8& & $<$2.0& & $<$2.0&&&\\ 1543+005& $<$0.1& &$<$0.2 & &$<$0.2& & $<$0.2& & $<$0.2& & $<$0.1&&&\\ 1732+094& $<$0.2& &$<$0.3 & & 1.5& +16& 1.4& -1 & $<$0.3& & $<$0.3&&&\\ 2015+657& 4.5& -53 & 3.6 & -71 & 3.2& +54& 3.1& +55& 4.3& +51& 4.3& +51&30&\\ 2021+614& 0.6& -53 & 1.0 & -51 &$<$0.2& & $<$0.2& & $<$0.1& & $<$0.1&&&\\ 2050+364& $<$0.2& & 0.4 & -77 &$<$0.2& & $<$0.1& & $<$0.1& & $<$0.1&&&\\ 2137+209& $<$0.2& &$<$0.1 & &$<$0.2& & $<$0.2& & 0.4& +68& $<$0.3&&&\\ 2337+264& $<$0.2& &$<$0.3 & &$<$0.3& & $<$0.2& & $<$0.2& & $<$0.2&&&\\ \noalign{\smallskip} \hline \noalign{\smallskip} \end{tabular} \end{flushleft} \end{table*} \subsection{Extended emission} We have detected extended emission (both diffuse and compact) close to the compact radio source in some cases at 21 cm. In the remaining sources, our upper limits on extended emission is typically 1 mJy/beam at 21 cm. 0248+430 (Figure 7) has a compact emitting region 15 arcsec east of the main component and a hint of weak emission 5 arcsec to the south. 0528+134 is resolved, showing an extension in the NW direction (Figure 8). Murphy \etal\ (1993) present an image of 0738+313 at 20 cm showing 2 emitting regions resembling 2 weak hot-spots and lobes on the opposite sides of the dominant component. In our image (Figure 9) these 2 weak components are almost completely resolved out and only a hint of emission has been detected 30 arcsec north and south of the compact region. 0941-080 shows a slightly resolved secondary component 20 arcsec east of the main one (Figure 10). 2134+004 has very weak and diffuse emission around the strong compact component (Figure 11). 2223+210 has a secondary component 4 arcsec away from the main one in the SW direction in our image at 1.35 GHz (Figure 12); the main component itself is resolved in a core-jet structure oriented NE with a possible counter jet in the image at 5 GHz (Figure 13). 2230+114 at 1.35 GHz (Figure 14) shows an elongated structure in the NW-SE direction with a hint of emission bending to SW, while in the 4.9 GHz image (Figure 15) the elongated structure turns out to be a core-jet structure with the possible presence of a counter jet. Stanghellini \etal\ (1990) report several cases of extended emission around GPS radio sources. Of the objects presented here showing extended emission, 0528+134 and 2223+210 are not true GPS objects (see also section 4.4 for a discussion of the case of 0528+134). In a couple of sources (0248+430, 0941-080 both belonging to the complete sample) it is difficult to say whether the secondary emission is related to the GPS radio source and further observations are probably needed. The extended emission found around 0738+313, 2134+004, and 2230+114 (the first 2 objects belong to the complete sample) is likely to be related to the GPS object. In the complete sample, 0108+388 is known to have extended emission, so there are 3 to 5 objects out of 33 with known extended emission so far. This percentage of 9 to 15$\%$ is slightly smaller but consistent with that previously claimed by Stanghellini \etal\ (1990). It is clear that the vast majority ($\sim 90\%$) of the GPS sources appear to be truly isolated and have no emission beyond the kpc scale at the current limits. \subsection{Variability} Waltman \etal\ (1991) presented monitoring observations at 2.7 and 8.1 GHz for several GPS sources covering the time range from 1979 to 1988. Some sources as 0237-233, 1245-197, 1345+125 were found to be very stable in flux density. Others were found to be variable: 0552+398 shows a variation of $\sim$ 30 $\%$ at 8.1 GHz. 2134+004 has a variability of about 15-20$\%$ at 8.1 GHz. 2352+495 has a variability below 10$\%$ at 8.1 GHz. Also 0742+103 is slightly variable. Wehrle \etal\ (1992) also report variability for some GPS objects in the time range 1985-1991 at 4.8, 8, and 14.5 GHz, from the University of Michigan Radio Astronomy Observatory monitoring program for several sources, some of which are GPS objects. 0552+398 shows an increase in flux density exceeding 50$\%$ at 8.4 and 14.5 GHz, 1127-145 shows a quasi periodical variability of approximately 1 Jy at all the 3 frequencies. The source 2230+114 also shows a rather remarkable flux density variability at all the 3 frequencies with an amplitude of 0.5-1 Jy and is a well known low frequency variable source (Bondi \etal\ 1996). The variability of 1404+286 has been discussed by Stanghellini \etal\ (1996). In conclusion, we find that some GPS sources (mainly quasars) show mild to high flux density variability at cm and mm wavelengths. However, without uniform monitoring of the complete sample it will not be possible to determine how common this variability is. We also note a couple sources where the spectral shape is variable and at some times the spectrum was peaked, and at other times it was not, 0528+134, the well known gamma-ray source (Mukherjee \etal\ 1996), and 0201+110. Both of these sources show a rather flat spectrum in the VLA observations from the second or the third session and they would be very easily discarded as GPS radio sources. But 0528+134 has been included in the class of GPS radio sources because of its GPS-like spectrum from the literature (O'Dea \etal\ 1991), and 0201+113 really shows a convex spectrum in the data in the first VLA session published in O'Dea \etal\ (1990). This behavior is not surprising in highly variable radio sources as we may well expect that the presence of new radio components will change the spectral shape. Thus, there are sources which show a peaked spectrum only part of the time. This aspect of the GPS phenomenon deserves more attention as it could be related to the remarkably different properties found between GPS galaxies and (some?) GPS quasars. \subsection{Polarization} Due to the low level of observed polarization, the errors are dominated by the r.m.s noise on the polarization images (typically 0.5 mJy) and by contamination from residual unpolarized emission (estimated about 0.2-0.3$\%$ of the total flux density). The error in the polarized flux may then be calculated as $\sigma _P=\sqrt {0.5^2 + (0.003\times S_{mJy})^2}mJy$. We considered any measurement of the fractional polarization below 0.3$\%$ to be an upper limit (Tables 6 and 7). The errors in the position angle have a systematic contribution due to the uncertainty in the determination of the position angle of the calibrator (3C286), assumed to be around 2-3 degrees. This is the dominant contribution in the position angle error for most of the objects with detected polarized flux density. In Fig. 16 we show the histograms of the fractional of polarization (mostly upper limits) for the complete sample at 1.3, 4.9 and 8.5 GHz. The fractional polarization is in general low at all the frequencies. Only a few quasars have a fractional polarization above 1 $\%$ at 4.9 or 8.5 GHz. When polarized emission has been detected we attempted a linear fit to the polarization angles versus the squared observed wavelength, as is expected from the Faraday effect on polarized radiation propagating through a magnetized and ionized medium. The low level or even the lack of detection of polarized flux limited us to only a few sources (all quasars). The fits are generally rather good and are given in Table 6 and 7, and in Figures 17 and 18. Due to the better frequency coverage in the present observations, our estimated rotation measures supersede those reported by O'Dea \etal\ (1990), though we cannot rule out that some of the difference is due to variability. We find Faraday rotation measures in the rest frame above 1000 rad/m$^2$ for 5 quasars of the complete sample. We also found a very high value ($>10^4$ rad/m$^2$) for 0552+398 which does not belong to the complete sample but has a GPS shape. Sometimes the frequencies which give a good fit include those close to the turnover, and in the case of 0552+398 are all below the turnover. This implies that the region emitting the polarized emission is different from that responsible for the optically thick emission or that the turnover is not caused by synchrotron self absorption. \subsection{Summary and conclusions} We have presented a bright flux-density-limited complete sample of 33 GPS radio sources selected on the basis of their peaked radio spectra. The sample selection was based on observations with the VLA, WSRT, and other instruments. Additional GPS sources not belonging to the complete sample have also been observed. We present our results on the polarization and radio spectrum. We found remarkable differences in the properties of quasars and galaxies, the latter having lower turnover frequencies, mostly undetectable polarization and lower redshifts. In the few objects where polarization has been detected at many frequencies, the Faraday rotation measure in the rest frame often exceeds 1000 rad/m$^2$. In about 10\% of the sources we detect weak diffuse extended emission. In the remaining $\sim 90\%$ any extended emission has a peak surface brightness is less than about 1 mJy/beam at 21 cm. In a following paper we will discuss the implications of the properties of the complete sample in the framework of the scenarios proposed to explain the existence of the GPS radio sources. | 98 | 3 | astro-ph9803222_arXiv.txt |
9803 | astro-ph9803152_arXiv.txt | s{I discuss the interplay between inflation and microwave background anisotropies, stressing in particular the accuracy with which inflation predictions need to be made, and the importance of inflation as an underlying paradigm for cosmological parameter estimation.} | Because of calculational simplicity, and because it provides a good fit to the observational data, an initial spectrum of adiabatic density perturbations is normally assumed responsible for all the observed structures in the Universe, such as galaxy clusters and microwave background anisotropies. The inflationary cosmology provides a natural explanation for such an initial spectrum, and indeed the causal generation of large-scale adiabatic perturbations requires a period of inflation.\cite{L95} In studying the microwave background anisotropies, we believe we have a tool through which all of the cosmological parameters, such as the Hubble constant $h$ and the density parameter $\Omega_0$, can be probed. This is because the evolution of perturbations in the matter and radiation fields depends on all the cosmological parameters. On the other hand, the microwave background anisotropies are entirely due to linear perturbation theory, and hence entirely dependent on the initial perturbations you have in the first place. The only reason that one can indeed hope to extract cosmological parameters is because one believes that the initial spectrum takes on a simple form, perhaps a power-law, which can be parametrized simply and then those initial parameters thrown into the melting pot along with the cosmological ones and fitted by the data. One of the delights of inflation is that the initial spectrum is indeed typically predicted to take on a simple form, and furthermore one which can readily be calculated to high precision for a given inflationary model. | Inflation as a paradigm is both eminently and imminently testable by upcoming microwave background observations. For example, the prediction of a peak structure is extremely generic and quite specific to the situation where perturbations begin their evolution on scales much larger than the Hubble radius, and details such as the peak spacing promises a very strong test~\cite{HW}. Something as simple as an observed spectrum without multiple peaks appears sufficient to rule out inflation (see e.g.~Barrow and Liddle~\cite{BL97}). If inflation passes these tests, then detailed fitting to the observations promises startlingly high quality information about the inflationary mechanism. However, the main purpose of my presentation is to provide a reminder of the important role inflation has in underpinning the microwave background endeavour. I stressed at the start that we can only get highly quality constraints from the present radiation power spectrum if we have a simple form, preferably motivated by theory, for the initial perturbations. Since the observations aim to be accurate at the percent level, the input information needs this accuracy too, and inflationary theory is now in a position where predictions at this level of accuracy can be made for all known models. This can be contrasted with the situation for topological defect models, where it has proven much harder to make accurate theoretical predictions. Less accurate theoretical predictions will naturally lead to much more poorly determined cosmological parameters. [In fact, Pen (these proceedings) has also argued that in a defect model the observed spectrum is less sensitive to the cosmological parameters, implying poorer parameter estimation even if the theoretical calculations can after all be made more accurate.] If all goes well with the observations, and inflation proves to be right, we can indeed look forward to the tiny error bars one hears about so often. If inflation is not correct, the results will still be spectacular but yet after the hype the constraints may seem disappointing. | 98 | 3 | astro-ph9803152_arXiv.txt |
9803 | astro-ph9803291_arXiv.txt | We report on the analysis of the ASCA and BeppoSAX X--ray observations of the polar system BL Hyi, performed in October 94 and September 96, respectively. Emission from both poles is apparent from the folded light curves of both observations; the emission from the second pole varies from cycle to cycle, indicating non--stationary accretion there. The temperature of the post--shock region is estimated to be about 10 keV. Inclusion of both complex absorption and Compton reflection significantly improves the quality of the fit. No soft X--ray component is observed; the BeppoSAX/LECS upper limit to the soft component is in agreement with theoretical expectations for this low magnetic field system. | Polar systems, a subgroup of magnetic Cataclysmic Variables (mCVs), contain a highly magnetized white dwarf with polar field strengths ranging from $\sim$ 10 MG $\sim$ to 230 MG (see Beuermann 1997 and references therein), and accreting material from a late type main sequence star. The magnetic field of the white dwarf is strong enough to phase--lock its rotation with the orbital period. These systems are strong X-ray emitters in both soft and hard X-ray bands (see review by Cropper 1990). While hard X-rays are emitted from a standing shock above the white dwarf surface, soft X-rays originate from hot photospheric regions heated either by irradiation from the post-shock plasma (Lamb \& Masters 1979) or by dense plasma blobs carrying their kinetic energy deep into the atmosphere (Kuipers \& Pringle 1982). Irradiation is important primarily for flow rates sufficiently low for the shock to stand high above the surface. However, a large fraction of the reprocessed radiation appears in the UV rather than at soft X-rays, as shown quantitatively in the case of AM\,Her (G\"ansicke et al. 1995). The ``blobby'' accretion mode is increasingly important at higher field strengths, because the blobs become more and more compressed when arriving at the surface of the white dwarf (e.g. Beuermann \& Burwitz 1995; King 1995 and references therein). While the soft X--ray component is in general adequately fitted by a black--body spectrum with a temperature of a few tens of eV (even if more sophisticated models are sometimes needed: see e.g. Van Teeseling et al. 1994), it has recently become clear that a simple thermal plasma model is no longer adequate in reproducing the hard X--ray component of Polars. Reflection from the white dwarf surface, complex absorption and multi--temperature emission may contribute significantly to the spectrum above a few tenth of keV (e.g. Cropper et al. 1997 and references therein). \begin{table*} \centering \caption{ASCA and BeppoSAX observations log} \label{log} \vspace{0.05in} \begin{tabular}{cccccc} \hline \hline ~ & ~ & ~ & ~ & ~ & ~\cr Instrument & En. range & Date of obs. & Exp. Time$^a$ & Count rate & 2-10 keV flux$^b$ \cr ~ & (keV) & ~ & (ks) & (cts/s) & (erg/s/cm$^2$) \cr \noalign {\hrule} ~ & ~ & ~ & ~ & ~ & ~\cr ASCA/GIS2 & 0.8-10 & 1994 Oct 11-12 & 43 & 0.13 & 8.0$\times$10$^{-12}$ \cr ASCA/GIS3 & 0.8-10 & 1994 Oct 11-12 & 43 & 0.16 & ~ \cr ASCA/SIS0 & 0.5-10 & 1994 Oct 11-12 & 42 & 0.21 & ~ \cr ASCA/SIS1 & 0.5-10 & 1994 Oct 11-12 & 41 & 0.16 & ~ \cr BeppoSAX/LECS & 0.1-10 & 1996 Sep 27 & 3.7 & 0.02 & ~ \cr BeppoSAX/MECS & 1.5-10 & 1996 Sep 27 & 12.5 & 0.12 & 6.8$\times$10$^{-12}$ \cr ~ & ~ & ~ & ~ & ~ & ~\cr \hline \hline $^a$ after applying the selection criteria\cr $^b$ single--temperature plasma model \end{tabular} \end{table*} X--ray emission from the polar system BL Hyi (H01319-68) was first detected by HEAO-1/A-2, and later on by Einstein (Agrawal et al. 1983), EXOSAT (Beuermann \& Schwope 1989, hereafter B\&S89) and ROSAT (Ramsay et al. 1996; Schwope et al. 1997). In particular, the HEAO-1 and EXOSAT observations showed that copious soft X-ray emission originates from the secondary pole in states of high accretion rate, while the EXOSAT observations showed this pole to become increasingly less active at lower accretion rates. This suggests a picture in which the source switches from one-- to two--pole accretion in going from low to high states (B\&S89). Recent EUVE observations indicate that the soft X--ray emission is best reproduced by a 17 eV blackbody (Szkody et al. 1997). Despite all the above observations, the hard X--ray spectrum of BL Hyi was still poorly known before the launch of ASCA and BeppoSAX satellites. ASCA observed this source in October 1994, while BeppoSAX observed it in September 1996, in the context of a Core Program devoted to study flux and spectral variability of Polars on different timescales and on a wide energy range. In this paper we present temporal and spectral analysis of both ASCA and BeppoSAX observations. Sec. 2 describes the observations and data reduction, while in Sec. 3 and Sec. 4 we present temporal and spectral results, respectively, which are then summarized and discussed in Sec. 5. | During the ASCA and BeppoSAX observations the source was significantly brighter than during the `intermediate state' of October 1985, when it was observed by EXOSAT (B\&S89). Assuming a 10 keV thermal plasma spectrum, the 2--10 keV phase--averaged flux, estimated with the {\sc pimms} package, was in fact $\sim$3$\times$10$^{-12}$ erg cm$^{-2}$ s$^{-1}$ during that EXOSAT observation, while it was $\sim$2.3 and 2.7 times brighter during the BeppoSAX and ASCA observations, respectively. The ASCA flux is in turn slightly lower than that measured by Ariel V ($\sim$10$^{-11}$ erg cm$^{-2}$ s$^{-1}$, Agrawal et al. 1983). Comparing both ASCA and BeppoSAX folded light curves with the EXOSAT ones (and in particular with that of October 1985, for which the statistics were the best), one important difference is apparent: there is significant emission from the second pole, while this emission during the EXOSAT observation was less intense, even if still present. As clearly evident from the BeppoSAX light curve, the emission from the second pole is highly variable, changing from cycle to cycle, and then suggesting non--stationary accretion. The ASCA and BeppoSAX hard X--ray spectra are the first of sufficient quality to reveal spectral complexities. A simple thermal plasma model, even if formally acceptable, is unsatisfactory due to both the presence of features (mostly at the lowest energies) in the residuals and the high temperature and metal (mainly iron) abundance obtained. The inclusion of complex absorption cures the low energy features. A partial covering model, with a column density of about 3$\times$10$^{21}$ cm$^{-2}$ and a covering fraction of $\sim$0.4 provides an acceptable description of the spectrum. It is conceivable that the absorber in front of the hard X-ray component is the free-falling matter feeding the accretion spot. For example in V834 Cen, this material was identified by its Zeeman component shifted in wavelength as expected from the free-fall velocity (Schwope \& Beuermann 1990). As the hard X-ray emitting post-shock plasma is probably located in front of the soft X-ray emitting sections of the photosphere, this absorbing material should affect the soft component as well. Due to substantial pre-ionization, however, it may be partially transparent to soft X-rays. On the other hand, the fact that Zeeman absorption from this material was seen in V834 Cen indicates that in that system neutral material is also present. Hence, the soft component may not emerge unhindered and its spectral properties may carry the signatures of the intervening absorber as well as those of the emitter. This has been taken into account when we have estimated the intensity of the soft component from the BeppoSAX data. The upper limit so obtained is about ten times that of the hard luminosity, if a black--body temperature of 17 eV (Szkody et al. 1997) is assumed; if, instead, a temperature of 25 eV (Ramsay et al. 1995) is adopted the upper limit is much more tight, reducing to about the hard luminosity. We also computed the upper limit to the black--body to hot plasma flux ratio in the 0.1--2.4 keV, to make comparison with the theoretical expectation (see Fig.~8 in Beuermann 1997). The limits are 5 and 1.3 for the adopted black--body temperatures, consistent with the comparatively hard X-ray spectra encountered in low-field systems (Beuermann \& Schwope 1994). This is especially true if the field strength of 12 MG measured by Schwope et al. (1995) is characteristic of the main accreting pole. Even a field strength of 23 MG reported by Ferrario et al. (1996) would still qualify BL Hyi as a low-field system, characterized by a strong hard X-ray component. Only at still higher field strengths is bremsstrahlung effectively suppressed by increasing cyclotron cooling (Woelk \& Beuermann 1996, see also Beuermann 1997). Finally, the inclusion in the spectral model of the reflection component from the white dwarf surface provides a significant improvement in the statistical quality of the fit and reduces the abundance to roughly the solar value. This component has been already detected in a number of Polars (e.g. Done et al. 1995; Beardmore et al. 1995; Ishida et al. 1997; Cropper et al. 1997; Done \& Magdziarz 1997) and is now recognized as an important ingredient of their X--ray spectrum. With all these components included, the best fit post--shock temperature results to be about 10 keV (with about 20\% uncertainty) in both observations, consistent with the EXOSAT results (B\&S89) for the same source and with the values usually found in Polars. | 98 | 3 | astro-ph9803291_arXiv.txt |
9803 | astro-ph9803258_arXiv.txt | We derive a self similar solution for the propagation of an extreme relativistic (or Newtonian) radiative spherical blast wave into a surrounding cold medium. The solution is obtained under the assumption that the radiation process is fast, it takes place only in the vicinity of the shock and that it radiates away a fixed fraction of the energy generated by the shock. In the Newtonian regime these solutions generalize the Sedov-Taylor adiabatic solution and the pressure-driven fully radiative solution. In the extreme relativistic case these solutions generalize the Blandford-McKee adiabatic solution. They provide a new fully radiative extreme relativistic solution which is different from the Blandford-McKee fully radiative relativistic solution. This new solution develops a hot interior which causes it to cool faster than previous estimates. We find that the energy of the blast wave behaves as a power law of the location of the shock. The power law index depends on the fraction of the energy emitted by the shock. We obtain an analytic solution for the interior of the blast wave. These new solutions might be applicable to the study of GRB afterglow or SNRs. | Many astrophysical phenomena (SNRs, $\gamma$-ray bursts - GRBs, AGN hot spots, etc.) are believed to involve radiative shock waves. These shocks accelerate particles which emit the observed radiation. In particular, it is widely accepted that the recently discovered GRBs afterglow results from an emission by relativistic shocks, created by the interaction between an initial ejecta and the interstellar medium. The recent observations of GRB afterglow have lead to numerous attempts to model this phenomena. If the radiative mechanisms are slow compared to hydrodynamics time scale, the blast wave evolution is adiabatic. The propagation of such a blast wave is described be a self similar solution: The Sedov-Taylor solution (\cite{sedov,taylor}) describes the Newtonian regime and the Blandford \& McKee (1976) solution describes the extreme relativistic regime. We call the solution radiative if the radiative mechanisms are fast compared to the hydrodynamic time scale. A fully radiative solution is one in which all the energy generated by the shock is radiated away. \cite{OM88} have shown that if a fully radiative blast wave emits its energy only in the vicinity of the shock, it can be described by one of two possible self-similar solutions: The pressure driven snow-plow (PDS) or the momentum conserving snow-plow (MCS) solution. In the PDS solution a thin shell ``snow plows'' through the external medium, driven by the pressure of its hot, roughly isobaric, interior. In the MCS solution the interior has been cooled and a thin shell slows down while conserving momentum. \cite{bm} have found an extreme relativistic fully radiative solution. This solution describes a thin shell with a cold interior and it can be considered as the relativistic generalization of the Newtonian MCS solution. However, momentum is not conserved in this solution, as in the relativistic case one has to consider the momentum carried by the emitted radiation. These solutions are either adiabatic or fully radiative. It is likely that in some cases not all the energy would be radiated away, even though the cooling is fast. Such is the situation if the shock distributes the internal energy between the electrons and protons and there is no coupling between the two afterwards. Since only the electrons cool, only a fraction of the internal energy will be radiated. It is likely that at least in the initial phases of a GRB afterglow this would be the case. We consider here this ``partially radiative'' scenario. The goal of our paper is to find an analytic solution for the evolution of Newtonian and relativistic partially radiative blast waves. This, we believe, will eventually lead to a physical description of the evolution of GRB afterglows in their non-adiabatic stage, and of other astrophysical phenomena. We find a self similar solution under the assumption that the radiative mechanism is fast (compared to the hydrodynamics time scale). We parameterize the radiation by a dimensionless (the non dimensionality is essential for a self-similar solution) parameter, $\epsilon$, which describes the fraction of the energy produced by the shock that is radiated away. We recover the extreme relativistic adiabatic Blandford-McKee solution and the Newtonian Sedov-Taylor solution in the corresponding limit when $\epsilon = 0$. In the Newtonian case when $\epsilon \to 1$ we recover the PDS solution, but we do not reproduce the MCS solution. Similarly in the relativistic limit of $\epsilon \rightarrow 1$ the solution approaches a new relativistic PDS solution which is different from the radiative Blandford-McKee solution. We describe our model in Sec. \ref{s:model}. It is composed of an adiabatic shock followed by a narrow radiative region and a wide self-similar adiabatic regime. We proceed in Sec. \ref{s:newtonian_cond} by calculating the radiative shock conditions for Newtonian blast waves. We describe the self similar equations and their solution in Sec. \ref{s:newtonian_self}. Using the same method we turn in Sec. \ref{s:rel} to the calculations of relativistic blast waves. Finally, in Sec. \ref{s:others_results} we discuss the limiting solutions obtained by other authors and compare them with our solution. | We have solved the hydrodynamical evolution of blast waves in which each shocked particle emits a fixed fraction of the energy it gains in the shock. We have divided the blast wave into three regions: the adiabatic shock, the radiative layer, and the adiabatic interior. The adiabatic shock and the radiation layer were combined to form a ``radiative shock'', which set the boundary conditions for the adiabatic interior. For fast cooling cases we have obtained a solution for a planar radiative layer with arbitrary shock velocities, independent of the cooling process. We have found that in the shock frame the fluid cools, slows down and its pressure {\it increases} during the cooling process. We have obtained self similar solutions for adiabatic interior for the Newtonian and the extreme relativistic cases. We find that radiation can change the hydrodynamics considerably. Because of this change in the blast wave evolution the luminosity decays slower (in time) than what was estimated earlier assuming that the self similar profiles are independent of the radiated energy. We have also found that as a blast wave becomes more radiative, its matter concentrates near the shock and forms a dense shell. However, the internal pressure does not drop to zero. In the fully radiative limit of Newtonian blast waves we have reached the pressure-driven solution. We have obtained a new extreme relativistic solution for fully radiative blast wave, which resembles the Newtonian PDS solution. Our self similar solutions reach this modified solution in the fully radiative limit. This solution does not correspond to the Blandford-McKee radiative solution. The pressure must be continuous within the adiabatic interior and it cannot develop self a similar shock or a rarefraction wave (apart from the main strong shock with the ISM). Therefore, even in the Newtonian PDS limit with isobaric interior, the solution contains a self similar transition layer between the ``radiative shock'' conditions and the internal pressure. The recently discovered X-ray, optical and radio emission following a GRB, so called ``Afterglow'' is widely believed to be the result of the deceleration of a relativistic material that collides with the surrounding matter. According to the common model, the shock wave produced by the collision accelerates electrons to relativistic velocities. These electrons, that carry a fixed fraction of the internal energy produced by the shock emit synchrotron radiation which is the observed afterglow. For reasonable parameters, the electron cooling is fast (at least in the early stages of the afterglow evolution), so that the electrons loose most of their energy. Since the density behind the shock is low, the coupling between the electrons and the protons is negligible. Therefore, the protons energy is not radiated away. This leads to a partially radiative rather than a fully radiative shock, followed by an adiabatic flow. This is {\it exactly} the scenario that leads to the partially radiative blast waves derived in this paper. Our new self similar partially radiative blast wave can serve as the basic hydrodynamic solution from which spectra and light curves of the afterglow can be calculated. In particular we derive a new $\Gamma(R)$ relation (Eqs. \ref{r:m_of_eps}, \ref{G_of_t} ) between the Lorentz factor and the radius. This is the critical relation that determines most afterglow observations. This research was supported by a US-Israel BSF grant 95-238 and by a NASA grant NAG5-3516. Re'em Sari thanks the Clore Foundations for support. | 98 | 3 | astro-ph9803258_arXiv.txt |
9803 | hep-ph9803500_arXiv.txt | The supersymmetric generalization of the limiting and Gaussian QCD spectra is obtained. These spectra are valid for $x \ll 1$, when the main contribution to the parton cascade is given by gluons and gluinos. The derived spectra are applied to decaying superheavy particles with masses up to the GUT scale. These particles can be relics from the Big Bang or produced by topological defects and could give rise to the observed ultrahigh energy cosmic rays. General formulae for the fluxes of protons, photons and neutrinos due to decays of superheavy particles are obtained. | The spectra of hadrons produced in deep-inelastic scattering and $e^+e^-$ annihilation are formed due to QCD cascading of the partons. In the Leading Logarithmic Approximation (LLA) which takes into account $\ln(Q^2)$ terms this cascade is described by the Gribov-Lipatov-Altarelli-Parisi-Dokshitzer (GLAPD) equation \cite{GLAPD}. The Modified Leading Logarithmic Approximation (MLLA) takes into account additionally $\ln(x)$ terms, where $x=k_{\parallel}/k_{\parallel}^{max}$ and $k_{\parallel}$ is the longitudinal momentum of the produced hadron. Color coherence effects are described in MLLA. Two approximate analytic solutions to the MLLA evolution equations have been obtained. % These are the {\em limiting spectrum\/} \cite{MLLA} and the {\em Gaussian spectrum\/} \cite{Mu,DFK}, in which we include the {\em distorted Gaussian spectrum\/} \cite{distG,dkmt,esw}. The limiting spectrum is the most accurate one among them. In fact, it describes well the experimental data at large $x$, too, and this is natural, though accidental (for an explanation see Ref.~\cite{dkmt}). The limiting spectrum has a free normalization constant, $K_{\rm lim}$, which cannot be calculated theoretically and is found from comparison with experimental data. This constant has to be considered as a basic parameter of the theory, and it can be used at all energies, where the physical assumptions, under which the limiting spectrum is derived, are valid. For detailed calculations of hadron spectra in $e^+e^-$ annihilation and comparison with experimental data see Ref.~\cite{comparison}. Up to energies of existing $e^+e^-$ colliders, $\sqrt{s} \lsim 183$~GeV, the limiting spectrum and the distorted Gaussian spectrum describe well the available data. At large energies $\sqrt{s}\gsim 1$~TeV the production of supersymmetric particles might substantially change the QCD spectra. Apart from future experiments at LHC, supersymmetry (SUSY) might strongly reveal itself in the decays of superheavy particles. They can appear as relics of the Big Bang, or be produced by topological defects (TD), and can be the sources of the observed ultrahigh energy cosmic rays (UHECR) at $E \gsim 1\cdot 10^{10}$~GeV. The range of masses, $m_X$, of interest for UHECR goes from the GUT scale ($m_X \sim 10^{16}$~GeV) or less down to $m_X \sim 10^{12} - 10^{14}$~GeV. In this paper we obtain the generalization of the limiting and Gaussian QCD spectra for the SUSY case. Although the influence of SUSY on, e.g., the evolution of parton distributions or the running coupling constant was considered in many works in the 80's \cite{split}, this is, to the best of our knowledge, the first time that fragmentation spectra of hadrons are examined for large $\sqrt{s}$ up to the GUT scale in SUSY-QCD. As application we use these spectra for calculations of the fluxes of UHECR produced in the decays of superheavy particles. | 98 | 3 | hep-ph9803500_arXiv.txt |
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9803 | astro-ph9803203_arXiv.txt | The white dwarf stars WD\,1614$+$136 and WD\,1353$+$409 are not sufficiently massive to have formed through single star evolution. However, observations to date have not yet found any evidence for binarity. It has therefore been suggested that these stars are the result of a merger. In this paper we place an upper limit of $\approx 50\kms$ on the projected rotational velocities of both stars. This suggests that, if these stars are the results of a merger, efficient angular momentum loss with accompanying mass loss must have occurred. If the same process occurs following the merging of more massive white dwarf stars, the predicted rate of Type~Ia supernovae due to merging white dwarfs may have been greatly over-estimated. Further observations to determine binarity in WD\,1614$+$136 and WD\,1353$+$409 are therefore encouraged. | The finite age of the universe limits the time available for stars to evolve. The lowest mass stars have had insufficient time to evolve to the point where the remnant is seen as a white dwarf star. There is a monotonically increasing relationship between the initial mass of the star and the resulting white dwarf star and so this implies a minimum mass for white dwarf stars. The exact value is uncertain, but is around 0.55\Msolar (Bragaglia, Renzini \& Bergeron, 1995). Nevertheless, white dwarf stars are found with masses well below this limit. These white dwarfs are the consequence of binary star evolution in which the evolution of a star through the red giant phase is interrupted by a common--envelope phase in which a companion star is engulfed by the expanding envelope and then rapidly spirals in towards the core, ejecting the envelope. This arrests the formation of the degenerate red giant core resulting in an anomalously low mass white dwarf star. Dramatic evidence for this scenario was provided by Marsh, Dhillon \& Duck \shortcite{Marsh95}. They observed 7 DA white dwarfs selected for the low mass derived for them by Bergeron, Saffer \& Liebert \shortcite{Berg92} from their spectra. Marsh et~al.\ were able to measure radial velocities with accuracies of a few \kms\ by using the narrow core of the \halpha\ absorption line. Periodic radial velocity variations showed at least 5 of the 7 stars to be binary stars. No evidence for binarity was found for WD\,1614$+$136 or WD\,1353$+$409. It should be emphasized that the observations of Marsh et~al.\ cannot rule out the possibility that these two white dwarf stars are binaries. Although a main--sequence companion more massive than 0.1\Msolar\ can be ruled out in both cases, the presence of another cool white dwarf, very low mass M dwarf or a brown dwarf, perhaps in a low inclination orbit, cannot be ruled out. With such strong evidence for the scenario outlined above it would seem to be inevitable that these white dwarfs were once members of binary systems. The nature of any companion, or its fate if it is no longer present, remain open questions. The failure of Marsh et~al.\ to detect binarity in WD\,1614$+$136 and WD\,1353$+$409 led Iben, Tutukov \& Yungleson \shortcite{Iben97} to suggest these stars are now single stars that are the result of a merger between a white dwarf and the companion responsible for the common--envelope phase. In this paper we show that the detection of a narrow core to the \halpha\ makes this suggestion very unlikely unless angular momentum loss from the merger product is extremely efficient. | If WD\,1614$+$136 and WD\,1353$+$409 are now genuinely single stars, their low projected rotational velocity implies efficient angular momentum loss following the merging of these stars with the companion star responsible for the common envelope phase. If the accompanying mass loss occurs following the merging of more massive white dwarfs, their role as the progenitors of Type~Ia supernovae is put into doubt. | 98 | 3 | astro-ph9803203_arXiv.txt |
9803 | astro-ph9803035_arXiv.txt | }[2]{{\footnotesize\begin{center}ABSTRACT\end{center} \vspace{1mm}\par#1\par \noindent {\bf Key words:~~}{\it #2}}} \newcommand{\TabCap}[2]{\begin{center}\parbox[t]{#1}{\begin{center} \small {\spaceskip 2pt plus 1pt minus 1pt T a b l e} \refstepcounter{table}\thetable \\[2mm] \footnotesize #2 \end{center}}\end{center}} \newcommand{\TableSep}[2]{\begin{table}[p]\vspace{#1} \TabCap{#2}\end{table}} \newcommand{\TableFont}{\footnotesize} \newcommand{\TableFontIt}{\ttit} \newcommand{\SetTableFont}[1]{\renewcommand{\TableFont}{#1}} \newcommand{\MakeTable}[4]{\begin{table}[htb]\TabCap{#2}{#3} \begin{center} \TableFont \begin{tabular}{#1} #4 \end{tabular}\end{center}\end{table}} \newcommand{\MakeTableSep}[4]{\begin{table}[p]\TabCap{#2}{#3} \begin{center} \TableFont \begin{tabular}{#1} #4 \end{tabular}\end{center}\end{table}} \newenvironment{references}% { \footnotesize \frenchspacing \renewcommand{\thesection}{} \renewcommand{\in}{{\rm in }} \renewcommand{\AA}{Astron.\ Astrophys.} \newcommand{\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \newcommand{\ApJ}{Astrophys.\ J.} \newcommand{\ApJS}{Astrophys.\ J.~Suppl.~Ser.} \newcommand{\ApJL}{Astrophys.\ J.~Letters} \newcommand{\AJ}{Astron.\ J.} \newcommand{\IBVS}{IBVS} \newcommand{\PASP}{P.A.S.P.} \newcommand{\Acta}{Acta Astron.} \newcommand{\MNRAS}{MNRAS} \renewcommand{\and}{{\rm and }} {We present a new distance determination to the Large and Small Magellanic Clouds using the newly developed red clump stars method (Paczy\'nski and Stanek 1998). This new, single-step, Hipparcos calibrated method seems to be one of the most precise techniques of distance determination with very small statistical error due to large number of red clump stars usually available. The distances were determined independently along four lines-of-sight located at opposite sides of each Magellanic Cloud. The results for each line-of-sight are very consistent. For the SMC we obtain the distance modulus: $m-M=18.56\pm0.03\pm0.06$~mag (statistical and systematic errors, respectively) and for the LMC: ${m-M=18.08\pm0.03\pm0.12}$~mag where systematic errors are mostly due to uncertainty in reddening estimates. Both distances will be refined and systematic errors reduced when accurate reddening maps for our fields are available. Distance moduli to both Magellanic Clouds are ${\approx0.4}$~mag smaller than generally accepted values. The modulus to the LMC is in good agreement with the recent determinations from RR~Lyrae type stars and upper limit resulting from the SN1987A echo. We suspect that the distance to the LMC and SMC is shorter by about 15\% than previously assumed: 42~kpc and 52~kpc, respectively. Calibrations of the period-luminosity relation for Cepheids which give overestimated distances to the LMC and SMC are probably incorrect and require urgent reanalysis. We also present our color-magnitude diagrams around the red clump for the LMC and SMC. We identify vertical red clump, first noted by Zaritsky and Lin (1997), in the color-magnitude diagram of both Magellanic Clouds and we interpret it as an evolutionary feature rather than unknown stellar population between the LMC and our Galaxy.}{Magellanic Clouds -- Galaxies: distances and redshifts -- Hertzsprung-Russel (HR) diagram} | % Two of our closest neighboring galaxies from the Local Group -- the Magellanic Clouds -- belong to the most important objects of the modern astrophysics. Hosting many objects commonly used as standard candles at practically the same distance they serve as ideal targets for calibrating the distance scale in the Universe. Therefore it is crucial to have well established distances to both Magellanic Clouds. Until recently, that is in the pre-Hipparcos era, it was generally accepted that the distance modulus to the Large Magellanic Cloud was in the range of ${m{-}M{=}(18.5{\div}18.6)}$~mag based mostly on the calibration of the period-luminosity (P--L) relation for Cepheids (\eg Laney and Stobie 1994, Gieren \etal 1997 and references therein) and measurements of the supernova SN1987A echo (\cf Panagia \etal 1991, Sonneborn \etal 1997). In 1992 Walker (1992) noted ${\approx0.3}$~mag discrepancy between absolute magnitudes of RR~Lyrae type stars from the Galaxy and the LMC at distance of 18.5~mag. Based on the new statistical parallax solution for RR~Lyrae, Layden \etal (1997) recalibrated the absolute magnitude-metallicity relation for these stars and found that the resulting distance modulus to the LMC might be significantly lower: ${m- M=18.28 \pm0.13}$~mag. The latest determinations of the distance to the LMC come from Hipparcos recalibration of the P--L relations for different type of pulsating stars. Both Cepheid and Mira recalibrations lead to larger distance moduli: Cepheids -- ${18.70\pm0.1}$~mag (Feast and Catchpole 1997) or ${18.57\pm0.11}$~mag (Madore and Freedman 1998) and Miras -- ${18.54\pm0.18}$~mag (van Leeuven \etal 1997). The value ${m-M=18.50\pm0.10}$~mag has been adopted by the HST Extragalactic Distance Scale Key Project team (\eg Ra\-wson \etal 1997 and other papers of the series) as the distance to the LMC to which extragalactic, Cepheid-based distance scale is tied. On the other hand Hipparcos recalibration for RR~Lyrae based both on direct parallax determination and statistical parallaxes gives results fully confirming Layden \etal (1997) determination: ${18.26\pm0.15}$~mag (Fernley \etal 1998). Similar results from statistical parallaxes were recently obtained by Popowski and Gould (1998). Thus it seems very likely that the RR~Lyrae stars give ${\approx0.25}$~mag shorter distance scale to the LMC. Meanwhile, recent reanalysis of the supernova 1987A echo by Gould and Uza (1998) also suggests smaller distance modulus than previously derived ({\it cf.} Sonneborn \etal 1997, Lundquist and Sonneborn 1997) with the upper limit as low as ${m-M<18.37\pm0.04}$~mag. The distance to the Small Magellanic Cloud is poorly known. The distance modulus to the SMC is assumed to be about 18.9~mag (Westerlund 1990). Massey \etal (1995) obtained ${19.1\pm0.3}$~mag from spectroscopic parallaxes. The most recent determination based on the Cepheid period-luminosity relation gives similar value: ${m-M=18.94}$~mag (Laney and Stobie 1994). Summarizing this short review, the distance to the Large Magellanic Cloud is at least controversial and the generally accepted value $m-M\approx18.5$~mag cannot be now treated as established beyond any doubt. The distance to the Small Magellanic Cloud is assumed to be ${m-M\approx18.9}$~mag, still poorly known, and therefore any new, independent information about both distances is of the greatest importance. Recently Paczy\'nski and Stanek (1998) proposed a new, single step method of distance scale determination to the objects in the Galaxy and neighboring galaxies. The method bases on the fact that the red clump giant stars seem to be a very homogeneous group of objects and their mean magnitude in the {\it I}-band is practically independent of their color. Therefore the red clump stars might be considered as excellent standard candle candidates. The red clump stars are very numerous compared with other objects used as standard candles so far (typically several orders of magnitude more numerous than \eg pulsating stars). Thus, their mean magnitudes can be obtained very precisely with small statistical error making the method very attractive for precise determination of distances in the Universe. The mean absolute magnitude of the red clump stars from the solar vicinity was determined by Paczy\'nski and Stanek (1998) who analyzed luminosities of about 600 such stars with precise parallaxes (accuracy better than 10\%) measured by Hipparcos (${M^{\rm loc}_{I_0}=-0.185\pm0.016}$~mag). Paczy\'nski and Stanek (1998) applied the method for determination of the distance to the center of the Galaxy comparing $M^{\rm loc}_{I_0}$ with the mean {\it I}-band magnitude of the red clump stars from the Baade's Window obtained from photometry of the first phase of the Optical Gravitational Lensing Experiment (OGLE). The new method was also applied by Stanek and Garnavich (1998) for determination of the distance to M31. Stanek and Garnavich (1998) also refined the maximum of the local red clump stars luminosity function by limiting the Hipparcos sample of the red clump stars to be volume rather than luminosity limited and derived ${M^{\rm loc}_{I_0}=-0.23\pm 0.03}$~mag from 228 red clump stars located within the distance ${d<70}$~pc. Very large number of red clump stars in the Hipparcos sample makes the calibration very sound comparing with calibrations of other standard candles. It should be noted that this calibration might be even more precise in the future when good {\it VI} photometry is derived for more Hipparcos stars. Although employing the red clump stars as distance indicator seems to be very attractive and can be possibly one of the most precise methods of distance determination one should be aware of potential systematic errors which can lead to errors in determined distance scale. The most important is proper determination of the interstellar extinction which can affect both target red clump luminosity as well as the local Hipparcos sample. However, the latter sample is very likely extinction free for ${d<70}$~pc or it is affected negligibly, \eg that one analyzed by Stanek and Garnavich (1998). Nevertheless precise determination of extinction to the target stars is very important to avoid large systematic errors. The second source of errors can be the differences in populations -- ages, chemical composition etc., between the red clump stars of the target object and those in the solar neighborhood. This is, however, a common problem of any other method of distance determination when similar objects from different locations of the Universe are compared. In case of the red clump stars theoretical models predict that their luminosity is very weakly dependent on chemical composition and age (Castellani, Chieffi and Straniero 1992, Jimenez, Flynn and Kotoneva 1998). Nevertheless, the empirical fact of the very small dispersion of {\it I}-band magnitudes of the red clump stars as observed in the solar vicinity, Baade's Window and M31 (Paczy\'nski and Stanek 1998, Stanek and Garnavich 1998) should find theoretical explanation which could also answer the question about uncertainties introduced by comparison of different populations of the red clump stars. The natural consequence of the successful application of the red clump method for determination of distances to the Galactic center (Paczy\'nski and Stanek 1998) and M31 (Stanek and Garnavich 1998) is a distance determination to other objects in the Local Group. The most natural candidates are the Magellanic Clouds, which are targets of the microlensing surveys providing huge databases of photometric data for millions of stars in both Magellanic Clouds. The OGLE project is a long term microlensing survey which second phase -- OGLE-II started at the beginning of 1997 (Udalski, Kubiak and Szyma\'nski 1997). The Magellanic Clouds are the new targets of the OGLE-II phase. Unlike other microlensing surveys the OGLE project observations are carried out with the standard {\it BVI}-bands with majority (${\approx75\%}$) observations obtained in the {\it I}-band. Thus, the data can be precisely transformed to the standard {\it BVI} system and applied directly to many projects unrelated to microlensing. In particular the OGLE observations are very well suited for the above mentioned red-clump method as the most of data is collected in the {\it I}-band. In this paper we apply the red clump stars method for determination of distances to both Large and Small Magellanic Clouds. Although OGLE-II data cover large fractions of both Magellanic Clouds we present here distance determination for two most west- and eastward lines-of-sight toward both Clouds where we assume extinction to be constant in the first approximation. Thus, the results should be considered as preliminary and will be refined when accurate maps of extinction based on the OGLE-II data are derived. | % Results of distance determination to the Magellanic Clouds with the red clump stars method (Table~3) lead to the surprising conclusion. All of our determinations are consistently ${\approx0.4}$~mag smaller than generally accepted distance moduli to the Magellanic Clouds. There might be a few reasons of the discrepancy: -- error of the zero point of the OGLE photometry. This is, however, extremely unlikely. In Section~2 we discussed accuracy of the OGLE photometry and showed comparison with other reliable CCD photometries of the Magellanic Clouds (\eg Fig.~1). The OGLE mean magnitudes are the average of tens of individual observations and were tied with the standard system based on observations collected on many nights with hundreds of observations of standard stars. We believe that our photometry constitutes a huge set of secondary {\it BVI} standards in both Magellanic Clouds and large systematic errors are excluded. -- errors in reddening estimates for our fields. Certainly overestimating of the reddening leads to smaller distance moduli. However, one should note that even with the zero reddening the distance moduli derived with the red clump method are smaller than generally accepted for the Magellanic Clouds (by at least 0.1~mag). But extinction to the Magellanic Clouds does exist and for the LMC the foreground extinction is at least ${E(B-V)=0.075}$~mag, ${A_I=0.15}$~mag, and to the SMC ${E(B-V)=0.037}$~mag, ${A_I=0.07}$~mag, (Schlegel, Finkbeiner and Davis 1998). Reddening to the SMC is much smaller and our distance determination is less affected by the reddening estimate. But consistently the distance modulus to the SMC is also smaller by ${\approx0.4}$~mag. -- the red clump stars method. We addressed this possibility in the Introduction. It seems unlikely from our present understanding of the red clump stars, but in principle, the mean {\it I}-band magnitude of the red clump stars might be a function of stellar population (although it shows very small dispersion in all objects investigated so far: local sample, Galactic bulge, M31 and Magellanic Clouds). We stress here once again that this problem must be carefully investigated both theoretically and observationally to answer potential questions. As any of the above reasons does not seem to be satisfactory, we come to the very tempting conclusion that the distance determination to the Magellanic Clouds with the red clump method is correct, the distance moduli are ${\approx 0.4}$~mag smaller than those generally accepted and both Magellanic Clouds are located about 15\% closer to our Galaxy than previously assumed: the LMC at 42~kpc, and the SMC at 52~kpc. Our conclusion seems to be a strong argument in favor of the shorter distance to the LMC. Our determination is based on the independent, very straightforward, single-step method with minimum possibilities for the potential error sources. The new technique is calibrated with Hipparcos data with much better accuracy than calibration of any other standard candle (orders of magnitude more stars, very close objects with precise parallaxes). Independent determinations for four lines-of-sight located at different sides of each Magellanic Cloud yield very consistent results which makes them very sound. If we underestimated the reddening, the Magellanic Clouds are even closer to us. On the other hand if we assume only minimum, foreground reddening, which is certainly an underestimate, ($E(B-V)=0.075\pm0.015$~mag for the LMC and $E(B-V)=0.037\pm0.015$~mag for the SMC, Schlegel, Finkbeiner and Davis 1998) the upper limits for distance moduli are: $m- M<18.29\pm0.03\pm0.03$~mag for the LMC and $m-M<18.65\pm0.03\pm0.03$~mag for the SMC. Still a few sigma gap between distance moduli resulting from the red clump method and presently accepted values remains. Brocato \etal (1996) presented the CCD photometry of selected globular clusters in the LMC. We find there that one of the CMDs, showing the region around NGC~1786, contains a neat red clump. Unfortunately Brocato \etal (1996) do not provide their photometry, but even visual inspection of their Fig.~1 indicates that the mean magnitude of the red clump stars is ${I=18.2\pm0.1}$~mag. NGC~1786 is located only one degree north from our field LMC$\_$SC15N so it is comparably affected by extinction (it falls almost into the extinction map of Harris, Zaritsky and Thompson 1997). Thus resulting distance modulus from NGC~1786 field and independent photometry is in perfect agreement with our red clump determinations in Table~3. As we mentioned in the Introduction, RR~Lyrae-based distance determination also seems to favor smaller distance modulus to the LMC. All recent determinations from statistical parallaxes and Hipparcos calibrations of RR~Lyrae stars are within one sigma agreement with our determination. Also the upper limit for the distance to the LMC from the SN1987A light echo by Gould and Uza (1998) is compatible with our determination. The only arguments for the larger distance to the LMC come then from the Cepheid P--L calibration (and the Mira variables but because of much lower accuracy this determination is much less reliable). However, because of its complex nature, the calibration might be simply incorrect. An incorrect account of possible population effects, mainly metallicity, could result in overestimated distances obtained with that method (\cf Sekiguchi and Fukugita 1998, Sasselov \etal 1997). Our SMC distance determination also confirms this thesis: our distance modulus to the SMC is ${\approx0.4}$~mag smaller than that derived from Cepheid P--L (Laney and Stobie 1994). It should be also noted that there exists ${\approx0.3}$~mag discrepancy between Hipparcos calibrated Cepheid P--L distance to M31 (Feast and Catchpole 1997) and recent determination with the red clump stars (Stanek and Garnavich 1998) and globular clusters (Holland 1998). Thus we would like to stress at this point the necessity of urgent reanalysis of the Cepheid P--L relation. Very precise, {\it BVI} light curves of hundreds of Cepheids from both Magellanic Clouds will be provided by the OGLE project to the astronomical community later during this year and can be of great importance for such a project. Concluding, it seems very likely that Cepheid P--L relation gives distance moduli overestimated by about ${\approx0.4}$~mag (at least toward the Magellanic Clouds) with all astrophysical consequences -- distances beyond M31, up to several Mpc rely mainly on Cepheids and the LMC distance. Our independent distance measurement to the LMC confirms the recent calibrations of the absolute magnitude--metallicity relation for RR~Lyrae stars giving fainter absolute magnitudes contrary to results of Chaboyer \etal (1998) and therefore leading to older ages of globular clusters (Layden \etal 1996). On the other hand shorter distance to the LMC means shorter distance to Cepheid-based galaxies and in turn larger Hubble constant and shorter age of the Universe. The "age problem" gap increases. We believe that there exists at least one crucial test which can verify our present findings. Paczy\'nski (1997) proposed detached eclipsing binaries as an excellent method of distance determination. Having precise light curve and spectroscopic data one can determine the distance to the eclipsing, well detached system with accuracy of a few percent. The catalog of eclipsing stars in the Magellanic Clouds suitable for such a distance determination with three color light curves, periods etc.\ will be released by the OGLE project in a few months. Faintness of the Magellanic Cloud stars will require the largest, 10-m class telescopes to collect good quality spectroscopic data and a lot of telescope time. However, importance of the problem makes such a project the one of the highest priority and we hope it will be undertaken soon. \Acknow{We would like to thank Prof.\ Bohdan Paczy\'nski for encouraging discussions and help at all stages of the OGLE project. We are indebted to Dr Krzysztof Stanek for useful coments and remarks. The paper was partly supported by the Polish KBN grant 2P03D00814 to A.\ Udalski. Partial support for the OGLE project was provided with the NSF grant AST-9530478 to B.\ Paczy\'nski.} \bigskip {\bf Note:} Based on theoretical isochrones fitting to the main sequence and the red clump in four fields in the LMC, Beaulieu and Sackett (1998, revised version of their paper) found significantly better fit assuming smaller distance modulus to the LMC (${m-M=18.3}$) providing thus, additional argument in favor of the shorter distance to the LMC. \newpage | 98 | 3 | astro-ph9803035_arXiv.txt |
9803 | astro-ph9803345_arXiv.txt | IRAS~06562$-$0337 has been the recent subject of a classic debate: proto-planetary nebula or young stellar object? We present the first 2$\micron$ image of IRAS~06562$-$0337, which reveals an extended diffuse nebula containing approximately 70 stars inside a 30$\arcsec$ radius around a bright, possibly resolved, central object. The derived stellar luminosity function is consistent with that expected from a single coeval population, and the brightness of the nebulosity is consistent with the predicted flux of unresolved low-mass stars. The stars and nebulosity are spatially coincident with strong CO line emission. We therefore identify IRAS~06562$-$0337 as a new young star cluster embedded in its placental molecular cloud. The central object is likely a Herbig Be star, $M \approx 20 M_{\odot}$, which may be seen in reflection. We present medium resolution, high S/N, 1997 epoch optical spectra of the central object. Comparison with previously published spectra shows new evidence for time variable permitted and forbidden line emission, including \ion{Si}{2}, \ion{Fe}{2}, [\ion{Fe}{2}], and [\ion{O}{1}]. We suggest the origin is a dynamic stellar wind in the extended, stratified atmosphere of the massive central star in IRAS~06562$-$0337. | Garcia-Lario, Manchado, Sahu, and Pottasch (1993, hereafter GMSP) present the first detailed analysis of IRAS~06562$-$0337. They argue that it is a proto-planetary nebula (PPN) undergoing final mass-loss episodes. Their time-series of optical spectra, obtained over a 5 year period, show the onset of forbidden line emission and the possible evolution of the central star toward hotter temperatures. They derive a Zanstra temperature of 2$\times$10$^{4}$ K, with a slight increase over a two year interval. The effective temperature of the exciting star, T$_{eff}$ $\approx$ 3.6$\times$10$^{4}$ K, also showed a slight increase in two years. The H$\alpha$ line profile changes in time, which GMSP interpret as variable high velocity winds associated with episodic mass-loss. The appearance of [\ion{O}{3}] emission lines in 1990 and the resulting 4363/(4959+5007) line ratio requires an ionizing region of high electron density, log($n_{e}$) $\approx$ 6.9. The absence of these lines in spectra obtained before and after 1990 is interpreted as collisional de-excitation due to changing densities in the ionized region effected by violent episodic mass-loss. From CO observations GMSP derive $V_{LSR}$ = 50 $\pm$ 1 km sec$^{-1}$, which agrees with the velocity derived from their high resolution optical spectra. Adopting a model galactic rotation curve, they estimate a distance of 4 kpc, which compares to a distance of 2.4 kpc estimated from the equivalent width of Na D absorption seen in their spectra. The IRAS colors fit with blackbodies show a trend of decreasing temperature with increasing wavelength which implies a gradient of dust temperatures. GMSP integrated the optical--IR spectral energy distribution of IRAS~06562$-$0337, yielding a luminosity of L = 7000 $L_{\odot}$ for their preferred distance of 4 kpc. Kerber, Lercher, and Roth (1996, hereafter KLR) describe an additional medium resolution, high S/N spectrum of IRAS~06562$-$0337 obtained in early 1996. [\ion{O}{3}] emission is still absent, but a wealth of \ion{Fe}{2} and [\ion{Fe}{2}] lines are found. These lines confirm the high electron density derived by GMSP from the [\ion{O}{3}] lines present in 1990. KLR argue that the spectrum also implies a considerable density gradient in the object, as [\ion{Fe}{2}] lines are collisionally suppressed at densities where \ion{Fe}{2} lines exist. They maintain the classification of IRAS~06562$-$0337 as a candidate PPN, designating it ``The Iron Clad Nebula''. Bachiller, Gutierrez, and Garcia-Lario (1998, hereafter BGG) present new mm and sub-mm observations of IRAS~06562$-$0337. They derive $V_{LSR}$ = 54.0 $\pm$ 0.2 km sec$^{-1}$ and adopt a different model Galactic rotation curve than GMSP to estimate a distance of 7 kpc. This distance yields a luminosity of 21000 $L_{\odot}$ and a cloud mass M $>$ 1000 $M_{\odot}$. From the strength of the CO emission and the presence of CS emission, BGG surmise that IRAS~06562$-$0337 is a ``young stellar object (or small cluster) still associated to its parent molecular cloud.'' BGG point out that IRAS~06562$-$0337 satisfies the three criteria for a Herbig Ae/Be star (Herbig, 1960) and the spectral energy distribution, which rises sharply in the far infrared (GMSP), is similar to Group II Herbig Ae/Be stars (Hillenbrand, 1992). They also note the presence of blue and redshifted wings in the CO emission indicating a bipolar outflow. The CO outflow may be driven by an eruptive ionized jet, which leads them to suggest the sporadic [\ion{O}{3}] emission seen by GMSP originates in a Herbig-Haro object. We present the first 2$\micron$ image of IRAS~06562$-$0337. Our image reveals a compact cluster of stars surrounding a bright, central object. We independently confirm the result also discovered by BGG that IRAS~06562$-$0337 is a young stellar object. In Section 2 of this paper, we describe our near-infrared observations and stellar census of the IRAS~06562$-$0337 cluster. We also compare the CO(2$\rightarrow$1) map of BGG with our image. In Section 3, we describe our new spectroscopic observations and summarize the resulting 1997 epoch emission line data. We make a detailed comparison with the 1996 epoch emission line data of KLR. In Section 4, we present our conclusions. | The bright central object in IRAS 06562$-$0337 is probably a single Herbig Ae/Be class star, with $M \approx 20 M_{\odot}$. There is some evidence that it is seen at least partially in reflection. For example, the central star is poorly fit by the stellar point spread function derived from other stars in our K$^{\prime}$ image. In addition, the centroids in the visible and at K$^{\prime}$ are offset by $\Delta\alpha \sim$0.6$\arcsec$. A comparison of our 1997 epoch optical spectra to the 1996 epoch spectrum presented by KLR has provided new evidence for spectral line variability in IRAS~06562$-$0337 that is not easily interpreted as simple changes in temperature or density. This argues against previous suggestions that the already-published spectral data reflect real-time evolutionary changes of the central star. Understanding the immediate environs of the central star is an essential first step to better understanding the emission line variability that has been observed. We note that our analysis of the 1996--1997 spectral variability of IRAS 06562$-$0337 using simple zones and changes in one parameter (such as the electron temperature or electron density) is only intended as an interpretive tool and should be treated with some caution. The inconsistencies of this analysis suggest the complexity of the actual emission line region. If the central star is analagous to a Herbig Be star, then its ``fully ionized zone'' should have a very high electron density. Thus, the low density derived with the [\ion{O}{1}] emission implies a stratification of densities, likely in the extended atmosphere of the central star in IRAS~06562$-$0337. The anti-correlation of [\ion{O}{1}] and [\ion{Fe}{2}] emission is not easily understood if they arise from a common partially ionized zone (perhaps located at the edge of a compact \ion{H}{2} region around the central star), but is likely consistent with an origin in a dynamic stratified atmosphere. Spectral variability is also consistent with a stellar wind origin. Stellar winds from Herbig Be type stars are highly variable, extremely complex, and poorly understood (e.g., they are possibly modulated by magnetic fields or non-radial pulsations; Catala et al. 1993). Wind line strengths and profiles can show nightly variations, some of which are correlated with temperature changes of the star (Scuderi et al. 1994). The appearance of [\ion{O}{3}] emission in IRAS 06562$-$0337 during 1990 (GMSP) can be understood as shocks in stellar jets from a variable stellar wind (Hartigan \& Raymond 1993). Additionally, we emphasize that the size of the ionized region in IRAS 06562$-$0337 is small. Six cm radio observations yielded a surprising non-detection given the strong H$\beta$ flux, and allow an upper limit angular diameter of $3.4\times10^{-4}$ arcsec to be set for the ionized region (GMSP). If we assume the 7 kpc distance estimate of BGG, then the ionized region in IRAS 06562$-$0337 has an upper limit diameter of 2--3 AU. We conclude that stellar winds in the extended, stratified atmosphere of a very young, $M \approx 20 M_{\odot}$ star is the most plausible explanation for the puzzling spectral variability of IRAS~06562$-$0337, the ``Iron Clad Nebula.'' A valuable time series of spectra of the central object in IRAS 06562$-$0337, covering a 10 year interval, are now available in the literature. A useful complementary study would obtain more densely sampled spectroscopic observations, to investigate the short timescale behavior of the emission. In particular, correlated line variability and characteristic timescales could be used to constrain the dimensions of the ionized region, which we have suggested is likely to be an extended stellar atmosphere. We emphasize the need for high resolution and high S/N spectroscopic observations in the near-UV, and of the entire Balmer series, as these data can be used to study the extended atmospheres of Be stars (Burbidge \& Burbidge 1953). It is important to build a consistent physical picture of the environs around the central star in order to answer questions such as ``is the central star seen in reflection?'' This problem in particular could be addressed with high resolution imaging and polarization data. Lastly, as a young star cluster, IRAS~06562$-$0337 is interesting and notable for its richness. Multi-color near-infrared photometric observations would allow an accurate measurement of the IMF, and an investigation into problems such as mass segregation and star formation efficiencies (e.g., Barsony et al. 1991). | 98 | 3 | astro-ph9803345_arXiv.txt |
9803 | astro-ph9803109_arXiv.txt | Using time-resolved two-dimensional aperture photometry we have put upper limits on the pulsed emission from two proposed optical counterparts for PSR B1951+32. Our pulsed upper limits of $m_{v,\rm{pulsed}} > 23.3$, $m_{b,\rm{pulsed}} > 24.4$, for the first candidate and $m_{v,\rm{pulsed}} > 23.6$, $m_{b,\rm{pulsed}} > 24.3$ for the second, make it unlikely that either of these is, in fact, the pulsar. We discuss three further candidates, but also reject these on the basis of timing results. A search of a $5\farcs5 \times 5\farcs5$ area centred close to these stars failed to find any significant pulsations at the reported pulsar period. | PSR B1951+32 is a fast pulsar in the peculiar combination supernova remnant (SNR) CTB 80. The characteristic spin-down age of the pulsar $(P/2\dot P)$ and the dynamical age of the SNR (Koo et al. 1990) are both $\sim10^5$ yrs. Its low surface magnetic field strength, of $~5\times10^7$ T, is unusual amongst the selection of young and middle-aged pulsars detected so far but it has been suggested that this is more representative of young pulsars in general (\cite{lyne83}). It is clearly important to test pulsar models with as wide a variety of pulsar parameters as possible. Steady emission from the pulsar has so far been detected at radio (\cite{strom87}; \cite{kul88}) and X--ray (\cite{beck82}; \cite{wang84}) wavelengths. Two possible optical counterparts have been proposed by Blair \& Schild (1985) and Fesen \& Gull (1985). Since the initial discovery of the 39.5-ms pulsar (\cite{kul88}), evidence for pulsed emission has also been found in X--rays (\cite{safi95}) and $\gamma$--rays (Ramanamurthy et al. 1995) with upper limits in the infrared (Clifton et al. 1988). To-date no upper limits on pulsed emission in the optical have been published. Optical observations are essential in determining the relative contributions of thermal and magnetospheric processes to pulsar emission. The continuing improvements in optical-detector sensitivities means that pulsed emission from 6 isolated neutron stars has already been detected. In this paper we present time-resolved optical observations of the central region of the supernova remnant CTB 80 and, in particular, of the two proposed pulsar candidates. | On the basis of luminosity, colour and timing we conclude that neither of the proposed optical counterparts are the pulsar PSR B1951+32. We note that the slight blue extension to candidate 1 shows the expected characteristics of a hot object. We would not expect to be sensitive to thermal radiation from the neutron star surface itself, as we estimate that a pulsar with T $\approx 10^6$ K at the distance of CTB 80 would have a magnitude in the range $m_v$=30--31. We may, however, be able to see emission from any hot circumstellar material. From phenomenolgical models of pulsar optical emission (Pacini and Salvati 1983, 1987) we can derive estimated magnitudes for the emission. For PSR 1951+32 this leads to values in the range 24--26. However recent observations of other middle-aged pulsars, 0656+14 and Geminga (\cite{car94}, \cite{shear98}), have indicated a magnitude considerbly in excess of that predicted by the Pacini and Salvati model. It is clearly very important to establish a positive optical identification for PSR B1951+32 in order to compare emission from pulsars of similar ages but with a range of other parameters such as $\dot{P}$ and magnetic field. These observations will be crucial if we are to distinguish between the various models for high-energy emission. In particular, the pulse shape and fraction can be used to differentiate the polar cap and outer gap models. We suggest observations with the VLT, Keck or the HST in order to definitively identify the optical counterpart of PSR B1951+32. | 98 | 3 | astro-ph9803109_arXiv.txt |
9803 | astro-ph9803279_arXiv.txt | The measurement of shape parameters of sources in astronomical images is usually performed by assuming that the underlying noise is uncorrelated. Spatial noise correlation is however present in practice due to various observational effects and can affect source shape parameters. This effect is particularly important for measurements of weak gravitational lensing, for which the sought image distortions are typically of the order of only 1\%. We compute the effect of correlated noise on two-dimensional gaussian fits in full generality. The noise properties are naturally quantified by the noise autocorrelation function (ACF), which is easily measured in practice. We compute the resulting bias on the mean, variance and covariance of the source parameters, and the induced correlation between the shapes of neighboring sources. We show that these biases are of second order in the inverse signal-to-noise ratio of the source, and could thus be overlooked if bright stars are used to monitor systematic distortions. Radio interferometric surveys are particularly prone to this effect because of the long-range pixel correlations produced by the Fourier inversion involved in their image construction. As a concrete application, we consider the search for weak lensing by large-scale structure with the FIRST radio survey. We measure the noise ACF for a FIRST coadded field, and compute the resulting ellipticity correlation function induced by the noise. In comparison with the weak-lensing signal expected in CDM models, the noise correlation effect is important on small angular scales, but is negligible for source separations greater than about $1^{\prime}$. We also discuss how noise correlation can affect weak-lensing studies with optical surveys. | The measurement of source morphologies in two-dimensional images is a fundamental problem in astronomy. The measurements of shape parameters are usually performed while assuming that the underlying noise is uncorrelated. However, spatial correlation of the noise is always present to some degree, and can significantly affect the derived parameters. In experimental situations, noise correlation can be produced by various effects such as convolution of background light with a beam (or point-spread function), interferometric imaging techniques, CCD readouts, etc. This effect is particularly important for measurements of weak gravitational lensing. Weak lensing provides a unique opportunity to measure the gravitational potential of massive structures along the line-of-sight (for reviews see Schneider et al. 1992; Narayan \& Bartelmann 1996). This technique is now routinely used to map the potential of clusters of galaxies (see Fort \& Mellier 1994; Kaiser et al. 1994, for reviews). Detections of the more elusive effect of lensing by large-scale structure have been reported by Villumsen (1995) and Schneider et al. (1997) in small optical fields. The search for a strong detection of the effect on larger angular scales is currently being attempted with present and upcoming wide-field CCDs in the optical band (e.g. \markcite{ste95}Stebbins et al. 1995; \markcite{kai96}Kaiser 1996; \markcite{ber97}Bernardeau et al. 1997), and with the FIRST radio survey (\markcite{kam98}Kamionkowski et al. 1998; \markcite{ref98}Refregier et al. 1998). The main challenge comes from the fact that weak lensing induces image distortions of only about 1\%, and thus requires high-precision measurements of source-shape parameters. Correlated noise can produce spurious image distortions and correlations in the shapes of neighboring sources, and therefore must be carefully accounted for in weak lensing studies. The effect of correlated noise is also particularly important for radio surveys performed with interferometric arrays. In such surveys, images are produced by Fourier inversion of the visibilities from a set of antenna pair. As a result of the incomplete visibility coverage, part of the noise on the image plane contains long range correlations which can extend across an entire field. The effect of these correlations is particularly relevant in the context of recent large radio surveys such as FIRST (\markcite{bec95}Becker et al. 1995; \markcite{whi96}White et al. 1996) and NVSS (Condon et al. 1997). While the present analysis was motivated by our attempt to measure weak lensing by large-scale structure with the FIRST radio survey, most of the following is general and can be applied to any wavelength and imaging technique. \markcite{con97}Condon (1997) computed the errors in gaussian fit parameters in the absence of noise correlation. He also gave a semi-quantitative treatment of the noise correlation which relies on simulations. In the present paper, we generalize his approach to include a general analytic treatment of the noise correlation. We show how the spatial correlation of the noise can be naturally characterized by the noise Auto-Correlation Function (ACF). We explicitly derive the effect of noise correlation on source parameters for a general two-dimensional fit and for a two-dimensional gaussian fit. We also consider correlations induced in the parameters of nearby sources. In particular, we focus on the ellipticity correlation functions, which are used in searches for weak lensing by large-scale structure. We apply our formalism to the case of the FIRST radio survey. After measuring the noise ACF for one FIRST coadded field, we compare the noise-induced ellipticity correlations to those expected for weak lensing by large-scale structure. We show that, while they are important on small angular scales, noise correlation effects are negligible for source separations greater than about $1^{\prime}$. This paper is organized as follows. In \S\ref{noise_characterization}, we define the noise ACF and show how it can be measured in practice. In \S\ref{general_fit}, we consider a general two-dimensional least-square fit. We derive the bias produced by the noise correlation on the mean, variance and covariance of the fit parameters, and the correlation of the parameters for two neighboring sources. In \S\ref{fit_gaussian}, we apply these results to the case of two-dimensional gaussian fits and compute the error matrix. We explicitly derive the ellipticity correlation function induced by the noise correlation. In \S\ref{first}, we consider the concrete case of the FIRST radio survey. We measure the noise ACF for a FIRST field, compute the induced ellipticity correlation function, and compare the latter to that expected for weak lensing. In \S\ref{optical}, we discuss the implications of this effect for optical surveys and related searches for weak lensing by large-scale structure. Finally, \S\ref{conclusion} summarizes our conclusions. | \label{conclusion} We have studied the effect of noise correlation on the shape parameters in two-dimensional images. The noise correlation is conveniently described by the noise ACF which can easily be measured in practice. We derived the magnitude of the effect for a general two-dimensional least-square fit. The noise correlation can produce a bias, and affect the variance and the covariance of source parameters. In addition, it can produce systematic correlation in the parameters of pairs of sources. We find that the effect is of second order in the inverse SNR of the sources. We applied these general results to the case of most practical interest, namely a two-dimensional gaussian fit. We computed the relevant matrices explicitly. In addition, we explicitly derived the systematic bias produced by this effect on the ellipticity correlation function of source pairs. This is particularly relevant for weak lensing studies. As a concrete example, we studied the effect of correlated noise on the shape of sources in the FIRST radio survey. We measured the noise ACF and found long range features which extend beyond the central maximum. We computed the resulting systematic effect on the ellipticity correlation function. We find that, in the FIRST survey, noise-induced ellipticity correlations dominate over the expected weak lensing signal for $\theta \lesssim 0.5^{\prime}$, but are negligible for $\theta \gtrsim 1^{\prime}$. We discussed the consequences of noise correlation for optical surveys. In optical images, noise correlation can arise from various effects such as the convolution of background light with the PSF, CCD read outs, shift-and-add preprocessing, drift-scanning, etc. Because the effect is quadratic in the source SNR, the effect could be overlooked if systematic distortions are monitored solely with bright stars. The effect of noise correlation could thus be important for searches of galaxy-galaxy lensing and of weak lensing by large-scale structure in optical surveys. | 98 | 3 | astro-ph9803279_arXiv.txt |
9803 | astro-ph9803086_arXiv.txt | A multi-domain spectral method for computing very high precision 3-D stellar models is presented. The boundary of each domain is chosen in order to coincide with a physical discontinuity (e.g. the star's surface). In addition, a regularization procedure is introduced to deal with the infinite derivatives on the boundary that may appear in the density field when stiff equations of state are used. Consequently all the physical fields are smooth functions on each domain and the spectral method is absolutely free of any Gibbs phenomenon, which yields to a very high precision. The power of this method is demonstrated by direct comparison with analytical solutions such as MacLaurin spheroids and Roche ellipsoids. The relative numerical error reveals to be of the order of $10^{-10}$. This approach has been developed for the study of relativistic inspiralling binaries. It may be applied to a wider class of astrophysical problems such as the study of relativistic rotating stars too. | One of the most promising source of gravitational waves is the coalescence of inspiralling compact binaries. The recent development of interferometric gravitational waves detectors ({\sl e.g. GEO600, LIGO, TAMA} and {\sl VIRGO}) gives an important motivation for studying this problem. Such a study requires a relativistic formalism to derive the equations of motion and then an accurate and tricky method to solve the resulting system of partial differential equations. We have recently \cite{BonGM97a} proposed a relativistic formalism able to tackle the problem of co-rotating {\sl as well as} counter-rotating binaries system, these latter being more relevant from the astrophysical point of view. We present now a very accurate approach based on multi-domain spectral method that circumvents the Gibbs phenomenon to numerically solve this problem and which can be applied to a wide class of other astrophysical situations. Various astrophysical applications of spectral methods have been developed in our group (for a review, see \cite{BonGM97b}), including 3-D gravitational collapse of stellar core \cite{BonM93}, neutron star collapse into a black hole \cite{Gou91,GouH93,GouHG95,Nov98}, tidal disruption of a star near a massive black hole \cite{Mar96}, rapidly rotating neutron stars \cite{BonGSM93,SalBGH94a,SalBGH94b,HaeBS95}, magnetized neutron stars \cite{BocBGN95,HaeB96} and their resulting gravitational radiation \cite{BonG96}, spontaneous symmetry breaking of rapidly rotating neutron stars \cite{BonFG96,BonFG98}, proto-neutron stars evolution \cite{Gou97,GouHZ97,GouHZ98}. In computational fluid dynamics, spectral methods are known for their very high accuracy \cite{GotO77,CanHQZ88}; indeed for a ${\cal C}^\infty$ function, the numerical error decreases as $\exp(-N)$ ({\sl evanescent error}), where $N$ is the number of coefficients involved in the spectral expansion, or equivalently the number of grid points in the physical domain. This is much faster than the error decay of finite-difference methods, which behaves as $1/N^q$, with $q$ generally not larger than $3$. For this reason, spectral methods are particularly interesting for treating 3-D problems --- such as binary configurations --- a situation in which the number of grid points is still severely limited by the capability of present and next generation computers. Spectral methods lose much of their accuracy when non-smooth functions are treated because of the so-called {\em Gibbs phenomenon}. This phenomenon is well known from the most familiar spectral method, namely the theory of Fourier series: the Fourier coefficients $(c_n)$ of a function $f$ which is of class ${\cal C}^p$ but not ${\cal C}^{p+1}$ decrease as $1/n^p$ only. In particular, if the function has some discontinuity, its approximation by a Fourier series does not converge towards $f$ at the discontinuity point: there remains a gap which is of the order 10\%. The multi-domain spectral method described in this paper circumvents the Gibbs phenomenon. The basic idea is to divide the space into domains chosen so that the physical discontinuities are located onto the boundaries between the domains (Sect.~\ref{s:mapping}). The simplest example is the case of a perfect fluid star, where two domains may be distinguished: the interior and the exterior of the star. The boundary is then simply the surface of the star. The second ingredient of the technique is a mapping between the domains defined in this way and some simple mathematical domains, which are cross products of intervals: $[a_1,a_2] \times [b_1,b_2] \times [c_1,c_2]$. The spectral expansion is then performed with respect to functions of the coordinates spanning these intervals (Sect.~\ref{s:multi}). The method of resolution of a basic equation, namely the Poisson equation, is exposed in Sect.~\ref{s:resol_Poisson}. For stiff equations of state, the above procedure is not sufficient to ensure the smoothness of all the functions. Indeed, for a polytrope with an adiabatic index greater than 2, the density field has an infinite derivative on the surface of the star. We present in Sect.~\ref{s:regul} a method for regularizing the density and recover the spectral precision. The power of the multi-domain spectral method is illustrated in Sect.~\ref{s:illustr}, where comparisons are performed between numerical solutions obtained by an implementation of the method and analytical solutions (ellipsoidal configurations of incompressible fluids). Finally Sect.~\ref{s:conclu} concludes the article by discussing the great advantages of the multi-domain spectral method for dealing with relativistic binary neutron stars. | \label{s:conclu} We have presented a new numerical approach capable of handling the surface discontinuities of stellar configurations, provided these discontinuities are star-like, which covers a wide range of astrophysically relevant situations. When used along with spectral methods this adaptive-domain technique ensures that no Gibbs phenomenon can appear. This results in a very high precision (evanescent error), as demonstrated in Sect.~\ref{s:illustr} by comparison with exact analytical solutions. The relative error for 3-D configurations can reach $10^{-10}$ with a relatively small number of degrees of freedom ($N_r\times N_\theta \times N_\varphi = 37\times 19\times 18$ in each domain). Let us recall that very high accuracy is required for a lot of astrophysical problems such as numerical stability analysis. Among these problems let us mention the study of symmetry breaking of rapidly rotating stars and the determination of the orbital frequency of the last stable orbit of a neutron stars binary system. The multi-domain spectral method is particularly well adapted to the computation of relativistic binary neutron star system. Three sets of domains can be used in this problem (see Fig.~\ref{f:m-grille}): two sets of (three or more) domains centered on each star and a third set of (two or more) domains centered at the intersection between the rotation axis and the orbital plane. This latter set of domains which reaches spatial infinity is required to compute the gravitational field of relativistic configurations. When needed, the quantities computed on one of the three domain sets are evaluated at the collocation points of another set by means of the method presented in Sect.~\ref{s:code}. We are currently applying this numerical method to the computation of steady-state configurations of relativistic counter-rotating (i.e. irrotational with respect to an inertial frame) neutron star binaries, following the formulation developed in \cite{BonGM97a}. We will report on the astrophysical results in a forthcoming paper. An interesting byproduct of the present technical paper is the following one. In a previous work \cite{BonGSM93}, we had been able to demonstrate that the virial error is representative of the true error (measured by direct comparison with analytical solutions) only in the spherically symmetric case. We had inferred that this remains valid in the axisymmetric and 3-D cases. In the present work, we have confirmed this conjecture, thanks to the ability of the present method to treat incompressible fluids, for which 3-D analytical solutions are available. \begin{figure} \centerline{\epsfig{figure=ng.eps,height=4cm}} \caption[]{\label{f:m-grille} Representation of the numerical domains that we use to compute relativistic steady-state configurations of binary neutron stars systems. The external domain extends to spatial infinity in order to compute the exact gravitational potentials. Due to the symmetry of the problem, only the $z > 0$ part of space is taken into account. } \end{figure} | 98 | 3 | astro-ph9803086_arXiv.txt |
9803 | astro-ph9803092_arXiv.txt | X-ray emission from the distant lensing cluster CL2236-04 at $z$ = 0.552 was discovered by ASCA and ROSAT/HRI observations. If the spherical symmetric mass distribution model of the cluster is assumed, the lensing estimate of the cluster mass is a factor of two higher than that obtained from X-ray observations as reported for many distant clusters. However, the elliptical and clumpy lens model proposed by Kneib {\it et al.}(1993) is surprisingly consistent with the X-ray observations assuming that the X-ray emitting hot gas is isothermal and in a hydrostatic equilibrium state. The existence of the cooling flow in the central region of the cluster is indicated by the short central cooling time and the excess flux detected by ROSAT/HRI compared to the ASCA flux. However, it is shown that even if the AXJ2239-0429 has a cooling flow in the central region, the temperature measured by ASCA which is the mean emission-weighted cluster temperature in this case, should not be cooler than and different from the virial temperature of the cluster. Therefore, we conclude that the effect of the clumpiness and non-zero ellipticity in the mass distribution of the cluster are essential to explain the observed feature of the giant luminous arc, and there is no discrepancy between strong lensing and X-ray estimation of the mass of the cluster in this cluster. | It is one of the prime objectives in observational cosmology to determine mass distributions and total masses of clusters of galaxies at various redshifts. It provides a clue to constrain the unknown cosmological parameters:the species of dark matter, the nature of cosmological mass density fluctuations, the average density in the universe and the cosmological constant.\\ \indent Recently, gravitational lensing has become a powerful tool to measure the mass distribution in clusters. Strong lensing events such as multiple images and/or giant luminous arcs make it possible to model the central region of clusters. However, many clusters show a mass discrepancy with a total mass deduced from strong lensing observation being larger by a factor of 2--3 than that deduced from X-ray observations (Loeb \& Mao 1994, Miralda-Escud\'e \& Babul 1995, Kneib {\it et al.} 1995, Schindler {\it et al.} 1995, \& 1997, Wu \& Fang 1997, Ohta, Mitsuda, Fukazawa 1997). However, there are some clusters whose masses deduced in these two ways agree very well. In such cases, the multi-phase of the intra-cluster gas (Allen, Fabian \& Kneib 1996, Allen 1997) or asymmetric mass distribution (Piere {\it et al.} 1996, B\"ohringer {\it et al.} 1997) has been taken into account in the lens modeling. This indicates that careful studies for a large number of clusters are necessary before reaching the conclusions that the lens modeling gives higher mass estimation than that from X-rays. CL2236-04 was discovered by Melnick {\it et al.} (1993) with a giant luminous arc. The redshift of the arc was measured to be $z_{arc}$=1.116 (Melnick {\it et al.} 1993) confirming the gravitational lens picture very satisfactorily. The arc is rather straight and velocity structure across the arc is measured (Melnick {\it et al.} 1993). These nature give strict constraints on the lensing mass estimation. Further, CL2236-04 is one of the best and rare clusters at high redshift of $z>0.5$ for which comparison of the mass distribution deduced by two individual methods is possible. We have performed the first deep ASCA and ROSAT/HRI pointing observations toward CL2236-04. In this paper, we report the discovery and the properties of X-ray emission from CL2236-04 and the mass distribution deduced from the X-ray observation is compared to the lensing mass. | A new X-ray source in the direction of the distant lensing cluster, CL2236-04, is discovered by ASCA, named AXJ2239-0429. The source extent is resolved by the following ROSAT/HRI observation and it confirms that the X-ray source is the cluster of galaxies, CL2236-04. The mass distribution of the cluster derived by X-ray observations and by the observed nature of the giant luminous arc is compared. It is found that within the frame work of a spherical isothermal $\beta$ model, the required mass to explain the observed location of the giant luminous arc is larger than that expected from the X-ray observations as reported for many giant luminous arc clusters (Wu \& Fang 1997). An introduction of a large cosmological constant reduces the lensing mass but the total reduction is only 20-30\%. Therefore, the cosmological constant can not be the main solution for the discrepancy found in many clusters. Cooling flow can be one of the other possible solutions (Allen, Fabian \& Kneib 1996, Allen 1997). According to the cooling flow model, the hot gas in the cluster central region is in the multi-temperature phase. The temperature of the hottest phase gas which represents a real gravitational potential depth, obtained by the multi-phase model fitting becomes much higher than the temperature obtained by the single phase model fitting because the last one is the average temperature of the multi-temperature gas. Allen (1997) is claiming that the most of the discrepancy reported before can be explained by the existence of the cooling flow. There are several evidences of the existence of the cooling flow in the central region of AXJ2239-0429 as shown in Sec.2. It may be possible that the measured temperature of AXJ2239-0429 underestimates the cluster potential depth. If the ASCA temperature underestimates the cluster potential depth by a factor of two due to the cooling flow, the Einstein ring radius of the cluster for a source at $z=1.116$ becomes as large as the distance between the giant luminous arc and the cluster center. However, the straight nature of the arc and a lack of the bright counter image can not be explained by a spherically symmetric lens. Further, the perturbation by the gravitational potential of galaxy A is not negligible. The absolute luminosity of galaxy A in visual band is $3.3\times 10^{11}h_{50}^{-2}L_{\odot}$ (Melnick et al. 1993). This luminosity is translated into the the line of sight velocity dispersion of $\sim 370{\rm km/s}$ by using the Faber-Jackson law (Binney \& Tremaine 1987). If the potential depth is about a factor of two deeper than that inferred from the ASCA temperature, it is impossible to construct the model which is able to explain the location and the observed feature of the giant luminous arc due to the large perturbation of the galaxy A. Therefore, we conclude that even if the AXJ2239-0429 has a cooling flow in the central region, the temperature measured by ASCA which is the mean emission-weighted cluster temperature in this case, should not be cooler than and different from the virial temperature of the cluster. The most plausible solution is effect of the clumpiness of the gravitational potential due to galaxy A and non-zero ellipticity in the mass distribution of the cluster. KMG proposed the cluster mass distribution model by taking into account these effects which explain the observed features of the giant luminous arc. The ellipticity, position angle and velocity dispersion of the cluster predicted by KMG is surprisingly consistent with the results of the X-ray observations of the cluster as shown in the previous section. The velocity dispersion of galaxy A assumed by KMG is also consistent with that deduced by the Faber-Jackson law. Therefore, we conclude that the effect of the clumpiness of the gravitational potential due to galaxy A and non-zero ellipticity in the mass distribution of the cluster are essential to explain the observed feature of the giant luminous arc and the cluster mass distribution required by the strong lensing effect is totally consistent with the X-ray observations assuming that the X-ray emitting hot gas is isothermal and in a hydrostatic equilibrium state. | 98 | 3 | astro-ph9803092_arXiv.txt |
9803 | astro-ph9803167_arXiv.txt | We present Washington system $C,T_1$ color-magnitude diagrams of 13 star clusters and their surrounding fields which lie in the outer parts of the LMC disk ($r>4\arcdeg$), as well as a comparison inner cluster. The total area covered is large ($2/3^{\Box^\circ}$), allowing us to study the clusters and their fields individually and in the context of the entire galaxy. Ages are determined by means of the magnitude difference $\delta$T$_1$ between the giant branch clump and the turnoff, while metallicities are derived from the location of the giant and subgiant branches as compared to fiducial star clusters. This yields a unique dataset in which ages and metallicities for both a significant sample of clusters and their fields are determined homogeneously. We find that in most cases the stellar population of each star cluster is quite similar to that of the field where it is embedded, thus sharing its mean age and metallicity. The old population (t$\geq$10 Gyr) is detected in most fields as a small concentration of stars on the horizontal branch blueward and faintward of the prominent clump. Three particular fields present remarkable properties: (i) The so far unique cluster ESO121-SC03 at $\approx$9 Gyr has a surrounding field which shares the same properties, which in turn is also unique in the sense that such a dominant old field component is not present elsewhere, at least not significantly in the fields as yet studied. (ii) The field surrounding the far eastern intermediate age cluster OHSC\,37 is noteworthy in the sense that we do not detect any evidence of LMC stars: it is essentially a Galactic foreground field. We can thus detect the LMC field out to $>11\arcdeg$ (the deprojected distance of ESO121SC03), or $\sim 11$ kpc, but not to $13\arcdeg (\sim 13$ kpc), despite the presence of \cls at this distance. (iii) In the northern part of the LMC disk the fields of SL388 and SL509 present color-magnitude diagrams with a secondary clump $\approx$0.45 mag fainter than the dominant intermediate age clump, suggesting a stellar population component located behind the LMC disk at a distance comparable to that of the SMC. Possibly we are witnessing a depth effect in the LMC, and the size of the corresponding structure is comparable to the size of a dwarf galaxy. The unusual spatial location of the cluster OHSC37 and the anomalous properties of the SL\,388 and SL\,509 fields might be explained as debris from previous LMC interactions with the Galaxy and/or the SMC. The mean metallicity derived for the intermediate age outer disk clusters is $<$[Fe/H]$>$$=-0.7$ and for their surrounding fields $<$[Fe/H]$>$$=-0.6$. These values are significantly lower than found by Olszewski {\it et al.} (1991, AJ, 101, 515) for a sample of clusters of similar age, but are in good agreement with several recent studies. A few clusters stand out in the age--metallicity relation in the sense that they are intermediate age clusters at relatively low \met ([Fe/H]$\approx -1$). | Unveiling the star formation history and chemical enrichment of galaxies is critical for understanding how they form and evolve. In this respect, Local Group galaxies play a fundamental r\^ole (Hodge 1989). The proximity of the Large Magellanic Cloud allows one to probe its stellar population properties with different techniques, namely photometry and spectroscopy of individual stars in clusters and the field, and integrated methods in the case of star clusters. Such studies for nearby galaxies are important to better understand very distant galaxies, whose stellar populations can only be probed by means of integrated properties. Concerning field color-magnitude diagram (CMD) studies, ground-based observations have allowed the accurate study of the brighter evolutionary sequences, sampling relatively large fields throughout the LMC bar and disk (e.g. Butcher 1977, Hardy {\it et al.} 1984, Bertelli {\it et al.} 1992, Westerlund, Linde \& Lyng\aa\ 1995, Vallenari {\it et al.} 1996). A prominent intermediate age (1-3 Gyr) stellar population is universally present in these fields, together with varying amounts of young blue main sequence (MS) stars. Hubble Space Telescope (HST) observations are limited to a small viewing area, but in turn allow deeper photometry, to well below the old MS turnoff. Several such fields have been studied with the V and I bands. One, at $\approx$4$^o$ north of the bar (Gallagher {\it et al.} 1996) and another near the SE end of the bar (Elson, Gilmore \& Santiago 1997) show evidence for the major star-forming event to have ocurred $\approx$2 Gyr ago, in agreement with the ground-based studies. More recent HST studies (Holtzman {\it et al.} 1997, Geha {\it et al.} 1998) have investigated three LMC fields at 3$^o$ to 4$^o$ from the bar center, one located in the north-east and two in the north-west. Surprisingly, they are finding many more faint MS stars than expected and suggest that there has been more star formation in the past than previously believed. Their models, assuming a standard IMF slope, suggest that fully one half of the stars in these fields were formed more than 4 Gyr ago. Although the latter studies refer to their fields as ``outer", we point out that in the present work we deal with genuine outer disk fields, well beyond any of these HST studies. Integrated photometry of large star cluster samples of all ages have shown differences in the spatial distribution of age groups both in the bar region (Bica, Clari\'a \& Dottori 1992) and the entire LMC (Bica {\it et al.} 1996). Differences in the spatial distribution among young groups have provided insight on the formation process and subsequent dynamical evolution of star cluster generations (Dottori {\it et al.} 1996). This integrated photometry cluster sample has been compared with integrated star cluster color models and has provided constraints on the cluster formation history (Girardi \& Bica 1993, Girardi {\it et al.} 1995). CMDs of LMC star clusters have also revealed a large intermediate age population (1-3 Gyr), which is separated by a pronounced age gap from the old stellar population as denoted by a few genuine globular clusters (see Da Costa 1991, Suntzeff {\it et al.} 1992, Olszewski, Suntzeff \& Mateo 1996 for reviews). Recently, CMDs in the Washington system of a sample of candidate old clusters selected from the Bica {\it et al.} (1996) and Olszewski {\it et al.} (1991) studies revealed them to instead be of intermediate age (Geisler {\it et al.} 1997, hereafter Paper I). This study increased considerably the known sample of 1-3 Gyr old clusters with accurate age determinations, and reinforced the conclusion that a major formation epoch was preceded by a quiescent period of many Gyr, or that cluster dissipation has been more effective than generally believed (e.g. Olszewski 1993). The objective of the present paper is to compare the properties of outer LMC clusters with those of their surrounding fields by using the same observational technique, and to infer the age-metallicity relation, and whether it depends on the spatial distribution throughout the LMC disk. In order to achieve this we employ Washington system C, T$_1$ bands and construct CMDs using the data of Paper I. Ages are inferred from the difference $\delta$T$_1$ between the giant branch clump and the turnoff (see also Paper I), and from the occurrence of particular stellar evolutionary sequences in the CMDs. The giant and subgiant branches allow one to derive metallicities using a technique analogous to that of Da Costa and Armandroff (1990) for VI photometry. However, our combination of Washington system filters is three times more \met sensitive than the VI system (Geisler \& Sarajedini 1996, 1998), allowing us to obtain accurate \mets for both the clusters and their fields. The cluster/field sample and the observations are described in Section 2. The cluster and field CMDs are described in Section 3. Ages and metallicities are derived in Section 4. In Section 5 we discuss the chemical enrichment of the outer disk, and in Section 6 the presence of dual clumps in two fields is noted and the possibility of a depth effect in this portion of the LMC disk is discussed. Finally, the conclusions of this work are given in Section 7. | By using Washington system color-magnitude diagrams of a large sample of IACs and their surrounding fields in the LMC outer disk we were able to derive ages and metallicities. The cluster stellar population is a major component of the field where it is embedded, thus sharing its mean age and metallicity properties, except in three cases where the clusters are $\approx$0.3 dex more metal poor than the field, suggesting that the chemical enrichment was not globally homogeneous in the LMC. Our data is consistent with a scenario in which local star formation events generated both the clusters and a significant part of their surrounding stellar fields. Thus during the last 1-3 Gyr the dynamical evolution of the disk has not significantly taken them apart, which provides information on the diffusion time-scale and mixing of stellar generations in the disk. The old population (t$\geq$10 Gyr) is detected in most fields as a small concentration of stars on the horizontal branch blueward and faintward of the prominent clump. The unique cluster ESO121-SC03 at $\approx$9 Gyr has a surrounding field which shares the same stellar \pop. No other field so far studied is dominated by such an old population. One possibility would be that a field and cluster population coupling might last that long. Alternatively, we might be dealing with a building block recently accreted by the LMC in the form of a dwarf galaxy. One IAC cluster (OHSC37) is so far from the LMC body that no surrounding LMC field is detected. The present observations suggest that the LMC stellar disk extends out to between $\sim 11-13$ kpc in deprojected radius. In the northern part of the LMC outer disk the fields of SL388 and SL509 present evidence of a depth effect with a secondary component located behind the LMC disk at a distance comparable to that of the SMC. A background layer of stars in the LMC was possibly detected, and its size is at least $\approx$1 kpc, comparable to that of a dwarf galaxy. The peculiar location of the cluster OHSC37 and the depth effect in the SL\,388 and SL\,509 fields might be explained as debris from previous interactions of the LMC with the Galaxy and/or the SMC. The average metallicity derived for the present outer disk IACs is $<$[Fe/H]$>$$=-0.71\pm0.17$ and for their surrounding fields $<$[Fe/H]$>$$=-0.61\pm0.11$. A few clusters stand out in the age--metallicity relation in the sense that they are intermediate age clusters at [Fe/H]$\approx -1$. The outer LMC disk (clusters and fields), at least of this age range (1-3 Gyrs) seem to be more metal-poor than previously generally regarded. The authors would like to thank Cristina Torres for providing data critical to the \met calibration in advance of publication. E. Geisler, as always, was an inspiration. This research is supported in part by NASA through grant No. GO-06810.01-95A (to DG) from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. This work was partially supported by the Brazilian institutions CNPq and FINEP, the Argentine institutions CONICET and CONICOR, and the VITAE and Antorcha foundations. | 98 | 3 | astro-ph9803167_arXiv.txt |
9803 | astro-ph9803021_arXiv.txt | The {\it COBE} FIRAS data contain foreground emission from interplanetary, Galactic interstellar dust and extragalactic background emission. We use three different methods to separate the various emission components, and derive the spectrum of the extragalactic Far InfraRed Background (FIRB). Each method relies on a different set of assumptions, which affect the FIRB spectrum in different ways. Despite this, the FIRB spectra derived by these different methods are remarkably similar. The average spectrum that we derive in the $\nu = 5 - 80~$cm$^{-1}$ (2000 -125 \um) frequency interval is: $I(\nu) = (1.3\pm0.4)\times10^{-5}\ (\nu /\nu_0)^{0.64\pm.12}\ P_{\nu}(18.5\pm 1.2~$K), where $\nu_0=100~$cm$^{-1}$ ($\lambda_0=100\ \um$) and $P$ is the Planck function. The derived FIRB spectrum is consistent with the DIRBE 140 and 240 \um\ detections. The total intensity received in the 5 - 80~cm$^{-1}$ frequency interval is 14 nW m$^{-2}$ sr$^{-1}$, and comprises about 20\% of the total intensity expected from the energy release from nucleosynthesis throughout the history of the universe. | The FIRAS (Far~Infrared~Absolute~Spectrophotometer) instrument aboard the Cosmic Background Explorer ({\it COBE}) satellite (Boggess \etal\ 1992 and references therein; Mather, Fixsen, \& Shafer 1993) was designed to measure the spectrum of the Cosmic Microwave Background (CMB), the Galaxy, and the Far InfraRed Background (FIRB). The FIRAS instrument and calibration are discussed by Mather, Fixsen \& Shafer (1993) and Fixsen \etal\ (1994b). The FIRAS observations of the CMB are discussed by Mather \etal\ (1990, 1994), and Fixsen \etal\ (1994a, 1996, 1997a). The FIRAS observations of the Galaxy are discussed in Wright \etal\ (1991), Bennett \etal\ (1994), and Reach \etal\ (1995). Puget \etal\ (1996) used the Pass 3 FIRAS observations to make a tentative background determination simular to one of the methods used here. At frequencies below $\sim$ 20 cm$^{-1}$ (500 \um), the FIRB is overwhelmed by the CMB. However, the CMB can be subtracted from the data since its emission is spatially uniform, and it has a Planck spectrum. The dipole anisotropy of the CMB must be included but the other anisotropy is not large enough to be a problem. At frequencies above $\sim$ 100 cm$^{-1}$ (100 \um),beyond the FIRAS range, the observed spectrum is dominated by the zodiacal emission, which is much more complex, since its intensity and spectrum are spatially varying. In between these frequencies, the observed spectrum is dominated by Galactic emission, which can be identified by its distinct (from zodiacal) spectrum, and its spatial variation over the sky. After the subtraction of these foreground emissions, any isotropic residual emission may still contain a uniform component from the solar system or the Galaxy, which must be estimated and separated from the FIRB. In this paper we examine three methods, which give a consistent result for the FIRB. | The three methods just described have different assumptions, weaknesses and strengths. The first method assumes that the Galactic spectrum is fixed, and that only its intensity varies with position. Temperature variations are observed to occur on various scales (Reach et al. 1995; FIRAS Exp Sup 1997; Lagache et al. 1998). Based on the DIRBE analysis, we expect these variations to have a only a small effect on our results. The DIRBE high latitude data at 240 \um\ were analyzed using both a single- and a variable-spectrum model to describe the Galactic foreground emission (Arendt et al. 1998), with no significant difference in the derived background. The second method does not assume a fixed Galactic spectrum, but almost all of the power in the fit is in a single component indicating that there is little systematic correlation of temperature with intensity of the H~I 21 cm line. This method solves directly for the dust emission from the dominant neutral and ionized phases of the ISM. It therefore avoids a potential problem of the third method (which relies on the DIRBE analysis) using backgrounds derived from the IR-to-H~I correlation only. Two problems with this method are, that at the FIRAS resolution, the correlation between the IR and the H~I line intensity is non-linear and that it ignores the molecular regions. The first was treated by using a quadratic fit to the relation. This choice is not unique, but it is supported by observations of the IR to H~I correlation in the Galaxy (Dall'Oglio et al. 1985; Reach, Koo, \& Heiles 1994). Molecular clouds are less abundant at the high Galactic latitudes that are used here. The third method relies on the accuracy of the DIRBE determination of the FIRB. The main uncertainty in the background determination are those associated with the determination of the foreground emission, and they are discussed in detail by Hauser et al. (1998). The use of a DIRBE template for the FIRAS background determination introduces another uncertainty, namely the consistency between the FIRAS and DIRBE calibration. Fixsen et al. (1997b) show that the most significant difference in the calibration is in the 240 \um\ DIRBE band, but using the FIRAS instead of the DIRBE calibration for that band introduces a very small change in the background, from a value of 13.6 to 12.7 nW m$^{-2}$ sr$^{-1}$ (Hauser et al. 1998). Overall, the three methods yield a consistent spectrum for the FIRB, an encouraging result considering the substantial differences in the three approaches (see fig 4a). The weaknesses of each method are compensated by the other methods. The color method assumes a single spectrum, but the other methods allow for color variation in the Galactic foreground. The line emission method uses a quadratic fit over 25\% of the sky, but the other methods use linear fits which are insensitive to the fraction of sky. The DIRBE method only uses the correlation between H~I and the foreground in small regions, but the line emission method includes ionized gas and the color method makes no assumptions about the gas. In all three FIRB spectra the background peaks at $\sim 50~$cm$^{-1}$ ($\sim 200$ \um), and exhibits a definite turnover at the higher frequencies. The higher frequencies are more affected by both noise and systematic effects. The average of the three spectra can be fit by: \begin{equation} I_\nu=(1.3\pm 0.4)\times 10^{-5}(\nu/\nu_\circ)^{.64\pm.12}\ P_\nu(18.5\pm1.2~K), \end{equation} in the 5 to 80 cm$^{-1}$ frequency range ($\lambda$ between 125 and 2000 \um) where $\nu_\circ=100~$cm$^{-1}$, and $P_\nu$ is the familiar Planck function. The uncertainties are highly correlated, with correlations of .98 for the intensity and index, $-$.99 for the intensity and temperature. and $-$.95 for the index and temperature. Figure 4b compares the analytical fit to the FIRB derived here, to the tentative background derived by Puget et al. (1996). The spectra differ significantly at $\nu \sim 35 - 60~$cm$^{-1}$, probably the result of the difference in subtraction of dust emission related to H$^+$. The crosses in the Figure represent the nominal DIRBE detections at 140 and 240 \um, while the diamonds represent the DIRBE detections using the FIRAS calibration. While the effect of the FIRAS calibration is larger at 140 \um, the uncertainty in the calibration is larger at this wavelength. So, the FIRB derived here is consistent with that derived by the DIRBE, within the uncertainty of the DIRBE$-$FIRAS calibration. The range of values for the FIRB at 100 \um\ represent the upper limit derived by Kashlinsky, Mather, \& Odenwald (1996) from a fluctuation analysis of the 100 \um\ DIRBE maps and the lower limit is derived by Dwek et al. (1998) from the 140 and 240 \um\ DIRBE detections. Dwek \etal\ (1998) show that the uniform DIRBE residuals cannot be produced by any local (solar system or Galactic) emission sources. Hence, the uniform residual derived here is most likely of extragalactic origin. The total flux received in this wavelength region is 14 nW m$^{-2}$ sr$^{-1}$, or about 20\% of the total expected flux of about 70 nW m$^{-2}$sr$^{-1}$ associated with the production of metals throughout the history of the universe in some models (Dwek \etal\ 1998). | 98 | 3 | astro-ph9803021_arXiv.txt |
9803 | astro-ph9803217_arXiv.txt | We study numerically the evolution of an adiabatic relativistic fireball expanding into a cold uniform medium. We follow the stages of initial free expansion and acceleration, coasting and then deceleration and slowing down to a non-relativistic velocity. We compare the numerical results with simplified analytical estimates. We show that the relativistic self similar Blandford-McKee solution describes well the relativistic deceleration epoch. It is an excellent approximation throughout the relativistic deceleration stage, down to $\gamma \sim 5$, and a reasonable approximation even down to $\gamma \sim 2$ though the solution is rigorous only for $\gamma \gg 1$. We examine the transition into the Blandford-McKee solution, and the transition from the solution to the non-relativistic self similar Sedov-Taylor solution. These simulations demonstrate the attractive nature of the Blandford-McKee solution and its stability to radial perturbations. | Sedov (1946), Taylor (1950) and Von Neumann (1947) discovered, in the forties, a self similar solution of the strong explosion problem, in which a large amount of energy is released on a short time scale in a small volume. This solution is known today as the ``Sedov-Taylor'' self similar solution. It describes a shock wave propagating into a uniform density surrounding. The shock wave and the matter behind it decelerate as more and more mass is collected. This solution describes well the adiabatic stage of a supernova remnant evolution. Blandford and McKee (1977) have later established a self similar solution describing the extreme relativistic version of the strong explosion problem. In this solution the Lorentz factor of the shock and the fluid behind it is much larger than unity. Such high Lorentz factors arise if the rest mass contained within the region where the energy $E$ was released is much smaller than $E$. In other words when the region containing the energy is ``radiation dominated'' rather than matter dominated. Such a region was later termed a ``fireball''. Cavallo and Rees (1978) have considered the physical processes relevant in a radiation dominated fireball as a model for gamma-ray bursts (GRBs). Goodman (1986) and Paczy\'nski (1986) have then considered the evolution of such a fireball. They have shown that a initially the radiation-pair plasma in a purely radiative fireball behaves like a fluid and it expands and accelerated under its own pressure until the local temperature drops to $\sim 20$keV, when the last pairs annihilate and the fireball becomes optically thin. Later Shemi and Piran (1990) considered matter contaminated fireballs. They have shown that, quite generally, all the initial energy will be transfered to the baryons in such fireballs whose final outcome is a shell of relativistic freely expanding baryons. Piran, Shemi and Narayan (1993) and M\'esaz\'aros Laguna and Rees (1993) have later carried these calculations in greater details. These works show that an initially homogeneous fireball will first accelerate while expanding and then coast freely as all its internal energy was transformed to kinetic energy. The surrounding matter will eventually influence the fireball after enough external matter has been collected and most of the energy has been transfered from the shell to the ISM. If the surrounding matter is diluted enough, then this will take place only after the initial free acceleration phase. This influence was considered by M\'esaz\'aros and Rees (1992), and Katz (1993), who suggested that the GRB is produced during this stage. The detailed shock evolution was later studies by Sari and Piran (1995). It is only after these stages, that the fireball have given the ISM most of its energy and then the self similar deceleration solution of Blandford-McKee applies. When the shock decelerates enough so that it is no longer relativistic, it is described by the Sedov-Taylor solution. Today it is widely accepted that GRBs involve relativistic expanding matter of this kind. While the GRB itself is produced, most likely, via internal shocks (Narayan, Paczy\'nski \& Piran, 1992; Rees \& M\'esaz\'aros, 1994; Sari \& Piran, 1997). The observed GRB afterglow corresponds, on the other hand to the slowing down of this relativistic flow. This has led to an increasing interest in the fireball solution and in its various regimes. In this paper we study the evolution of a homogeneous fireball focusing on its interaction with the surrounding matter. We do not consider here internal shocks, which arise due to interaction within the relativistic flow and require nonuniform velocity. We have developed a spherically symmetric relativistic Lagrangian code based on a second order Gudnov method with an exact Riemann solver to solve the ultra-relativistic hydrodynamics problem. With this code, it is possible to track the full hydrodynamical evolution of the fireball within a single computation, from its initial ``radiation dominated'' stage at rest through its acceleration, coasting, and shock formation, up to the relativistic deceleration and finally to the Newtonian deceleration. With typical parameters this computation spans more than eight orders of magnitudes in the size of the fireball and more than twenty orders of magnitudes in its density. We describe these computations here. We show that, quite generically, the solution converges during the relativistic deceleration phase to the Blandford-McKee solution and then it transforms to the Sedov-Taylor solution. Even though the attractive nature of the Blandford-McKee solution suggests that it is stable we explore explicitly the stability of this solution and we show that it is stable to radial perturbations. We review, first, in section \ref{sec_Analytic} the current analytic understanding of the fireball evolution thorough the following stages: (i) free acceleration and coasting, (ii) energy transfer (iii) relativistic self similar solution and (iv) Newtonian Sedov-Taylor solution. We discuss the numerical results in section \ref{numerical}. In section \ref{perturbation} we examine the evolution of perturbations to the Blandford-McKee solution. We discuss the implications of these results in section \ref{conclusions}. | 98 | 3 | astro-ph9803217_arXiv.txt |
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9803 | astro-ph9803117_arXiv.txt | We investigate the formation of virialized halos due to the gravitational collapse of collisionless matter using high-resolution N-body simulations. A variety of formation scenarios are studied, ranging from hierarchical clustering to monolithic radial collapse. The goal of these experiments was to study departures from the universal density profiles recently found to arise in cosmological settings. However, we found that even for models which exhibit quite a different formation history, the density and velocity dispersion profiles of the virialized halos are strikingly similar. Power law density profiles do not result even in models with initial power law profiles and no initial substructure or non-radial motions. Such initial conditions give rise to a radial orbit instability which leads to curved velocity dispersion and density profiles. The shapes of the density profiles in all our models are well parameterized by the profiles of halos formed in a generic cosmological setting. Our results show that the universality of dark halo density profiles does not depend crucially on hierarchical merging as has been suggested recently in the literature. Rather it arises because apparently different collapse histories produces a near universal angular momentum distribution among the halo particles. We conclude that the density and velocity dispersion profiles of virialized halos in an expanding universe are robust outcomes of gravitational collapse, nearly independent of the initial conditions and the formation history. | The hierarchical collapse of dark matter into virialized halos is likely to have played a key role in the formation of galaxies and clusters of galaxies. Several aspects of this process have been studied in recent years, in particular the role of initial conditions and of ongoing mergers in shaping the final structure of the dark matter halos. However, it remains an open question whether violent relaxation erases all information about the initial conditions and the formation history of halos. If it does not, then the density and velocity profiles of dark matter halos may bear imprints of the initial power spectrum as well as the cosmological density parameter $\Omega$ and the cosmological constant $\Lambda$. Such a dependence between the initial power spectrum and the density profiles of virialized objects was pointed out by Hoffman \& Shaham (1985) and Hoffman (1988). Their conclusions were based on the secondary infall model of Gunn \& Gott (1972), and the self-similar solution of Fillmore \& Goldreich (1984) and Bertschinger (1985). These analytic calculations, however, were based on simplifying assumptions (\eg spherical symmetry) and, therefore, could not consider all aspects of gravitational collapse. N-body simulations, which begin with generic initial conditions, provide an attractive alternative. Though the effects of gas dynamics are neglected, they are likely to represent a realistic description of the formation of galaxy clusters and the outer parts of galactic halos. In recent years, multi-mass simulation techniques have been developed which allow one to investigate the formation of individual halos with high numerical resolution. Navarro, Frenk \& White (1996, 1997) (hereafter NFW) have investigated the structure of dark matter halos which form in a cold dark matter (CDM) universe. They found that the density profiles of halos do not follow a power law, but tend to have a slope $\alpha = d \ln \varrho/d \ln r$ with $\alpha=-1$ near the halo center and $\alpha = -3$ at large radii. Over more the four orders of magnitude in mass, the density profiles follow a universal law, which can be parameterized by \begin{equation} \frac{\varrho(r)}{\varrho_b} = \frac{\delta_n}{\frac{r}{a_n} (1 + \frac{r}{a_n})^2}. \label{nfw} \end{equation} The two parameters are the scale radius $a_n$ which defines the scale where the profile shape changes from slope $\alpha >-2$ to $\alpha <-2$, and the characteristic overdensity $\delta_n$. They are related because the mean overdensity enclosed within the virial radius $r_{vir}$ is $180$. Equation~(\ref{nfw}) differs slightly in its asymptotic behavior at large radii from the profile which was proposed by Hernquist (1990) to describe the mass profiles of elliptical galaxies. This profile, which goes like $r^{-4}$ at large radii, was used by Dubinsky \& Carlberg (1991) to fit the density distribution of halos which were formed in their CDM-like simulation. Cole \& Lacey (1996), Tormen, Bouchet \& White (1996), Navarro, Frenk \& White (1997), Huss, Jain \& Steinmetz 1997 (hereafter HJS) and Moore \etal (1997), among others, have extended the original results of NFW to other initial spectra, to other cosmologies and to higher spatial and mass resolution. Most of the above studies show that halos which form in a variety of cosmological models are well described by equation~(\ref{nfw}). The scale radius and the central overdensity appear to be related to the formation time of the halo (Navarro \etal 1997). These results suggest that the density profile found by NFW is quite generic for any scenario in which structures form due to hierarchical clustering. The power spectrum and cosmological parameters only enter by determining the typical formation epoch of a halo of a given mass and thereby the dependence of the characteristic radius $a$ on the total mass of the halo. All these studies were based on initial particle distributions drawn from a Gaussian random field. Since cold dark matter scenarios have initial perturbations on all scales, the formation of halos generically proceeds by hierarchical merging of dark matter clumps. In this paper we instead use a set of artificial initial conditions, gradually increasing in complexity, to isolate the physical effects that determine the final properties of the halo. We control the degree of substructure by varying the velocity dispersion of the particles. The merging history of the halos varies between the two extremes of monolithic radial collapse and hierarchical merging. | We have simulated the collapse of dark matter halos with varying initial conditions. As shown in figure 1 the models are designed to vary the amount of sub-structure and merging in the collapsing halo. We found that even with a smooth initial particle distribution and minimal merging, the virialized density and velocity dispersion profiles of the halos are very similar to the halos formed in a cold dark matter cosmology. The density profile has logarithmic slope $-2$ at a characteristic radius and a slope shallower in the inner regions and steeper in the outer regions. It is well fit by the form proposed by NFW. The only way we could get a power law density profile was to turn off the non-radial components of the gravitational force completely. Since this artificial condition can only be satisfied in a simulation, we conclude that dark matter halos resulting from gravitational collapse do not have power law profiles. Instead the density and velocity dispersion profiles have a nearly universal characteristic shape, independent of the initial conditions or the formation history. Our results suggest that the key to universal density profiles does not lie in the hierarchical merging history of cold dark matter halos, as recently argued by Syer \& White (1996). It is a more generic feature of gravitational collapse, which can be derived from the near universal behavior of the velocity field of the halo particles. The velocity dispersion profile grows with time in nearly the same way for halos in all the models. The mechanism by which angular momentum is acquired can vary, e.g. model II has little merging and instead shows a strong bar instability, but the resulting velocity dispersion profiles have similar shape. NFW used hierarchical merging as embodied in the Press-Schechter scenario to model the variation of the scale radius $a_n$ with halo mass and the cosmological model. To the extent that their results are empirically verified, they provide a useful guide to the variations of halo profiles. But our results suggest that the merging history is not a fundamental determinant of halo profiles, since the emergence of the characteristic scale radius occurs even in models with negligible merging. It would be interesting to pursue this issue further with detailed dynamical studies. | 98 | 3 | astro-ph9803117_arXiv.txt |
9803 | astro-ph9803267_arXiv.txt | We report the discovery of a quasi-periodic oscillation (QPO) at 61.0$\pm$1.7 Hz with the Rossi X-Ray Timing Explorer in the low-mass X-ray binary and persistently bright atoll source GX 13+1 (4U 1811--17). The QPO had an rms amplitude of 1.7$\pm$0.2\% (2--13.0 keV) and a FWHM of 15.9$\pm$4.2 Hz. Its frequency increased with count rate and its amplitude increased with photon energy. In addition a peaked noise component was found with a cut-off frequency around 2 Hz, a power law index of around --4, and an rms amplitude of $\sim$1.8\%, probably the well known atoll source high frequency noise. It was only found when the QPO was detected. Very low frequency noise was present with a power law index of $\sim$1, and an rms amplitude of $\sim$4\%. A second observation showed similar variability components. In the X-ray color-color diagram the source did not trace out the usual banana branch, but showed a two branched structure. This is the first detection of a QPO in one of the four persistently bright atoll sources in the galactic bulge. We argue that the QPO properties indicate that it is the same phenomenon as the horizontal branch oscillations (HBO) in Z sources. That HBO might turn up in the persistently bright atoll sources was previously suggested on the basis of the magnetospheric beat frequency model for HBO. We discuss the properties of the new phenomenon within the framework of this model. | Based on their correlated X-ray timing and spectral behavior the brightest low-mass X-ray binaries (LMXBs) can be divided into two groups; the atoll sources and Z sources (Hasinger \& van der Klis \markcite{hk89}1989, hereafter HK89; van der Klis \markcite{kl95}1995). Atoll sources show two states: the island and the banana state, after the tracks they produce in an X-ray color-color diagram (CD). Z sources on the other hand trace out a Z-like shape in a CD, with usually three branches: the horizontal, the normal, and the flaring branch. The power spectra of atoll sources can be described by two noise components plus, sometimes, a Lorentzian component to describe quasi-periodic oscillations (QPOs). The first noise component, the very low frequency noise (VLFN) has a power law shape $P\propto\nu^{-\alpha}$, with $1\leq\alpha\leq1.5$. The other, the high frequency noise (HFN), can be described by a power law with an exponential cut-off $P\propto\nu^{-\alpha} e^{-\nu/\nu_{cut}}$, usually with $0\leq\alpha\leq0.8$ and $0.3\leq\nu_{cut}\leq25$ Hz. The HFN sometimes has a local maximum (``peaked noise'') around 10-20 Hz (see van der Klis 1995). Yoshida et al. (1993) found peaked noise around 2 Hz. Several broad QPO(-like) peaks were found with the Rossi X-ray Timing Explorer (RXTE) around 20 Hz (Strohmayer et al. \markcite{st96}1996; Ford et al. \markcite{fo97}1997; Yu et al. \markcite{yu97}1997; Wijnands \& van der Klis \markcite{wikl97}1997), and one at 67 Hz (simultaneously with one at 20 Hz, see Wijnands et al. \markcite{wi97b}1997b;). In addition to QPOs below 100 Hz, QPOs between 300 and 1200 Hz, the so-called kHz QPOs, have been found. The power spectra of Z sources show three broad noise components: VLFN with $1.5\leq\alpha\leq2$, HFN with $\alpha\sim0$ and $30\leq\nu_{cut}\leq100$ Hz, and low frequency noise (LFN), which has the same functional shape as the HFN with $\alpha\sim0$ and $2\leq\nu_{cut}\leq20$ Hz. Note that despite having the same name, HFN in Z sources is not the same phenomenon as HFN in atoll sources. Z sources show three types of QPOs: the normal/flaring branch QPO (N/FBO) with centroid frequencies from 6 to 20 Hz, the horizontal branch QPO (HBO) from 15 to 60 Hz, and the kHz QPOs in the same range as observed in atoll sources. (For an extensive review on the power spectra of atoll and Z sources we refer to van der Klis \markcite{kl95}1995. For kHz QPOs we refer to van der Klis \markcite{kl97}1997.) GX 13+1 has been classified as an atoll source, although of all atoll sources it shows properties which are closest to that seen in the Z sources (HK89). Moreover, Schulz, Hasinger \& Tr\"umper\markcite{sc89} (1989) put GX 13+1 among the luminous sources that have been classified as Z sources. Together with GX 3+1, GX 9+1, and GX 9+9, GX 13+1 forms the subclass of the persistently bright atoll sources. In the CD they have only been seen to trace out banana branches and their power spectra can be described by relatively strong ($\sim$3.5\% rms) VLFN and weak ($\sim$2.5\% rms) HFN, as compared to other atoll sources. No QPOs have been found before in these sources, neither at frequencies below 100 Hz nor at kHz frequencies (Wijnands, van der Klis \& van Paradijs\markcite{wi97a} 1997a; Strohmayer et al.\markcite{st97} 1997). For GX 13+1, so far no HFN has been observed. Its banana branch resembled a more or less straight strip in the CD, whereas the other three sources showed more curved banana branches (HK89). Stella, White \& Taylor \markcite{st85}(1985) have reported bimodal behavior of GX 13+1 in the hardness-intensity diagram (HID). In one state the spectral hardness was correlated with count rate, while in the other it was anticorrelated. The transition between the two states occurred within one hour. The main difference between atoll and Z sources is the mass accretion rate, $\dot{M}$. Atoll sources accrete at $\dot{M}\la0.5\dot{M}_{Edd}$, whereas Z sources accrete at near Eddington rates. It was proposed by HK89 that a second difference lies in the strength of the magnetic field strength $B$ of the accreting neutron star. Recent spectral modeling by Psaltis \& Lamb \markcite{ps96}(1996) suggests that indeed atoll sources have $B<10^9$ Gauss and Z sources $B\sim10^9$--$10^{10}$ Gauss. The bright atoll sources are found to have a higher inferred $B$ than the low luminosity atoll sources, making them the best candidates to show Z source HBO. In this Letter we report the discovery of a 57--69 Hz QPO and of a two branched structure in the CD and HID of GX 13+1. We suggest that the QPO is the same phenomenon as Z source HBO. | We have discovered a QPO in GX 13+1 between 57 and 65 Hz and around 69 Hz. It is the first time a QPO has been found in a persistently bright atoll source. The rms amplitude of the QPO increased with photon energy. Together with the QPO, a peaked noise component was found, probably atoll source HFN, with a cut-off frequency of $\sim$2 Hz. The HFN was only detected when the QPO was present. In the October data the QPO was only found in the upper branch, i.e. at high count rates. On the upper branch the QPO frequency increased with count rate from $\sim$57 Hz to $\sim$65 Hz. In the November data the QPO could only be detected at high count rates. Assuming that $\dot{M}$ and count rate were positively correlated, the QPO frequency increased with $\dot{M}$, at least on the upper branch in the October observation. However, during the November observation we found the QPO at $\sim$69 Hz; the mean count rate was then lower than during the October observation of the $\sim$65 Hz QPO. The pattern traced out by GX 13+1 in the CD and HID in our observations does not resemble the patterns traced out by other atoll sources. Atoll sources trace out islands and/or a banana branch. During EXOSAT observations GX 13+1 traced out a banana branch (HK89) which is quite different from what we observed with RXTE. We also do not find evidence for the bimodal behavior found by Stella et al. \markcite{st85}(1985) in the sense that both our branches in the HID show a positive correlation between hard color and count rate. The relatively sharp turn in the CD and HID may be related to one of the vertices in the patterns traced out by Z sources. Comparison with EXOSAT observations, using our RXTE energy spectra folded with the EXOSAT response matrix, shows that the source was $\sim$30\% brighter during the RXTE observations. (Note that in the RXTE ASM lightcurve the source shows intrinsic variations of $\sim$50\%; during our observations the ASM count rate was near average.) This might explain why the two branched structure has not been seen before in GX 13+1. In any case EXOSAT was not sensitive enough to have detected the QPO reported in this Letter. Recently Stella \& Vietri \markcite{st98}(1998) proposed that at least some of the $<100$ Hz QPOs observed in atoll sources are due to Lense-Thirring precession of the inner part of the accretion disk. In order to test this model one needs the frequencies of the simultaneously observed kHz QPOs. Since no kHz QPOs have been observed in GX 13+1 it is not possible to test this model. The upper limits on kHz QPOs are comparable with those found in the other members of the subclass (Wijnands et al. \markcite{wi97a}1997a; Strohmayer et al. \markcite{st97}1997) The QPO properties (frequency, dependence on $\dot{M}$, energy spectrum, and the simultaneous presence of a band limited noise component [the HFN]) are all similar to those found for HBO in Z sources. On the basis of the above described similarities we suggest that the QPO we found is the same phenomenon as the HBO, and that the HFN component is related to the QPO in a similar way as Z source LFN to HBO. The identification of atoll source HFN with Z source LFN was previously proposed by van der Klis \markcite{kl94}(1994). The HBO in Z sources can be explained by the magnetospheric beat frequency model (Alpar \& Shaham\markcite{al85}, 1985; Lamb et al.\markcite{la85}, 1985) according to which the observed QPO frequency is the difference between the neutron star spin frequency and the frequency at which blobs of matter orbit the neutron star at the magnetospheric radius. In this model the frequency increases with $\dot{M}$ and decreases with neutron star spin frequency and $B$. The HBO is observed at frequencies between 15 and 60 Hz. LFN is found with a cut-off frequency between 2 and 20 Hz; it usually appears and disappears together with the HBO. In quantitative models the LFN is naturally produced as an extra component to the HBO, with a total power that is comparable to that in the HBO (Shibazaki \& Lamb \markcite{sh87}1987). Both components get stronger with photon energy. According to HK89 the properties of atoll and Z sources are determined by $\dot{M}$ and $B$. Spectral modeling based on this picture shows that of all atoll sources, the persistently bright subclass, especially GX 13+1, have $\dot{M}$ and $B$ closest to those of the Z sources (Psaltis \& Lamb \markcite{ps96}1996). On the basis of the magnetospheric beat frequency model one might therefore expect to see HBO-like phenomena in these sources. The detection of a HBO-like phenomenon in GX 13+1 seems to confirm these expectations. The fact that the QPO frequency in GX 13+1 is at the high end of the HBO frequency range in Z sources can be explained by a slower spinning neutron star or/and by a $B$ that is lower than in the Z sources, but still high enough to produce HBO. A higher $\dot{M}$ than Z sources is unlikely in view of the luminosity. | 98 | 3 | astro-ph9803267_arXiv.txt |
9803 | astro-ph9803098_arXiv.txt | Under the assumption that the gas mass fraction of galaxy clusters estimated out to an outer hydrostatic radius is constant, it is possible to constrain the cosmological parameters by using the angular diameter distance relation with redshift. We applied this to a sample of galaxy clusters from redshifts of 0.1 to 0.94, for which published gas and total masses are available from X-ray data. After scaling the gas fraction values to the $r_{500}$ radius (Evrard, Metzler, \& Navarro 1996), we find an apparent decrease in gas fractions at high redshifts, which can be scaled back to the mean gas fraction value at z $\sim$ 0.1 to 0.2 of (0.060 $\pm$ 0.002) $h^{-3/2}$, when $\Omega_m = 0.55^{+0.35}_{-0.23}$ ($\Omega_m + \Omega_{\Lambda} =1$, 1-$\sigma$ statistical error). However, various sources of systematic errors can contribute to the change in gas mass fraction from one cluster to another, and we discuss such potential problems in this method. | In recent years, measurements of gas mass fraction in galaxy clusters together with the universal value for the cosmological baryon density as derived from nucleosynthesis arguments have been used to constrain the cosmological mass density of the universe. There are two basic assumptions used in such an analysis: (1) that the galaxy clusters are the largest virialized systems in the universe and based on hierarchical clustering models that clusters represent composition of the universe as whole, and (2) that the gas mass fraction when measured out to a standard (hydrostatic) radius is constant. Evrard (1997) applied this argument to a sample of galaxy clusters studied by David, Jones \& Forman (1995) and White \& Fabian (1995), with redshifts up to 0.2. The gas mass fraction values were scaled to the $r_{500}$ radius first defined by Evrard, Metzler \& Navarro (1996), which has shown to be a conservative estimate of the outer hydrostatic radius based on numerical simulations. When both gas and total (virial) masses are estimated out to this radius, then the assumption that all clusters should have the same gas mass fraction is well justified. The mean gas mass fraction, as derived by Evrard (1997), for the low redshift clusters is $\bar{f}^{\rm X-ray}_{\rm gas} (r_{500}) = (0.060 \pm 0.003)\, h^{-3/2}$. Recently, Myers {\it et al.} (1997) observed the Sunyaev-Zel'dovich (SZ) effect (Sunyaev \& Zel'dovich 1980) in a sample of low redshift clusters, and derived a mean SZ gas mass fraction, $\bar{f}^{\rm SZ}_{\rm gas}$, that ranges from (0.061 $\pm$ 0.011) $h^{-1}$ to (0.087 $\pm$ 0.030) $h^{-1}$ (we refer the reader to Birkinshaw 1998 for a recent review on the SZ effect). Using the mean cluster gas fraction and the nucleosynthesis derived value for the $\Omega_b$, both Evrard (1997) and Myers {\it et al.} (1997) put constraints on the cosmological mass density of the universe, $\Omega_m$, with some dependence on the Hubble constant. If cluster gas fractions are indeed constant in a sample of galaxy clusters when calculated out to the outer hydrostatic radius, then it is possible to constrain the cosmological parameters based on their dependence in the angular diameter distance relation with redshift. In Section 2 of this paper, we present the potential possibility of using cluster gas fraction as a {\it standard candle} and apply this to a sample of $\sim$ 53 clusters that range in redshifts from 0.1 to 0.94, for which published data on masses are available from literature. In Section 3, we discuss various systematics uncertainties and possible selection effects in this method. | Under the assumption that the baryonic (gas) fraction in galaxy clusters are constant, we have constrained the cosmological parameters using the angular diameter distance relation with redshift. However, the present X-ray based data from various studies may be providing a biased result due to unknown selection effects. When the angular diameter distance dependence on the traditional Hubble constant measurements using SZ and X-ray data are included with the present method involving cluster gas mass fraction, it is likely that tighter constraints on the cosmological parameters are possible. We hope to explore this possibility in a future paper. Given the biases that may go in to the presented gas mass fraction vs. redshift diagram, we suggest that results from a well defined and random sample of galaxy clusters, such as that would be available from AXAF in X-ray and from interferometric observations in SZ, be used to constrain the cosmological parameters using the angular diameter distance relation with redshift. | 98 | 3 | astro-ph9803098_arXiv.txt |
9803 | astro-ph9803051_arXiv.txt | We present the results from {\it ASCA} observations of NGC~3783 carried out during 1993 and 1996. Variability is observed both between the two epochs and within the individual observations, the latter due to a non-stationary process. We find the spectra at both epochs to contain features due to absorption by ionized material, predominantly due to O{\sc vii} and O{\sc viii}. We find the opacity for such material to decrease by a factor $\sim$2 in the $\sim 0.7$--1.0~keV band between the epochs while the photoionizing flux increases by $\sim$25\%. By comparison with detailed photoionization models we show this behaviour is inconsistent with that expected from a single-zone of photoionized gas. A deficit in the data compared to such models, consistent with excess absorption by O{\sc viii}, supports the suggestion that the material consists of two or more zones Significant Fe $K$-shell emission is also observed during both epochs. We show this emission is dominated by the asymmetric line profile expected from the innermost regions of the accretion flow. We find no evidence that the Fe emission varied in either shape or equivalent width between the two epochs. | \label{Sec:intro} The almost face-on spiral (SBa) galaxy NGC~3783 ($z = 0.0097$) is one of the brighter Seyfert galaxies in the sky in all wavebands (e.g. Alloin et al. 1995 and references therein). At X-ray wavelengths, NGC 3783 was first detected in the {\it Ariel-V} sky-survey and has subsequently been observed by all major X-ray instruments (e.g. Malizia et al 1997 and references therein). Evidence for deep absorption features in the 0.5--1.5~keV band was first obtained in a {\it ROSAT}\ observation (Turner et al. 1993) and confirmed in a subsequent {\it ASCA}\ observation (George, Turner \& Netzer 1995, hereafter G95). The latter workers demonstrated the opacity was dominated by $K$-shell absorption edges due to O{\sc vii} and O{\sc viii} within a screen of material (with an effective hydrogen column density $\sim 10^{22}\ {\rm cm^{-2}}$) within the column--of--sight to the central source of X-ray emission. Subsequent {\it ASCA}\ observations have revealed such screens of photoionized material (so--called "warm absorbers'') to be a common features in Seyfert galaxies (Reynolds 1997; George et al 1998a, hereafter G98). G95 also suggested the tentative detection of the emission predicted to arise from within such a screen of ionized gas if it subtends a significant solid angle at the central source. Here we present the results from the analysis of four new {\it ASCA}\ observations of NGC~3783 performed in 1996 July, along with a re-analysis of the two previous observations carried out in 1993 December (and as reported previously by G95, Reynolds 1997 and G98). The observations performed in 1996 were each separated by $\sim$2 days, with the hope (based on the behaviour seen during 1993) of measuring the reaction of the ionized screen to changes in the illuminating continuum, and hence better constrain the location of and physical conditions within the screen. Unfortunately, the photoionizing source did not co-operate, exhibiting only relatively low-amplitude variations during the 1996 observations. However, we as we shall show, useful insights regarding the nature of the ionized, circumnuclear material can still be obtained by comparision of the 1996 and 1993 observations. The observations and data-screening employed are described in \S\ref{Sec:data_anal}. In \S\ref{Sec:spatial} we describe the results of our spatial and temporal analysis. The results of our spectral analysis of the absorption due to the ionized gas are presented in \S\ref{Sec:spectral-cont}, confirming the presence of the ionized material during both 1993 and 1996. However we find the opacity of the ionized material to decrease by a factor $\sim$2 between these epochs, and discuss physically realistic explanations. In \S\ref{Sec:nldl} we discuss the constraints which can be placed on the emission due to Fe $K$-shell processes in the 5--7~keV band. Finally, in \S\ref{Sec:Discuss} we review our findings in the context of the results from similar objects and future observations. | \label{Sec:Discuss} \subsection{The ionized-absorber} \label{Sec:disc-ionized} Our observations confirm the presence of substantial opacity due to ionized material within the column--of--sight to the X-ray emitting region in NGC~3783. Furthermore we have found this opacity in the $\sim 0.7$--1.0~keV band to decrease by a factor $\sim$2 in the 18 months between the observations carried out in 1993 and 1996. This large change in opacity is inconsistent with that expected if the ionized material is simply responding to changes in the illuminating continuum. Such behaviour has been seen in at least two other objects (MCG-6-30-15, Fabian et al 1994; NGC~3227, Ptak et al 1994, George et al 1998b). During both epochs we find the bulk of the opacity to be reasonably well modelled by a screen of ionized material. However, in all cases, we find a deficit is present in the data/model residuals. In \S\ref{Sec:93vs96} we briefly explored a number of possible explanations for such a deficit. We suggest that the most attractive explanation is in terms of there being two or more zones of photoionized material (\S\ref{Sec:2zone}). In this case the column variability can be attributed to a change in column of just one of the zones, which might then be inferred to exist as clouds moving into and out of the column--of--sight. Such an hypothesis has been suggested in other objects based on the differential variability in the depth of the O{\sc vii} and O{\sc viii} edges (e.g. NGC~3227, Ptak et al 1994; MCG-6-30-15, Otani et al 1996; NGC~4051, Guainazzi et al 1996), and/or, as here, spectroscopically based on detailed photoionization calculations (e.g. NGC~3516, Kriss et al 1996). Of course future observations will reveal whether the zone referred to as the 'ambient gas' in \S\ref{Sec:2zone} truly always lies within the column--of--sight. A viable alternative is that this zone is in fact simply a slower moving (or larger) structure traversing the the column--of--sight. Absorption features due to C{\sc iv}, N{\sc v} and possibly H{\sc i} have been observed in the ultraviolet (UV) spectrum of NGC~3783 (e.g. Shields \& Hamann 1997, and references therein). Furthermore there is good evidence that the implied column density in these ions is variable (e.g. Maran et al 1996). Based on the column densities implied by the 1993 {\it ASCA}\ observations of the O{\sc vii} and O{\sc viii} edges as given in G95, Shields \& Hamann demonstrated that the predicted column densities in the UV absorption lines to be in moderately good agreement with the observed range. Hence it was suggested that there might be a single zone of ionized material responsible for the absorption features in both the X-ray and UV regimes. Since the UV absorption due to C{\sc iv} is seen imprinted on the corresponding broad emission line, this would clearly place the location of the ionized material at a radius larger than the region emitting the bulk of the broad C{\sc iv} emission ($\gtrsim $ several light-days). This is in agreement with our estimate in \S\ref{Sec:nonequilibrium} that the material had to be at a radius $\lesssim 20$ light-days in order to be in ionization equilibrium. In Table~6 we list the column densities of the abundant Li-like ions predicted from our single-zone model of the ambient gas implied by the {\it ASCA}\ observations in 1996. We find column densities in the strong UV absorption lines within an order of magnitude of those observed, confirming the basic conclusion of Shields \& Hamann. However, as acknowledged by Shields \& Hamann and several other workers (e.g. Mathur 1997 and references therein), such comparisons are extremely sensitive to the form of the photoionizing continuum in the unobserved 13.6--300~eV band, and the assumed elemental abundances. Given these facts, and the observed variability, more detailed comparisons require simultaneous X-ray and UV observations of a wider species of ions such as may be possible with the launch of {\it AXAF}, {\it XMM} and {\it FUSE}. Unfortunately, due to the uncertainities in the calibration of the XRT/SIS instrument below 0.6~keV (and hence our exclusion of these data from our detailed spectral analysis) our data are unable to provide stringent constraints on the intensity of the emission features predicted to arise within the ionized material. However, as noted in \S\ref{Sec:ion_emis}, the limits provided by the current data are consistent with the material subtending a solid angle $\lesssim2\pi$ at the central source. \subsection{The Fe emission line} \label{Sec:disc-line} Our data confirm the presence of intense Fe $K$-shell emission within NGC~3783, with a total eqivalent width ($\sim 300$~eV) and broad, asymetric profile fairly typical of other Seyfert 1 galaxies (e.g. Nandra et al 1997b, and references within). In \S\ref{Sec:nldl} we suggest the line can be deconvolved into two components. The first is a narrow component with a rest-frame energy of 6.4~keV (corresponding to an ionization state $\lesssim$Fe{\sc xii}) with an equivalent width ($\sim 60$~eV) consistent with production within the distant structures within the circumnuclear environment of the central source (Ghisellini, Haadt, Matt 1994). The second component, referred to as the diskline component in \S\ref{Sec:nldl}, is that arising close to $6$ gravitational radii of a putative supermassive black hole. The extreme gravitiational effects suffered by the emission line photons at such radii give rise to an asymmetric profile. We find no compelling evidence that the Fe emission varied in either shape or equivalent width between the 1993 and 1996, indicating no significant changes in the distribution of Fe emissivity. | 98 | 3 | astro-ph9803051_arXiv.txt |
9803 | astro-ph9803045_arXiv.txt | This paper presents a new derivation of the Generalized Poisson distribution. The derivation is based on the barrier crossing statistics of random walks associated with the Poisson distribution. A simple interpretation of this model in terms of a single server queue is also included. In the astrophysical context, the Generalized Poisson distribution is interesting because it provides a good fit to the evolved, Eulerian counts-in-cells distribution measured in numerical simulations of hierarchical clustering from Poisson initial conditions. The new derivation presented here can be used to construct a useful analytic model of the evolution of clustering measured in these simulations. The model is consistent with the assumption that, as the universe expands and the comoving sizes of regions change as a result of gravitational instability, the number of such expanding and contracting regions is conserved. The model neglects the influence of external tides on the evolution of such regions. Indeed, in the context of this model, the Generalized Poisson distribution can be thought of as arising from a simple variant of the well-studied spherical collapse model, in which tidal effects are also neglected. This has the following implication: Insofar as the Generalized Poisson distribution derived from this model is a reasonable fit to the numerical simulation results, the counts-in-cells statistic must be relatively insensitive to such effects. This may be a consequence of the Poisson initial condition. The model can be understood as a simple generalization of the excursion set model which has recently been used to estimate the number density of collapsed, virialized halos. The generalization developed here allows one to estimate the evolution of the spatial distribution of these halos, as well as their number density. For example, it provides a framework within which the halo--halo correlation functions, at any epoch, can be computed analytically. In the model, when halos first virialize, they are uncorrelated with each other. This is in good agreement with the simulations. Since it allows one to describe the spatial distribution of the halos and the mass simultaneously, the model allows one to estimate the extent to which these halos are biased tracers of the underlying matter distribution. | Consider an initially Poisson distribution of particles that clusters gravitationally as the universe expands. In this paper, the initial Poisson distribution will also be called the initial Lagrangian distribution. As time passes, the particle distribution evolves, as, for example, tightly bound virialized clusters (called halos, or dark matter halos, in this paper) form. Thus, the evolved distribution is different from the initial Lagrangian distribution. In what follows, the evolved distribution will be called the Eulerian distribution. The goal of this paper is to use the properties of the initial Lagrangian distribution to derive a reasonable approximation to the form of the evolved Eulerian distribution. In the absence of a model relating the two distributions, the only constraint is that required by mass conservation: the number of particles in the initial and evolved distributions is the same, so the average density, $\bar n$, in the two distributions must be the same. In what follows, quantities measured in the Lagrangian space will have a subscript `0', while those in Eulerian space will not. In this notation, mass conservation implies that $\bar n_0=\bar n$. Studies of clustering from Poisson initial conditions (Itoh, Inagaki \& Saslaw 1993 and references therein) show that when the initial, Lagrangian distribution is Poisson, then the evolved Eulerian distribution is Generalized Poisson. This paper presents a model in which this is so. The model is consistent with three general hypotheses about the evolution of clustering. The first is the hypothesis that, in comoving coordinates, initially denser regions contract more rapidly than less dense regions, and that sufficiently underdense regions expand. The second assumption is that, as the universe evolves, the number of such expanding and contracting regions is conserved---only their comoving size changes. The third is that the influence of external tides on the evolution of such comoving regions can be neglected, if one is only interested in computing statistics such as the mass function of collapsed halos, or the distribution of counts in Eulerian cells. There are no compelling physical arguments for any of these assumptions, and initial particle configurations which violate some or all of these assumptions are relatively easy to construct. That the model predicts a counts-in-cells distribution which is a reasonable approximation to that measured in the numerical simulations suggests that, at least for clustering from Poisson initial conditions, these simple assumptions may also be reasonably accurate. \subsection{The Generalized Poisson distribution}\label{gpdsec} Since it plays a central role in this paper, various known properties of the Generalized Poisson distribution are summarized below. The Generalized Poisson distribution (Consul 1989) has the form \begin{equation} p(N|V,b) = {\bar N(1-b)\over N!}\, \Bigl[\bar N(1-b) + Nb\Bigr]^{N-1} {\rm e}^{-\bar N(1-b) - Nb}\!. \label{gpdf} \end{equation} Here $p(N|V,b)$ is the probability that a cell of size $V$ placed randomly within a particle distribution contains exactly $N$ particles. If $\bar n$ denotes the average density, then $\bar N\equiv \bar nV$. In this paper $0\le b<1$, and, for reasons discussed below, it will be supposed that $b$ is not a function of $V$. The case $b=0$ is the Poisson distribution. Equation~(\ref{gpdf}) is a Compound Poisson distribution (e.g. Saslaw 1989); it arises if point sized clusters, called halos in the following, have a Poisson spatial distribution, and the probability a randomly chosen halo contains exactly $n$ particles is \begin{equation} \eta(n,b) = {(nb)^{n-1}\,{\rm e}^{-nb}\over n!}. \label{borel} \end{equation} This is the Borel$(b)$ distribution (Borel 1942). In this paper, equation~(\ref{borel}) will be called the halo mass function. The Generalized Poisson distribution was first discovered in the astrophysical context by Saslaw \& Hamilton (1984) (also see Sheth 1995a). It provides a good fit to the distribution of particle counts in randomly placed cells, provided the particle distributions have evolved, as a result of gravitational clustering, from an initially Poisson distribution (Itoh, Inagaki \& Saslaw 1993 and references therein). In fact, the fits are significantly improved if $b$ is allowed to increase to an asymptotic value as $V$ increases. This scale dependence is simply a consequence of relaxing the assumption that Borel clusters are point sized, but still requiring that they have some finite size. The asymptotic value of $b$ is that which would have characterized the distribution, had the clusters been point sized (Sheth \& Saslaw 1994). For this reason, the asymptotic value of $b$ is fundamental, and the point sized idealization useful. This paper is mainly concerned with the point sized idealization, so that, in what follows, $b$ is independent of $V$. The point sized idealization is also motivated by the following observation. To a good approximation, the distribution of bound virialized halos in the numerical simulations is Borel$(b)$. Thus, to a good approximation, clustering from Poisson initial conditions evolves in such a way that, at all times, particles are bound up in Borel$(b)$ halos, and, at the time when they first virialize, these halos have a Poisson distribution. The evolution of clustering is parameterized by the time dependence of $b$; it is zero initially, and it increases as the universe expands (Zhan 1989; Sheth 1995b and references therein). Therefore, in the remainder of this paper, $b$ will be treated as a pseudo-time variable, and the Borel$(b)$ distribution will often be called the halo mass function at the epoch $b$. As $V\to 0$, most cells in the Lagrangian and Eulerian distributions will be empty. Equation~(\ref{gpdf}) shows that, in this limit, the probability that a cell is not empty is $\bar N(1-b)$, and the probability that a non-empty cell contains exactly $N$ particles is given by the Borel$(b)$ distribution. In other words, at the epoch $b$, the halo mass function is the same as the vanishing-cell-size limit of the Eulerian counts in cells distribution (Sheth 1996a). This fact will be useful later. The Borel$(b)$ distribution can be derived from a number of different constructions, all of which are related to the Poisson distribution (Epstein 1983; Sheth 1995b; Sheth 1996b; Sheth \& Pitman 1997). In the context of this paper, all these constructions can be thought of as providing models that allow one to compute the Eulerian space distribution, in the limit of vanishing cell size, given that the Lagrangian space distribution is Poisson. One of these constructions, based on the statistics of random walk barrier crossings associated with the Poisson distribution, is the excursion set model (Epstein 1983; Sheth 1995b). This paper shows how to derive the Generalized Poisson distribution from a simple generalization of this excursion set model. The generalization shows how to derive the Eulerian space Generalized Poisson distribution from the Lagrangian space Poisson distribution, for all cell sizes, and all times. \subsection{Outline of this paper} Section~\ref{const} summarizes the random walk, excursion set model which leads to the Borel$(b)$ distribution. Sections~\ref{shift} and~\ref{expdf} describe a generalization of this model which leads to a new derivation of the Generalized Poisson distribution. Section~\ref{halos} shows how to describe the spatial distribution of virialized halos within the context of this model. It shows that the model is consistent with the Compound Poisson interpretation of the Generalized Poisson distribution -- in the model, Borel$(b)$ halos have a Poisson distribution at the time when they first virialize. Moreover, in the model, the $V\to 0$ limit of the counts-in-cells distribution is, indeed, the halo mass function. This shows explicitly that the excursion set approach developed here is able to reproduce the known properties of the Generalized Poisson distribution. The relation between this model and the well-studied spherical collapse model (outlined in Appendix~\ref{scoll}) is discussed in Section~\ref{scm}. Section~\ref{queue} contains a brief digression which relates the excursion set model of the previous section to a simple single server queue system. Section~\ref{scale} discusses a scaling solution associated with the model that is analogous to the scaling solution found in Section~3.2 of Sheth (1995b). Section~\ref{twobar} discusses the associated two barrier problem. The solution of this problem may provide useful diagnostics in assessing the rate of evolution of the Eulerian statistics computed earlier in the paper. Clustering from more general initial conditions, using the techniques developed here, is treated in a forthcoming paper. | This paper presents a new derivation of the Generalized Poisson distribution. The derivation allows one to construct a useful model of hierarchical clustering from Poisson initial conditions. The resulting model is useful because the Poisson assumption allows one to solve many problems that, at present, have no solution if more realistic initial conditions are used. The model is a simple generalization of the excursion set model developed by Bond et al. (1991). Their approach allows one to estimate the mass function of collapsed halos; the generalization presented here allows one to describe the spatial distribution of these halos as well. The model can also be thought of as a simple variant of the spherical collapse model. In the model, initially denser regions contract more rapidly than less dense regions, sufficiently underdense regions expand, the influence of external tides on the evolution of such regions is ignored, and the number of expanding and contracting regions is assumed to be conserved. Strictly speaking, none of these assumptions can be justified physically. However, these simplifications mean that the model can be worked out relatively easily. Moreover, the Generalized Poisson distribution, derived after making these assumptions, is a reasonably accurate fit to the Eulerian counts-in-cells distribution measured in numerical simulations of clustering from Poisson initial conditions. This suggests that, at least for estimating the evolution of the counts-in-cells statistic from such initial conditions, these simplifications are justified. In the model, a collapsed halo occupies a vanishingly small volume. In the simulations, collapsed halos have non-zero sizes---any given halo virializes at some fraction, typically about one half, of its turnaround radius. This means that on scales smaller than that of a typical halo, the counts-in-cells distribution computed here will cease to be a good approximation to that measured in the simulations. As discussed in the introduction, the fact that halos have non-trivial density profiles means that the $b$ parameter in equation~(\ref{gpdf}) depends on scale. A reasonable approximation to the effects of this scale dependence can be computed from models, such as those proposed by Navarro, Frenk \& White (1996), of the density profiles of collapsed halos (see Sheth \& Saslaw 1994 for details). As Poisson initial conditions are not realistic anyway, this seems an unnecessary refinement to an already idealized model. As the basic model has worked out so easily, as it allows one to estimate the extent to which halos are biased tracers of the mass, and, most importantly, as it provides a reasonably accurate description of the evolution of clustering measured in numerical simulations, it seems worth extending it to describe clustering from more general initial conditions. This extension is in progress. | 98 | 3 | astro-ph9803045_arXiv.txt |
9803 | astro-ph9803273_arXiv.txt | We report on Australia Telescope Compact Array observations of the $\sim$$10^5$~yr old pulsar PSR~B0906--49. In an image containing only off-pulse emission, we find a weak, slightly extended source coincident with the pulsar's position, which we argue is best interpreted as a pulsar wind nebula (PWN). A trail of emission extending behind the pulsar aligns with the major axis of the PWN, and implies that the pulsar is moving north-west with projected velocity $\sim$60~\kms, consistent with its scintillation speed. The consequent density we infer for the pulsar's environment is $>2$~cm$^{-3}$, so that the PWN around PSR~B0906--49 is confined mainly by the high density of its surroundings rather than by the pulsar's velocity. Other properties of the system such as the PWN's low luminosity and apparent steep spectrum, and the pulsar's large characteristic age, lead us to suggest that this nebula is substantially different from other radio PWNe, and may represent a transition between young pulsars with prominent radio PWNe and older pulsars for which no radio PWN has been detected. We recommend that further searches for radio PWNe should be made as here: at low frequencies and with the pulsed emission subtracted. | The spin-down observed in most pulsars corresponds to a significant rate of energy loss, which manifests itself primarily in the form of a magnetized wind of relativistic particles (e.g. Rees \& Gunn 1974\nocite{rg74}). Under certain conditions the interaction between this wind and its surroundings is observable, in the form of a pulsar wind nebula (PWN). At radio frequencies PWNe fall into two basic classes, plerions and bow-shock nebulae: plerions (e.g.\ \cite{hb87}) are the filled-center components of supernova remnants (SNRs) in which the pulsar wind is confined by the pressure of hot gas in the SNR interior, while bow-shock nebulae are confined by the ram pressure associated with their pulsar's high velocity (e.g.\ \cite{fk91}; \cite{fggd96}). Both are characterized by significant levels of linear polarization, a centrally-peaked morphology and a flat spectral index ($-0.3 < \alpha < 0$; $S_\nu \propto \nu^{\alpha}$). While only a tiny fraction of a pulsar's spin-down energy goes into producing the radio emission from such nebulae, radio PWNe are valuable diagnostics of the properties of both the wind and of the pulsar itself (e.g.\ \cite{fggd96}). Insight into the workings of such PWNe is limited by the fact that only six pulsars, all with ages $\la10^5$~yr, have associated radio PWNe (\cite{fs97}, hereafter FS97). Attempts to detect radio nebulae around older pulsars have so far been unsuccessful (e.g. \cite{ccgm83}). In an attempt to increase the sample of radio PWNe, FS97\nocite{fs97} used the Very Large Array at 8.4~GHz to search for PWNe around 35 pulsars of high spin-down luminosity and/or velocity, but found no new PWNe down to a surface brightness $T_b \sim 1.2$~K. They concluded that only young, energetic pulsars produce observable radio nebulae, and that the properties of a pulsar's wind may change as the pulsar ages. However, a weak, compact or steep-spectrum nebula is difficult to detect using the approach of FS97, because the nebula is ``hidden'' by the emission from its associated pulsar. This has motivated us to attempt an alternative strategy for finding radio PWNe, whereby visibilities are recorded at high time resolution so that off-pulse images can be produced. The full results of this program are described elsewhere (\cite{sgjf98}); here we report on the discovery of an unusual PWN associated with PSR~B0906--49. PSR~B0906--49 ($l = 270\fdg3$, $b = -1\fdg0$) is a 107~ms pulsar (\cite{dmd+88}) whose characteristic age $\tau_c =$~112\,000~yr and spin-down luminosity $\dot{E} = 4.9 \times 10^{35}$~erg~s$^{-1}$ place it amongst the 5\% youngest and most energetic of all pulsars. \HI\ absorption has yielded distances to the pulsar in the ranges 2.4--6.7~kpc (\cite{kjww95}) and 6.3--7.7~kpc (\cite{sdw+96}), while a distance estimate using the pulsar's dispersion measure (\cite{tc93}) is $6.6^{+1.3}_{-0.9}$~kpc. In future discussion we assign it a distance $7d_7$~kpc. The pulse profile shows a strong main pulse and weaker interpulse, both of which are $\sim$90\% linearly polarized (\cite{wmlq93}; \cite{qmlg95}). Scintillation measurements imply a transverse velocity for the pulsar of $\sim$50$d_7^{1/2}$~\kms\ (\cite{jnk98}). | Using pulsar-gating observations at $\lambda=20$~cm, we have found an off-pulse, slightly extended source coincident with PSR~B0906--49, which we interpet as a faint pulsar wind nebula. An associated trail implies a projected velocity for the pulsar of $\sim$60~\kms\ along a position angle 315\arcdeg. The system is unusual in several ways: \begin{enumerate} \item the nebula has a lower luminosity and steeper spectrum than any other radio PWN yet discovered; \item PSR~B0906--49 is older than any other pulsar known to power a radio PWN, while only PSR~B1853+01 has a lower spin-down luminosity; \item the PWN appears to be generated by a slow ($\sim$60~\kms) pulsar moving through a dense ($>2$~cm$^{-3}$) medium. \end{enumerate} Thus PWN~B0906--49 appears to be very different from other radio PWNe. To account for this, we propose that either the low velocity and high density inhibits shock acceleration, or that the particle spectrum injected by a pulsar into its PWN steepens with age. In the latter case, PWN~B0906--49 may be in transition between the high luminosity, flat-spectrum PWNe seen around young pulsars, and the undetectable radio PWNe around older pulsars. We note that PWNe of low flux density and steep spectrum were heavily selected against in FS97's ungated, high frequency observations. We therefore recommend that future searches, particularly around intermediate-age pulsars, should be made as here: at low frequencies, and by imaging only the off-pulse emission. A collection of PWNe produced under a variety of conditions still needs to be accumulated before we can fully understand how the radio emission from a PWN is produced and how it depends on the properties of its pulsar. | 98 | 3 | astro-ph9803273_arXiv.txt |
9803 | astro-ph9803335_arXiv.txt | We report on the detection of four primordial galaxies candidates in the redshift range $4.55 \le z \le 4.76$, as well as 4 other candidates in the range $4.12 \le z \le 4.32$. These galaxies have been detected with the now common technique of the Lyman break, but with a very original instrumental setup, which allows the simultaneous detection on a single exposure of three narrow ($\mathrm{FWHM} \sim 200$~\AA) bands. Depending on between which pair of bands the Lyman break is located, it thus allows to detect two ranges of redshifts. Two of the the higher redshift candidates are located only $50\h50$~kpc apart, thus forming a possible primordial galaxy pair. | Introduction} The knowledge of the epoch when galaxies formed, that is when they formed the bulk of their stars, is a major cosmological question. The detection of such primordial galaxies would give constraints on their formation, their chemical evolution or the formation of large scale structures. All the attempts to detect these galaxies at redshifts $z\ge3$ by detecting the Ly$\alpha$ emission line, either by visible and IR spectroscopy (Koo \& Kron 1980, Schneider et al. 1991, Thompson \& Djorgovski 1995), narrow band imaging (De Propris et al. 1993, Parkes et al. 1994) or Perot-Fabry (Thompson et al. 1995) have failed, while the Ly$\alpha$ flux of a $z=4$ galaxy, $\sim 10^{-16\pm 1}$~erg~cm$^{-2}$~s$^{-1}$ (Thompson \& Djorgovski 1995), should have been easily observed with 3.6m telescopes. Only recently, Hu (1998) reported the detection of \lya\ emitters at redshifts as high as 4.5. However, it is now known that even small amounts of dust can absorb the great majority of the Ly$\alpha$ photons by resonant scattering (Charlot \& Fall 1993). The most promising method to observe these galaxies seems to use the Lyman break at 912~\AA. This break is present in all kinds of galaxies, for all reasonable stellar formation scenarii (Bruzual \& Charlot 1993), and is observed in the visible range for $z\ge 3.5$. The first very high redshift galaxies have been observed through multi band imaging, with broad filters located on each side of the Lyman break at $\left<z\right>=3.2$, selected using color criteria, and several of the candidates have recently been spectroscopically confirmed with the Keck telescope (Steidel \& Hamilton 1992,1993, Steidel et al. 1995, Pettini et al. 1997). The average projected density of these objects is expected to be 0.5~gal~arcmin$^{-2}$ at $R \sim 25$ and $\left<z\right> = 3.2$ (Steidel et al. 1995). In this paper, we present the results of the search for very high redshift galaxies ($z\sim4.5$ ), using the same technique, but with an original instrumental setup : we have designed multi band filters, with three narrow ($\mathrm{FWHM} \simeq 200$~\AA) bands, placed so that they avoid the most prominent night sky emission lines. This disposition and the very high transmission (above 95\%) of the filters compensate the narrowness of the bands, and reduces by a factor of three the amount of data to reduce, as well as telescope observing time. Observations are described in Sect.~\ref{data}, and our detections are discussed in Sect.~\ref{res}. | 98 | 3 | astro-ph9803335_arXiv.txt |
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9803 | astro-ph9803179_arXiv.txt | The existence of bimodal disks is investigated. Following a simple argument based on energetic considerations we show that stationary, bimodal accretion disk models in which a Shakura--Sunyaev disk (SSD) at large radii matches an advection dominated accretion flow (ADAF) at smaller radii are never possible using the standard slim disk approach, unless some extra energy flux is present. The same argument, however, predicts the possibility of a transition from an outer Shapiro--Lightman--Eardley (SLE) disk to an ADAF, and from a SLE disk to a SSD. Both types of solutions have been found. We show that for any given accretion rate a whole family of bimodal SLE--ADAF and SLE--SSD solutions can be constructed, each model being characterized by the location of the transition radius. | Observations of black hole X--ray binaries (BHXBs, e.g.~V404 Cyg, Nova Muscae, A0620-00, GRO J1655-40, Cyg X-1) suggest that a bimodal accretion disk may be present in these sources (Thorne \& Price \cite{thorneprice:1975}; Shapiro, Lightman \& Eardley \cite{sle:1976}; Narayan, McClintock \& Yi~\cite{narclintyi:1996}; Esin, McClintock \& Narayan \cite{esinmcclnar:1997}). The inner part is a hot, optically thin, quasi--spherical accretion flow, producing a non--thermal spectrum via synchrotron radiation and Comptonization, with a cut--off around 100 keV. The outer part is a geometrically thin, optically thick disk (Shakura \& Sunyaev \cite{shaksuny:1973}; Frank, King \& Raine \cite{frankkingraine:1992}) which emits a multi--color blackbody spectrum peaked around a few keV. By making plausible assumptions on how the transition radius $R_{\trans}$ depends on the accretion rate $\dot M$, one obtains a one--parameter family of spectra, once the black hole mass $M$ and the viscosity parameter $\alpha$ are specified. The different spectral states of the BHXBs can then be explained by varying the accretion rate $\dot M= \dot m\dot M_{Edd}$, where $\dot M_{Edd}$ is the Eddington accretion rate. Historically the hot, inner region was first thought to be a cooling--dominated, optically thin accretion disk, the SLE disk (Shapiro, Lightman \& Eardley \cite{sle:1976}, henceforth SLE). However, such disks are subject to a violent thermal instability, which makes them useless for constructing stationary bimodal models. An alternative to the SLE disk could be the recently discovered hot, advection dominated accretion flows (ADAFs: Narayan \& Yi \cite{narayanyi:1994}, \cite{narayanyi:1995}; Abramowicz et al. \cite{abrchen:1995}), which are stable against both thermal and viscous perturbations. ADAF models have been quite successful in reproducing the spectral properties of BHXBs in quiescence and of the Galactic center source Sgr A$^{*}$ (Narayan, Yi \& Mahadevan \cite{naryimah:1995}). A bimodal disk structure, in which an outer Shakura--Sunyaev disk (SSD) connects smoothly to the inner ADAF, seems very promising in explaining the whole range of spectral states of BHXBs, from the quiescent to the high state (Esin, McClintock \& Narayan \cite{esinmcclnar:1997}). Advection dominated accretion flows are dynamically very stable. The robustness of the ADAF solutions led Narayan \& Yi (\cite{narayanyi:1995}) to suggest that perhaps nature will always select this option whenever it becomes available. Yet, the precise mechanism through which the SSD material is converted into an ADAF remains a matter of debate. Observations of Cyg X-1 in the low state suggested that the transition from a cold to a hot disk might result from the secular instability of the radiation--pressure dominated inner part of the SSD (SLE; Ichimaru \cite{ichimaru:1977}). The even more violent thermal instability of the radiation--pressure dominated region (Pringle, Rees \& Pacholczyk \cite{pringreespac:1973}; Piran \cite{piran:1978}) might also be held responsible for such a transition. The drawback is that the transition radius for these models is usually quite close to the black hole, $\rad_{tr}\simeq 45\alpha^{2/21}\dot m^{16/21}m_{*}^{2/21}\rad_g$, where $m_{*}\equiv M/M_{\odot}$ and $\rad_g = 2GM/c^2$ is the Schwarzschild radius (SLE; Frank, King \& Raine~\cite{frankkingraine:1992}), whereas observations of many BHXBs (e.g.~Narayan, Barret \& McClintock \cite{narbarclint:1997}; Hameury et al.~\cite{hamlasclnar:1997}) seem to imply a transition radius $\rad_{tr}\sim 10^4 \rad_g$. Moreover, no self--consistent global solution for such a viscously/thermally driven transition has been found so far, and, more probably, the thermal instability results in a genuinely time--dependent behavior (e.g.~a limit cycle, see Meyer \& Meyer--Hofmeister \cite{meyer2hofm:1981}, Mineshige \& Wheeler \cite{minwheel:1989}; Honma, Matsumoto \& Kato \cite{honmatskat:1993}; Szuszkiewicz \& Miller \cite{szuszkmiller1:1997}, \cite{szuszkmiller2:1997}), rather than in a stationary transition to a hot flow. More promising in this respect are the evaporation models in which matter evaporates from the SSD forming an advection dominated corona (Meyer \& Meyer--Hofmeister \cite{meyermeyhof:1994}; Narayan \& Yi \cite{narayanyi:1995}). The question remains whether we understand why it seems not to be possible to convert an SSD directly into an ADAF at a certain radius. This is the question we address in this paper. We use a simple argument based on energetics to predict which kind of bimodal disk structures should be possible. We conclude that only bimodal accretion flows which are formed by a SLE disk outside and by an ADAF or SSD inside are allowed. We confirm our qualitative prediction by actually constructing global models for both bimodal SLE--ADAF and SLE--SSD solutions. | The goal of this investigation has been to shed some light on the physics of disk transitions. Using a simple argument based on energetics, we have shown that global, stationary bimodal solutions are possible if the outer disk satisfies a simple condition. Our main result is that only an SLE disk can match an inner ADAF (or SSD), while a bimodal accretion configuration formed by an outer SSD and an inner ADAF is not permitted, at least within the standard slim disk physics. The two allowed bimodal structures, an outer SLE disk glued to an inner ADAF or SSD, have been computed numerically. We found that, in both cases, the transition radius turns out to be arbitrary, so a whole family of solutions exists. This may sound odd, since there is no physical reason to ask the sound speed to take a precise value at an intermediate radius, as we did in finding our numerical solutions. In principle the condition should be placed at the outer edge. In doing so, however, one should be extremely accurate in fixing $c_s(\rad_{out})$ to find the transition at the desired radius since most bimodal solutions correspond to very nearly identical values of $c_s$ (or any other variable) at the outer or inner edge. It remains nevertheless true that the structure of bimodal disks is determined once a complete set of conditions is prescribed at the edges and at the sonic radius, without the need of any extra piece of physics. The existence of the SLE--ADAF and SLE--SSD bimodal solutions agrees with the energy argument of section \ref{sec-therm-trans}, and strengthens its plausibility. The same argument indicates that a SSD--ADAF model is not allowed. There is a caveat here, however. The essence of the energy argument lies in recognizing that the non--locality of the advective cooling term in the energy equation can be removed by performing a transformation to the fluid frame, and doing the thermal analysis in the frame comoving with the fluid. If, on the other hand, thermal conduction is introduced, the energy equation becomes a second order diffusion type equation and there is no way of removing the non--locality by changing the observer's frame. In this case our argument definitely does not apply. Honma (\cite{honma:1996}) has shown that SSD--ADAF structures are possible when the thermal heat flux from the hot ADAF into the SSD is accounted for. In such accretion flows the location of the transition follows self--consistently from the model, contrary to what we have found in this paper. This is because the heat flux will evaporate per unit time an amount of matter from the SSD which needs to be in balance with the actual accretion rate. The transition radius will therefore automatically adjust itself until this balance is reached. In concluding, we would like to state clearly that, because of the thermal instability of the SLE disk, the models presented here most probably do not exist in nature. | 98 | 3 | astro-ph9803179_arXiv.txt |
9803 | astro-ph9803209_arXiv.txt | We derive an analytic expression for the intensity of resonance-line radiation ``trapped'' in a semi-infinite medium. Given a source function and destruction probability per scattering, the radiation pressure due to trapped photons can be calculated by numerically integrating over analytic functions. We apply this formalism to a plane-parallel model stellar atmosphere to calculate the radiation pressure due to Lyman-$\alpha$ photons produced following absorption of UV and X-rays from an AGN. For low surface gravity stars near the AGN ($g \sim 10\:\cmss,\ r \sim 0.25 {\rm\:pc}$), we find that the pressure due to Lyman-$\alpha$ photons becomes an appreciable fraction of that required for hydrostatic support. If the broad emission line emitting gas in AGNs and QSOs consists of stellar outflows, it may be driven, in part, by Lyman-$\alpha$ pressure. | \label{intro} Interest in the effect of hard UV and X-ray radiation on stellar atmospheres and stellar evolution began in earnest with work on X-ray binaries (e.g., \cite{DOst73}, \cite{BSun73}, and \cite{Arons73}), but the subject has since been revisited in the context of active galactic nuclei (AGN). Consider a star of effective temperature $T_{\rm eff}$ and an AGN of luminosity $L = 10^{46}L_{46}\:\ergs$. The incident power per area from the AGN is equal to $\sigma T_{\rm eff}^4$ at the ``heating'' distance \begin{equation} d_h = \left(\frac{L}{4\pi\sigma T_{\rm eff}^4}\right)^{\onehalf} = 1.5\times 10^{17} L_{46}^{\onehalf}\left( \frac{5000 {\rm K}}{T_{\rm eff}}\right)^2 \:{\rm cm}; \label{eq:rheat} \end{equation} heating of the stellar surface by the AGN is important at distances $d \lesssim d_h$. Fabian (1979) was apparently the first to suggest that heating by AGN radiation might substantially affect the envelopes of nearby stars, resulting in enhanced mass loss. Edwards (1980) argued that when a star approached within $d \lesssim d_h$, the irradiated photosphere would develop supersonic horizontal winds carrying heat from the illuminated hemisphere to the shadowed side. Edwards conjectured that enhanced mass loss could then occur, resulting in a ``cometary star'' as the stellar outflow is accelerated radially away from the AGN\@. Matthews (1983) noted the importance of direct radiation pressure from the AGN, arguing that this could ablate matter tangentially from the stellar photospheres of giant or supergiant stars. Voit \& Shull (1988) discussed the effects of X-rays on the upper atmosphere of a star near the AGN\@. They argued that X-ray heating of a $\sim 1 M_{\sun},$ $R_{*} \approx 100 R_{\sun}$ star ($g\approx 3\:\cmss$) would result in a hot, ionized wind, with temperature at the critical point $T \approx 3\times 10^5 {\rm K}$; radiation pressure in C IV, N V, and O VI resonance lines would contribute to the acceleration of the wind. They concluded that stars which are already red supergiants could develop winds with $\dot M \approx 10^{-7}L_{46}^{0.9}d_{18}^{-2}R_{*,100}^2 M_{\sun} {\rm yr^{-1}}$, where the distance from the AGN is $d = 10^{18} d_{18}{\:\rm cm}$, and $R_{*,100} \equiv R_*/100R_{\sun}$. These mass loss rates exceed normal mass loss rates for red supergiants only within a distance of $d_{18} \lesssim 0.3 L_{46}^{0.45} R_{*,100}\:{\rm cm}$. They estimated the mass loss due to ablation by radiation pressure (\cite{Matthews83}) for a $\sim 1 M_\odot$ star to be $\dot M \approx 10^{-8}L_{46}d_{18}^{-2}R_{*,100}^{5/2}$ provided $d_{18} < 0.6L_{46}^{\onehalf} R_{*,100}$. There has also been some exploration of how AGN radiation might affect the evolution of stars (\cite{Matthews83}; \cite{Verbuntetal84}; \cite{Tout89}). Tout \etal discussed the evolution of stars immersed in a blackbody background with temperature $T = 10^{3.75 - 4} {\rm K}$, concluding that the main sequence evolution is largely unaffected, but that the star will expand to a radius $\sim 10$ times larger than usual during the ``red giant'' phase of evolution. Note, however, that the energy density of AGN radiation equals a $10^4{\rm K}$ blackbody at a distance of only $2\times 10^{16} L_{46}^{0.5}\:{\rm cm}$ from the AGN\@. Furthermore, it is not clear how their conclusions would have to be modified for the anisotropic and nonthermal radiation field of the AGN\@. These discussions of the effects of AGN radiation on stellar atmospheres and evolution have led some to consider a ``bloated stars scenario'' (BSS) to explain the origin of the AGN broad line region (BLR). The BSS proposes that the BLR consists of the stellar winds and/or expanded envelopes of stars irradiated with AGN radiation (e.g., \cite{Pen88}, \cite{Kaz89}). The attractiveness of this scenario stems from its efficient use of known resources: stars are believed to be present near the ``central engines'' of AGN, their gravity provides a possible ``containment'' mechanism for BLR clouds, there is ample mass with which to replenish the clouds, and the mass shed provides material to fuel the AGN\@. So far, however, no specific model has successfully explained how all these mechanisms might function, though pieces of the puzzle have been explored (\cite{AN94}). High signal/noise measurements of line profiles indicate that the number of discrete clouds or bloated stars must be large, perhaps exceeding $\sim 3\times 10^6$ for Mrk~335 (Arav \etal 1997) and $\sim 3\times 10^7$ for NGC~4151 (Arav \etal 1998). Such large numbers appear to challenge the BSS, but the difficulties may not be insurmountable. Recent observations and photoionization models of Baldwin \etal (1996) and Ferland \etal (1996) offer some concrete evidence for the BSS\@. They find that the BLR seems to contain several distinct components. One (component ``A'') has sharp (FWHM $\approx 1000\:\kms$), symmetric line profiles centered on zero velocity, while another (component ``B'') appears to be outward-flowing, peaked at zero velocity but with a long blue tail (down to $-11,000\:\kms$). They interpret ``A'' as the expanded envelopes of stars, and ``B'' as their radiatively accelerated, outflowing winds. The aim of the present paper is to call attention to the fact that irradiation of a star by hard X-rays from the AGN will result in the generation of Lyman-$\alpha$ (\lya) photons within the stellar atmosphere, and the pressure of these trapped resonance-line photons may contribute to mass loss. The possible importance of \lya\ photons was suggested previously by Puetter [unpublished, cited in Penston (1988) and Voit \& Shull (1988)]. Voit \& Shull rejected Puetter's suggestion, arguing that the \lya\ pressure could be estimated to be \begin{equation} P_{\rm Ly\alpha} \approx \left(\frac{aT^4}{3}\right)\frac{W_\lambda}{\lambda} \label{eq:PVoitShull} \end{equation} where $W_\lambda$ is the equivalent width of the \lya\ absorption profile of the gas between the point of interest and the surface, and $T$ is the gas temperature. Voit \& Shull then used $W_\lambda/\lambda < 1$ to obtain an upper limit $P_{\rm Ly\alpha} \lesssim aT^4/3$ which, for stars of interest, is insufficient to ``bloat'' the atmosphere. However equation (\ref{eq:PVoitShull}) presumes the radiation field within the gas to be close to a blackbody at the gas temperature $T$ --- an assumption which need not be valid when \lya\ photons are being generated within the gas by nonthermal processes, such as H(1s$\rightarrow$2p) excitation by photoelectrons and secondary electrons in a non-Maxwellian ``tail'' to the electron energy distribution function. It is therefore essential to estimate the rate of production of \lya\ photons, and to examine their diffusion in physical space and frequency space as well as the possibility of photon ``destruction'' by, for instance, collisional de-excitation of electronically excited H atoms. In \S\ref{sec-linetrans}, we give a brief treatment of resonance-line transfer, extending the detailed calculations of Neufeld (1990). In particular, we derive an expression for the case of a semi-infinite, absorbing slab of material. Since resonance-line photons created in a stellar atmosphere will almost certainly not be able to scatter through the entire star without being absorbed, this limit is an excellent approximation. In \S\ref{sec-model}, we apply our calculational method to a simple, plane parallel, irradiated model atmosphere. In \S\ref{sec-results}, we present the results. In \S\ref{sec-summ} we summarize and discuss implications and further directions for work. | \label{sec-summ} We have developed a formalism for computing the \lya\ pressure in a plane-parallel stellar atmosphere. The treatment of resonance-line trapping given in \S 2 can be incorporated into already developed stellar atmosphere codes. For conditions believed representative of the BLR (ionizing flux of $\sim 4\times 10^{10}\:\ergcmcms$), stars with $g\lesssim 10\:\cmss$ will develop an appreciable pressure due to trapped \lya\ photons in their atmospheres. We have neglected numerous complicating factors, all of which appear likely to enhance the rate of mass loss. First of all, the gravitational acceleration $g\propto r^{-2}$, rather than the assumed $g=\mbox{constant}$. As a result, the outer layers of the atmosphere will have a larger scale height and must ultimately go over into a thermal wind, even if the lower layers are essentially hydrostatic and stable. With the outer layers of the atmosphere photoionized and heated to $T_4 \gtrsim 1$, a fluid element will have positive enthalpy at $r \gtrsim 600\, T_4^{-1} (M/M_{\sun}) R_{\sun}$. Furthermore, we have neglected the radiation pressure due to resonance lines other than \lya, as well as the radiation from the star itself. Ferland (1997) reports that including radiation pressure from all resonance lines rather than just a select few sometimes increases the total radiation by an order of magnitude (using CLOUDY's escape probability approximation, of course). Since no part of our derivation in \S 2 depends explicitly on the resonance line being \lya, it is straight forward to extend our calculation to include other resonance lines. Most important, however, is the fact that except along the ``axis'' --- the line from the center of the star to the AGN --- the atmosphere will be subject to substantial tangential stresses due to three effects related to anisotropic irradiation by the AGN: \begin{enumerate} \item The momentum deposited by absorbed photons --- at anywhere other than the axis, the momentum will have a tangential component; \item Transverse gradients in the \lya\ pressure --- at any particular radial column density $N_{\rm H}$, the \lya\ production rate will be greater along the axis, where X-rays are entering radially; and \item Transverse gradients in the gas temperature --- the rate of X-ray heating will be largest along the axis, where the zenith angle is zero. \end{enumerate} These tangential stresses must result in significant tangential flows, with flow velocities that might approach escape speeds along the ``terminator,'' perhaps resulting in a ``cometary star'' as envisioned by Edwards (1980) and Matthews (1983). It would obviously be of great interest to calculate axisymmetric fluid-dynamical models of red giant atmospheres subject to intense X-ray irradiation in order to investigate the effects of tangential stresses. It should be noted, of course, that we lack a quantitative theory for mass loss even from unperturbed evolved stars, so we are hardly poised to develop a definitive treatment including this new complication. Nevertheless, it does appear likely that trapped \lya\ may help drive mass loss in low-$g$ stars subject to intense X-ray irradiation. The extra \lya\ radiation pressure and associated tangential stresses might induce such evolving stars to ``bloat'' further and lose mass prematurely, so that the population of stars near the AGN core could be skewed from the average throughout the galaxy. Finally, we note that none of our calculations are dependent on the gravitating object being a star. Recently, Walker and Wardle (1998) have suggested the existence of a population of cool, self-gravitating clouds in the galactic halo to account for ``Extreme Scattering Events'' which occur in monitoring the flux of compact radio sources. They propose clouds of approximately a Jovian mass with a radius of about an AU, which would correspond to a gravitational acceleration $g \sim 6\times 10^{-4}\:\cmss$ at the ``surface.'' With a total column density of $\sim 10^{26}\:{\rm cm}^{-2}$, the majority of the AGN energy flux would be absorbed well before reaching (and disrupting) the center of the cloud. The proposed particle number density of $\sim 10^{12}\:{\rm cm}^{-3}$ corresponds well to that inferred by Ferland \etal (1996) and Baldwin \etal (1996). Thus if these self-gravitating, sub-stellar clouds exist in host galaxies of AGN, they seem to be potential candidates for the origin of the BLR\@. It is unclear, though, how they could exist with a high metal content, as required by BLR observations, without rapidly cooling and collapsing (in much less than a Hubble time) prior to exposure to AGN irradiation. Although we conclude that X-ray induced \lya\ pressure would not be significant for main sequence stars near AGN, it could be important either for lower gravity stars which have evolved off the main sequence up the giant branch, or for substellar, self-gravitating clouds. It remains to be seen whether the resulting cometary structures can account for the required BLR covering factor and cloud population, and the AGN mass supply. | 98 | 3 | astro-ph9803209_arXiv.txt |
9803 | astro-ph9803252_arXiv.txt | We have demonstrated the techniques which we have developed in the RVSAO suite to create a system for the accurate, automated reduction of spectra for galaxy redshifts and stellar radial velocities. More than half of all published redshifts have been measured using these techniques, as well as a large number of stellar radial velocities. The correlation method for obtaining redshifts can be successfully extended from absorption line spectra to emission line spectra, with a substantial improvement in effectiveness over the previous method for obtaining emission line redshifts, automated line fitting. The reduction of emission-line spectra requires different reduction steps than absorption line correlations. Emission line correlation redshifts are susceptible to blunders due to the presence of cosmic rays. However, using automated line fitting ({\bf emsao}) and absorption line correlation velocities the blunder rate can be kept near zero, with the degree of automation kept high. We have developed new techniques for calibrating and characterizing the blunder rate and the individual errors in redshift measurements. The blunder rate for RVSAO reductions can be kept near zero by the use of some simple heuristics to identify possible mistakes. For typical redshift survey data from the FAST spectrograph the automation rate is 95\%. Our self calibrating internal error estimator is accurate to $\sim$ 20\%. Large, stable surveys enable development of more accurate and stable error estimators. We have developed new methods for creating, calibrating, and using galaxy redshift templates. We have created an emission line template, femtemp97, having the median properties of a large set of strong emission line spectra. We have created an absorption line template, fabtemp97, having the mean properties of a large set of absorption line spectra showing no sign of emission. These spectra arise from physically distinct processes, and can be used to form a pair of basis vectors to perform a 2-D spectral classification. We have developed a new method for establishing the zero point for redshift observations, a method which minimizes the systematic differences between emission and absorption line redshifts. This zero point is determined as accurately as we can establish the wavelength calibration using standard HeNeArFe lamps. We have shown a technique for measuring and eliminating differences between the instrumental zero point and the true zero point. We have shown improved techniques for a number of the substeps necessary to obtain accurate redshifts, including: removal of emission(absorption) lines when correlating against an absorption(emission) line template; suppression of the night sky lines; supression of the continuum; design of the Fourier filter; and zero padding of spectra. The rapid development of large aperture telescopes with multi-object spectrographs presents substantial challenges for redshift and radial velocity reductions. Reducing one or two orders of magnitude more spectra of objects one or two orders of magnitude fainter while maintaining high quality control standards and minimal personnel costs is clearly a difficult problem. RVSAO provides a solid methodological and software basis to meet these new challenges. | 98 | 3 | astro-ph9803252_arXiv.txt |
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9803 | astro-ph9803314_arXiv.txt | We report the detection of faint emission in the high-excitation [\ion{O}{iv}] 25.90$\mu$m line in a number of starburst galaxies, from observations obtained with the Short Wavelength Spectrometer (SWS) on board ISO. Further observations of \object{M~82} spatially resolve the [\ion{O}{iv}] emitting region. Detection of this line in starbursts is surprising since it is not produced in measurable quantities in \ion{H}{ii} regions around hot main-sequence stars, the dominant energy source of starburst galaxies. We discuss various models for the formation of this line. [\ion{O}{iv}] that is spatially resolved by ISO cannot originate in a weak AGN and must be due to very hot stars or ionizing shocks related to the starburst activity. For low-excitation starbursts like \object{M~82}, shocks are the most plausible source of [\ion{O}{iv}] emission. | Mid-infrared fine structure lines are powerful probes of dusty and obscured galactic nuclei, being able to penetrate extinctions up to the equivalent of A$_V\sim 50$. Using the Short Wavelength Spectrometer (SWS) on board the Infrared Space Observatory (ISO), it is possible to detect faint lines and sources. The rich observed spectra can be used for a detailed modelling of the ionizing spectra of starbursts (e.g. Rigopoulou et al. \cite{rigo96}, Kunze et al. \cite{kunze96}) and AGNs (Moorwood et al. \cite{moor96}). Clear differences between their spectra make these lines a valuable new tool for discriminating between AGN and starburst activity in visually obscured galaxies. AGN spectra include emission from highly ionized species and the so-called coronal lines, requiring photons up to $\sim$300\,eV for their creation. In contrast, starburst spectra are dominated by lines of low excitation species, because even hot, massive stars emit few ionizing photons beyond the \ion{He}{ii} edge at 54\,eV. Line ratios like [\ion{O}{iv}]\,25.9$\mu$m / [\ion{Ne}{ii}]\,12.8$\mu$m and [\ion{Ne}{v}]\,14.3$\mu$m / [\ion{Ne}{ii}]\,12.8$\mu$m have been used by Lutz et al. (\cite{lutz96a}) and Genzel et al. (\cite{genzel98}) to establish the dominant source of luminosity in ultraluminous infrared galaxies (ULIRGs). In some of the starburst templates studied, very faint [\ion{O}{iv}] emission was found, about two orders of magnitude weaker than in typical AGNs. Faint [\ion{O}{iv}] emission in starbursts is not relevant for establishing the power source of ULIRGs, but its origin poses an interesting problem because its creation ionization energy is slightly above the \ion{He}{ii} edge. In this letter we examine possible mechanisms for its production. | We have discussed various excitation mechanisms for faint [\ion{O}{iv}] emission from starburst galaxies. In general, starburst-related sources and in particular ionizing shocks provide the most plausible explanation. Weak buried AGNs may be plausible for individual sources but can be ruled out for the best studied case of \object{M~82} whose [\ion{O}{iv}] emitting region has been spatially resolved. In addition, the fairly small scatter in [\ion{O}{iv}] versus starburst luminosity favours a starburst-related origin, since no finetuning of two independent mechanisms is required. | 98 | 3 | astro-ph9803314_arXiv.txt |
9803 | astro-ph9803122_arXiv.txt | In the $z=1.94$ {\CIV} absorption line system in the spectrum of quasar Q1222+228 ($z_{em}=2.04$), we find two clouds which have contrasting physical conditions, although they are only at a 17~{\kms} velocity separation. In the first cloud {\SiII}, {\SiIV}, and {\CII} are detected, and {\AlII} and {\AlIII} column density limits in conjunction with photoionization models allow us to infer that this cloud has a large Si abundance and a small Al abundance relative to a solar abundance pattern. This pattern resembles that of Galactic metal--poor halo stars, which must have formed from such high redshift gas. The second cloud, in contrast, has detected {\AlII} and {\AlIII} (also {\SiIV} and {\CII}), but no detected {\SiII}. We demonstrate, using photoionization models, that Al/Si must be greater than (Al/Si)$_\odot$ in this unusual cloud. Such a ratio is not found in absorption profiles looking through Milky Way gas. It cannot be explained by dust depletion since Al depletes more severely than Si. Comparing to other Al--rich environments, we speculate about the processes and conditions that could give rise to this abundance pattern. | \label{sec:intro} One of the motives for studying quasar (QSO) absorption line systems is to document the photoionization structure, chemical composition, and kinematics of the ISM and halos of high redshift galaxies at a level of detail on par with studies of the Milky Way ISM and Halo. QSO lines of sight that pass through high redshift galaxies sample multiple gaseous structures having a variety of physical conditions such as found in Galactic {\HI} and {\HII} regions, supershells, infalling Halo gas, and material being processed at the Galaxy/Halo interface. Recognizing the abundance patterns, photoionization conditions, and kinematics associated with these various types of structures is a prerequisite for a detailed understanding of the evolution of galactic gas. Based upon low resolution spectra, early QSO absorption line efforts have been limited to simple curve of growth analyses of equivalent widths, providing only a weighted--average of the absorbing gas properties in each system. A great deal has been learned from these studies and a global (statistically based) picture of chemical and ionization evolution has been suggested [see Steidel (1993\nocite{steidel93}) and references therein]. The next logical step is to examine {\it variations\/} in the chemical and ionization conditions within single galaxies and to incorporate kinematic information. Using high resolution UV spectra, several researchers studied the Milky Way Disk and Halo at a detailed level (\cite{wel97}; \cite{savaraa}; \cite{snow96}; \cite{sf95}; \cite{fs94}; \cite{sf93}). Some recent efforts have focused on the cloud--by--cloud conditions in high redshift galaxies (cf.~\cite{tripp97}; \cite{pet94}). It is hoped that, ultimately, a statistical picture of the processes that give rise to the observed absorbing gas properties will improve our understanding of the present--epoch Milky Way and galactic evolution. In this {\it Letter\/} we study the cloud--to--cloud properties in a {\CIV} system at $z \sim 1.94$ along the line of sight toward the quasar Q$1222+228$ in order to: 1) demonstrate a large variation in ionization and/or abundance conditions between two kinematically adjacent clouds in the same absorber; and 2) present an unusual cloud that has an Al/Si abundance ratio enhanced by a factor of several relative to the solar ratio. | \label{sec:summary} Two kinematically adjacent, absorbing clouds in the same galaxy at $v\sim0$~{\kms} and $\sim 17$~{\kms} have very different abundance patterns, if they are in photoionization equilibrium. Based upon CLOUDY models, Cloud A is inferred to have Si enhanced relative to a solar abundance pattern, and Al underabundant relative to solar. This abundance pattern is characteristic of Milky Way Halo stars (\cite{lau96}; \cite{savaraa}) which could have formed from gas like that observed in this $z=1.94$ galaxy. We would expect such a pattern for many clouds in high redshift galaxies. Independent of the details, Cloud B has Al several times enhanced relative to Si (compared to the solar abundance pattern). This Al enhancement is highly unusual. If depletion onto dust grains is important we would expect a smaller Al/Si since Al depletes more readily than Si. To date, such an Al enhanced pattern has not been seen in absorption profiles looking through Milky Way interstellar gas (\cite{savaraa}; \cite{snow96}; \cite{sf95}; \cite{fs94}; \cite{sf93}). How could an enhancement of Al originate in a $z=1.94$ cloud? In an attempt to find clues we note three astrophysical environments in which Al enhancement is observed: 1) in the stellar photospheres of the most metal poor globular clusters (\cite{shet96}; \cite{smith96}); 2) in the broad line regions of some AGNs\footnote{We note that the absorber lies 30,000 {\kms} from the emission redshift of the QSO, which is not a BAL AGN.} (\cite{shields97}); and 3) in the photospheres of Milky Way Bulge stars (\cite{mr94}). The common theme for the enhancement of Al in these three seemingly different environments is novae. The novae could produce the Al either directly in their ejecta (\cite{smith96}) or indirectly by providing magnesium isotopes to the ISM that later deposit onto stellar photospheres (\cite{langer}). The $^{26}$Mg and $^{27}$Mg isotopes would then be converted to Al through proton capture in deep CNO convective mixing layers in metal poor stars (\cite{langer}). More generally, does this suggest that a concentration of novae contributed to enhancing the Al in this cloud? In the globular cluster environment one key is to have a large number of stars that were formed coevally. Another key, which may apply to all three environments, is to have a potential well large enough to retain the gas. These may be prerequisites in our case also. Another possibility for excess Al production is a particular class of supernova for progenitors over a narrow mass range. The predicted amount of enhancement relative to other elements depends on the specific supernovae model adopted (\cite{nom97}), however only a small subclass of models would give Al/Si in agreement with that inferred for this cloud. Just how unusual is this cloud with a large Al/Si ratio? Such a pattern has not been seen in absorption along dozens of lines of sight toward Milky Way disk and halo stars. Most high redshift clouds do not have large {\AlIII} and {\AlII} (relative to {\SiII}). However, in the $z = 2.14$ damped {\Lya} absorber in Q$0528-251$ (\cite{lu96}), we have identified a single outlying cloud (at $\sim -100$~{\kms}) for which the equivalent width of the {\AlII} $\lambda 1671$ transition is larger than that of the {\SiII} $\lambda 1527$ transition. The other clouds in the same system clearly have the opposite ratio. Finding more examples and establishing similarities between Al--rich environments are logical next steps in diagnosing the origin of this abundance pattern at high redshift and perhaps understanding the anomalous enhancements seen in metal--poor globular cluster and Bulge stars. | 98 | 3 | astro-ph9803122_arXiv.txt |
9803 | astro-ph9803158_arXiv.txt | s{ We estimate the dust polarized emission in our galaxy at high galactic latitudes, which is the dominant foreground for measuring CMB polarization using the high frequency instrument (HFI) aboard Planck surveyor. We compare it with the level of CMB polarization and conclude that, for angular scales $\le 1^{\circ}$, the scalar-induced CMB polarization and temperature-polarization cross-correlation are much larger than the foreground level at $\nu \simeq 100 \, \rm GHz$. The tensor-induced signals seem to be at best comparable to the foreground level.} | The forthcoming satellite CMB projects MAP and Planck surveyor hold great promise for detecting the CMB polarization. A major stumbling block this detection is the unknown level of galactic polarized foregrounds. In this paper, we attempt to estimate the level of dust polarized emission at high galactic latitudes. We model this emission using the three-dimensional HI maps of the Leiden/Dwingeloo survey at high galactic latitudes and the fact that the dust emission, for a wide range of wavelengths, has a tight correlation with the HI emission maps of this survey~\cite{boul}. Assuming the dust grains to be oblate with axis ratio $\simeq 2/3$, which recent studies support~\cite{hil}, we determine the intrinsic dust polarized emissivity. The distribution of magnetic field with respect to the dust grain distribution is quite uncertain, we thus consider three extreme cases (to be described below). | 98 | 3 | astro-ph9803158_arXiv.txt |
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9803 | astro-ph9803228_arXiv.txt | We investigate the implications of a very thick (scale height 1.5 - 3.0 kpc) disk population of MACHOs. Such a population represents a reasonable alternative to standard halo configurations of a lensing population. We find that very thick disk distributions can lower the lens mass estimate derived from the microlensing data toward the LMC, although an average lens mass substantially below $0.3\Msol$ is unlikely. Constraints from direct searches for such lenses imply very low luminosity objects: thus thick disks do not solve the microlensing lens problem. We discuss further microlensing consequences of very thick disk populations, including an increased probability for parallax events. | Current data from the MACHO collaboration \cite{MACHOmass} indicate that in the context of a spherical isothermal model with a Maxwellian velocity distribution, some significant fraction of the Galactic halo is composed of MACHOs with masses roughly in the range 0.1 to 1.0 $\Msol$. Such masses are consistent with several astrophysical candidates for MACHOs -- white dwarfs, neutron stars, and black holes -- each of which presents serious challenges for stellar formation and evolution theories. However, the MACHO component of the halo, if it is not the major component, as in Cold Dark Matter scenarios, may have a very different distribution from the typically assumed spherical isothermal model. The MACHO distribution may be in a significantly flattened halo and/or, due to dissipation, more centrally condensed. In addition, such a distribution might have a significant rotational component. These possibilities have strong implications not only for the MACHO fraction of the halo, but also for the mass estimates derived from the event durations. The velocity dispersion and mass density distribution which describe the halo model are crucial input parameters in extracting an estimate of the lens mass from the data. The event duration is given by the radius of the Einstein ring divided by the MACHO velocity transverse to the line of sight, where the Einstein ring radius is a function of the position of the MACHO along the line of sight and the lens mass. The masses of the lenses can only be determined statistically, in the context of an a priori assumption about the distribution and velocity dispersion of the lenses. Such an analysis, using a spherical isothermal model yields a central mass estimate of $\approx 0.4\Msol$. However, as discussed above, the MACHO distribution and velocity dispersion may be very different from that assumed in the standard halo model. Earlier work \cite{us_rotate} in exploring models with a highly flattened halo and a bulk rotational component has suggested that a very highly condensed model for the MACHO distribution, such as a thick disk, might reduce the MACHO mass estimate from the current data to a level consistent with brown dwarf candidates. Previous explorations \cite{us_nomacho} have indicated that such models may be able to reproduce the observed optical depths toward the Large Magellanic Cloud (LMC) and the Galactic bulge. The overall shape of the Galactic halo is unknown. Attempts to determine the shape of galactic halo potentials from flaring of the outer Galactic gas layer \cite{Olling,Sackett} point to a flattened halo; flattening of the potential is also supported by simulations of the cold dark matter halo formation \cite{nbody}. While it is highly unlikely that the entire halo is in a disk-like configuration, it is not unreasonable to assume that the MACHO component of the halo is significantly more flattened than the dark matter halo. The condensation of a gaseous halo component to form a very thick disk is likely to result in significant star formation and thus the production of a population of lens candidates. In this paper we explore in detail the consequences of such a lens population. We first define the density and velocity structure, and outline the constraints on these distributions. We then present predictions and observational consequences of such thick disks; in particular we determine the expected frequency of parallax events. | Microlensing studies have yielded much exciting data in the past few years and are continuing to survey different lines of sight through the Galaxy in order to probe the Galactic halo. However, the conclusions that can be drawn from the data to date are very model dependent -- assumptions about the distribution of the lenses and their velocity structure have a strong impact on their interpretation. Thus we need to examine a wide range of reasonable lens distributions. Very thick disks present a reasonable alternative to a halo population of lenses. If the lenses are stellar remnants, it seems likely that their configuration will be more condensed than that of a standard non-baryonic halo. While we have found that very thick disks cannot lower the lens mass estimate to the brown dwarf regime, they have the advantage that their total mass in MACHOs is somewhat less than that for a standard halo that is truncated at $50\kpc$ (and much less than a MACHO halo which traces the extended dark halo out to at least $100\kpc$.) A thick disk distribution cannot produce an optical depth toward the LMC of more than about $1.5 \times 10^{-7}$. However, it can explain (within the experimental uncertainties) all or a significant fraction of the current optical depth estimates. As more events are detected, it may be possible to distinguish between very thick disk and halo lens distributions. The most promising avenue for such a discriminant is the observation of parallax events. Because disk lenses would be both closer and on average slower than halo lenses we expect a higher rate of such events for a disk population. Although the survey experiment light curve measurements can only marginally discriminate between disk and halo distributions for reasonable numbers of events, a modest expenditure of telescope time to obtain one percent photometry on all the Magellanic cloud events should be capable of making the distinction very clearly. | 98 | 3 | astro-ph9803228_arXiv.txt |
9803 | hep-ph9803479_arXiv.txt | Contrary to naive cosmological expectations, all evidence suggests that the universe contains an abundance of matter over antimatter. This article reviews the currently popular scenario in which testable physics, present in the standard model of electroweak interactions and its modest extensions, is responsible for this fundamental cosmological datum. A pedagogical explanation of the motivations and physics behind electroweak baryogenesis is provided, and analytical approaches, numerical studies, up to date developments and open questions in the field are also discussed. | \label{Intro} The most basic distinction drawn between particles found in nature is that between particle and antiparticle. Since antiparticles were first predicted \cite{Dirac1,Dirac2} and observed \cite{Anderson1,Anderson2}, it has been clear that there is a high degree of symmetry between particle and antiparticle. This means, among other things, that a world composed of antimatter would behave in a similar manner to our world. This basic tenet of particle physics, the symmetry between matter and antimatter, is in stark contradiction to the wealth of everyday and cosmological evidence that the universe is composed almost entirely of matter with little or no primordial antimatter. The evidence that the universe is devoid of antimatter comes from a variety of different observations. On the very small scale, the absence of proton-antiproton annihilations in our everyday actions, constitutes strong evidence that our world is composed only of matter and no antimatter. Moving up in scale, the success of satellite launches, lunar landings, and planetary probes indicate that our solar system is made up of the same type of matter that we are, and that there is negligible antimatter on that scale. To determine whether antimatter exists in our galaxy, we turn to cosmic rays. Here we see the first detection of antimatter outside particle accelerators. Mixed in with the many protons present in cosmic rays are a few antiprotons, present at a level of around $10^{-4}$ in comparison with the number of protons (see for example \citeasnoun{SA 88}). However, this number of antiprotons is consistent with their secondary production through accelerator-like processes, $p+p\rightarrow 3p + {\bar p}$, as the cosmic rays stream towards us. Thus there is no evidence for primordial antimatter in our galaxy. Finally, if matter and antimatter galaxies were to coexist in clusters of galaxies, then we would expect there to be a detectable background of $\gamma$-radiation from nucleon-antinucleon annihilations within the clusters. This background is not observed and so we conclude that there is negligible antimatter on the scale of clusters (For a review of the evidence for a baryon asymmetry see \citeasnoun{GS 76}). A natural question to ask is, what is the largest scale on which we can say that there is no antimatter? This question was addressed by \citeasnoun{GS 76},and by \citeasnoun{FS 85}, and in particular has been the subject of a careful recent analysis by \citeasnoun{CRG 97}. If large domains of matter and antimatter exist, then annihilations would take place at the interfaces between them. If the typical size of such a domain was small enough, then the energy released by these annihilations would result in a diffuse $\gamma$-ray background and a distortion of the cosmic microwave radiation, neither of which is observed. Quantitatively, the result obtained by the latter authors is that we may safely conclude that the universe consists entirely of matter on all scales up to the Hubble size. It therefore seems that the universe is fundamentally matter-antimatter asymmetric. The above observations put an experimental upper bound on the amount of antimatter in the universe. However, strict quantitative estimates of the relative abundances of baryonic matter and antimatter may also be obtained from the standard cosmology. Primordial nucleosynthesis (for a review see \citeasnoun{CST 95}) is one of the most powerful tools of the standard cosmological model. The theory allows accurate predictions of the cosmological abundances of all the light elements, H, $^3$He, $^4$He, D, B and $^7$Li, while requiring only a single input parameter. Define $n_b$ to be the number density of baryons in the universe. Similarly define $n_{\bar b}$ to be the number density of antibaryons, and the difference between the two to be $n_B$. Then, if the entropy density in the universe is given by $s$, the single parameter required by nucleosynthesis is the baryon to entropy ratio \be \eta \equiv \frac{n_B}{s} = \frac{n_b-n_{\bar b}}{s} \ , \ee and one may conservatively say that calculations of the primordial light element abundances are correct if \be 1.5\times 10^{-10} < \eta < 7\times 10^{-10} \ . \label{nucleo} \ee Although the range of $\eta$ within which all light element abundances agree with observations is quite narrow (see figure \ref{bbnfig}), its existence at all is remarkable, and constitutes a strong confirmation of the standard cosmology. For recent progress in nucleosynthesis see \citeasnoun{TFB 96} and \citeasnoun{CH 97}. The standard cosmological model provides a complete and accurate description of the evolution of the universe from extremely early times (a few minutes) to the present day ($10-20$ billion years) given a host of initial conditions, one of which is the value of $\eta$. This standard picture is based on classical, fluid sources for the Einstein equations of General Relativity (GR). While the success of the standard cosmology is encouraging, there remains the question of the initial conditions. One approach is just to consider the initial values of cosmological parameters as given. However, the values required for many parameters are extremely unnatural in the sense that the tiniest deviation from them leads to a universe entirely different from the one that we observe. One well known example of this is the initial value of the mass density of the universe, the naturalness of which is at the root of the {\it flatness problem} of the standard cosmology. The philosophy of {\it modern} cosmology, developed over the last thirty years, is to attempt to explain the required initial conditions on the basis of quantum field theories of elementary particles in the early universe. This approach has allowed us to push our understanding of early universe cosmology back to much earlier times, conservatively as early as $10^{-10}$ seconds, and perhaps much earlier. The generation of the observed value of $\eta$ in this context is referred to as {\it baryogenesis}. A first step is to outline the necessary properties a particle physics theory must possess. These conditions were first identified by \citeasnoun{Sakharov} and are now referred to as the three {\it Sakharov Criteria}. They are: \begin{itemize} \item Violation of the baryon number ($B$) symmetry. \item Violation of the discrete symmetries $C$ (charge conjugation) and $CP$ (the composition of parity and $C$) \item A departure from thermal equilibrium. \end{itemize} The first of these is rather obvious. If no processes ever occur in which $B$ is violated, then the total number of baryons in the universe must remain constant, and therefore no asymmetry can be generated from symmetric initial conditions. The second Sakharov criterion is required because, if $C$ and $CP$ are exact symmetries, then one can prove that the total rate for any process which produces an excess of baryons is equal to the rate of the complementary process which produces an excess of antibaryons and so no net baryon number can be created. That is to say that the thermal average of $B$, which is odd under both $C$ and $CP$, is zero unless those discrete symmetries are violated. Finally, there are many ways to explain the third criterion. One way is to calculate the equilibrium average of $B$: \bea \langle B\rangle_T & = & \hbox{Tr}(e^{-\beta H}B) \nonumber \\ & = & \hbox{Tr}[(CPT)(CPT)^{-1}e^{-\beta H}B)] \nonumber \\ & = & \hbox{Tr}(e^{-\beta H}(CPT)^{-1}B(CPT)] \nonumber \\ & = & -\hbox{Tr}(e^{-\beta H}B) \ , \eea where in the third step I have used the requirement that the Hamiltonian $H$ commutes with $CPT$, and in the last step used the properties of $B$ that it is odd under $C$ and even under $P$ and $T$ symmetries. Thus $\langle B\rangle_T = 0$ in equilibrium and there is no generation of net baryon number. This may be loosely described in the following way. In quantum field theories in thermal equilibrium, the number density of any particle species, $X$ say, depends only on the energy of that species, through \be n_{eq}(X)=\frac{1}{e^{(E-\mu)/T} \pm 1}\ , \ee where $\mu$ is the chemical potential corresponding to baryon number. Since the masses of particle and antiparticle are equal by virtue of the CPT theorem, and $\mu=0$ if baryon number is violated, we have that \be N_{eq}(X)=\int\frac{d^3p}{(2\pi)^3}n_{eq}=N_{eq}({\bar X})\ , \ee and again there is no net asymmetry produced. The focus of this article is to review one popular scenario for generating the baryon asymmetry of the universe (BAU), as quantified in equation~(\ref{nucleo}), within the context of modern cosmology. In general, such scenarios involve calculating $n_B$, and then dividing by the entropy density \be s=\frac{2\pi^2}{45}g_* T^3 \ , \ee where $g_*$ is the effective number of massless degrees of freedom at temperature $T$. While there exist many attempts in the literature to explain the BAU (for a review see \citeasnoun{AD 92}), I will concentrate on those scenarios which involve anomalous electroweak physics, when the universe was at a temperature of $10^2\,$GeV (for earlier reviews see \citeasnoun{NTreview}, \citeasnoun{CKNreview}, and \citeasnoun{review}). The production of the BAU through these models is referred to as {\it electroweak baryogenesis}. In the next section I will describe baryon number violation in the electroweak theory both at zero and at nonzero temperature. In section \ref{CP} I shall move on to the subject of CP violation, explaining how this arises in the standard model and how it is achieved in some popular extensions. Section \ref{EWPT} contains an account of the electroweak phase transition, including a discussion of both analytic and numerical approaches. Having set up the framework for electroweak baryogenesis, I turn in section \ref{local} to the dynamics in the case where baryon production occurs close to a phase boundary during a phase transition. In section \ref{nonlocal}, I then extend these ideas to include the effects of particle transport, or diffusion. Section \ref{MSSM} contains a description of how baryogenesis is implemented in a popular extension of the standard model, the minimal supersymmetric standard model (MSSM). In section \ref{defects}, I explain how, in some extensions of the electroweak theory, baryogenesis may be mediated by topological defects, alleviating the constraints on the order of the phase transition. Finally, in section \ref{conclusions} I summarize the results and comment on open questions and future directions in the field. It is my hope that this article will fulfill its intended role as both a review of the background and basic material for beginners in the field, and as a summary of and commentary on the most recent results and directions in the subject. However, the focus of this article, as in any such endeavor, is quite idiosyncratic, and I apologize to any of my colleagues whose work has been omitted or incorrectly detailed. A different focus can be found in other accounts of the subject and, in particular, for a comprehensive modern review of numerical approaches I recommend \citeasnoun{review}. A note about conventions. Throughout I use a metric with signature $+2$ and, unless explicitly stated otherwise, I employ units such that ${\hbar}=c=k=1$ so that Newton's constant is related to the Planck mass through $G=M_{pl}^{-2}$. | \label{conclusions} Modern particle cosmology consists of the union of the hot big bang model with quantum field theories of elementary particles. Under mild assumptions, it is a consequence of this structure that when the universe was extremely young and hot, the net baryon number of the universe was zero. That is, the number of particles carrying a given baryon number in any region was equal on average to the number of the appropriate antiparticles carrying the opposite baryon number. However, on the other hand, it is a clear observational fact that the universe is maximally baryon - antibaryon asymmetric. This fact is quantified from the considerations of primordial nucleosynthesis. These calculations are perhaps the most impressive success of the standard cosmology, accurately predicting the abundances of the light elements from the single input parameter of the baryon to entropy ratio, which is constrained as in equation~(\ref{nucleo}). Until recently, there were two possible explanations for this. First, the universe as a whole could be baryon number symmetric, but baryon-antibaryon separation could have resulted in an apparent baryon asymmetry in the local universe. the second possibility is that some dynamical process took place as the universe evolved, causing baryons to be preferentially produced over antibaryons. A recent analysis \cite{CRG 97} has ruled out the former possibility over scales up to the size of the observable universe. It therefore appears that the latter option, {\it baryogenesis}, must have taken place. In this article I have tried to describe how a number of different physical effects, all present in the standard electroweak theory at nonzero temperature, can come together in the context of the expanding universe to implement baryogenesis. It is an amazing fact about the Glashow-Salam-Weinberg model and its modest extensions that they satisfy all three Sakharov criteria for producing a baryon excess. As a result, over the last decade, {\it electroweak baryogenesis} has been a very popular scenario for the generation of the BAU. There are four technical issues to be investigated when considering models of electroweak baryogenesis. these are \begin{enumerate} \item How is the departure from equilibrium realized? Is the electroweak phase transition strongly first order or are topological defects necessary? \item How is sufficient CP violation obtained? \item What is the rate of baryon number violating processes? Is this rate high enough in the unbroken phase and can washout of any asymmetry be avoided? \item What are the actual dynamics of baryon number production? \end{enumerate} I hope I have described how each of these issues has been addressed in the literature on electroweak baryogenesis. Detailed analyses of the phase transition, coupled with the smallness of the CP violation due to phases in the CKM matrix, have made it clear that an extension of the standard model is required to make the scenarios viable. While general two-Higgs doublet theories have been considered, perhaps the most appealing candidate from particle physics motivations is the minimal supersymmetric standard model. The implementations of electroweak baryogenesis in this model have recently been investigated and the range of parameters of the model for which an appreciable BAU can be generated have been calculated. This range should be accessible to future particle colliders and thus the scenario of EWBG in the MSSM should be experimentally testable. If the MSSM or two-Higgs models are not chosen by nature, then other models of electroweak baryogenesis may be relevant. If the topological structure of the relevant theory admits gauge solitons, then defect mediated electroweak baryogenesis may be important, irrespective of the order of the phase transition. Although I have argued that there exist viable scenarios of electroweak baryogenesis, particularly in the context of supersymmetric models, there remain a number or open questions and directions for future research. I shall list the ones that I feel are most important. \begin{enumerate} \item At present, the best quantitative understanding of the electroweak phase transition comes from lattice Monte Carlo simulations. While the results from these are impressive, they do not provide an intuitive understanding of the microphysics of the phase transition. The numerical results are partially supported by some analytical approaches but these often cannot be trusted in the physical range of Higgs masses. An analytic understanding of the nonperturbative dynamics of the phase transition would be an important step forward. \item The chemical potential and ``fluid'' approaches to nonlocal baryogenesis provide good analytical tools for understanding the nonlocal production of baryons. However, in the case of local baryogenesis, analytical methods which yield believable quantitative results have yet to be found. Again, it is encouraging that numerical calculations, for example of Chern-Simons number diffusion, are providing quantitative predictions but an appropriate analytical model is very desirable. \item While it will strongly support electroweak baryogenesis if supersymmetry is verified with parameters in the correct range, it is important to ensure that there is sufficient CP violation. To this end, phenomenological predictions of CP violation in SUSY models and the corresponding experimental tests are crucial hurdles that the scenario must pass. \end{enumerate} If electroweak baryogenesis is correct, then no matter what the relevant electroweak model is, the physics involved is extremely beautiful and diverse. The topology of gauge field theories, the physics of phase transitions, CP violation, plasma dynamics and thermal field theory all play a part in generating the BAU. However, perhaps the most attractive feature of electroweak baryogenesis scenarios is that they should be testable at the next generation of particle colliders. It is an exciting possibility that we may soon understand the origin of the matter that makes up our universe. | 98 | 3 | hep-ph9803479_arXiv.txt |
9803 | gr-qc9803084_arXiv.txt | We develop the general theory of stars in Saa's model of gravity with propagating torsion and study the basic stationary state of neutron star. Our numerical results show that the torsion force decreases the role of the gravity in the star configuration leading to significant changes in the neutron star masses depending on the equation of state of star matter. The inconsistency of the Saa's model with Roll-Krotkov-Dicke and Braginsky-Panov experiments is discussed. \noindent{PACS number(s): 04.40.Dg,04.40.-b,04.50.+h} | In recent years the interest in scalar-tensor theories of gravity has been renewed. One reason for this is the important role which these theories play in the understanding of inflantionary epoch. On the other hand the scalar-tensor gravitation (the so called "dilaton gravity") arises naturally from the low-energy limit of the super-string theory \cite{Witten}, \cite{DT}. The predictions of scalar-tensor theories may differ drastically from these of general relativity. For example such a phenomenon -- "spontaneous scalarization" was recently discovered by Damour and Esposito-Farese as a non-perturbative strong field effect in a massive neutron star \cite{Damour}. Other interesting phenomenon is "gravitational memory" of black holes proposed in \cite{Barrow}. The "gravitational memory" in the case of boson stars was investigated in \cite{TLS} (see also \cite{CS}).Their stability through cosmic history using catastrophe theory was investigated in \cite{TSL}. Many theories of gravity with propagating torsion involving a scalar field have been proposed in the last decades, too \cite{HRRS}, \cite{HRR}, \cite{SG}. In such theories contrary to the usual Einstein-Cartan gravity \cite{Hehl1}-\cite{Hehl3}, there are long-range torsion mediated interactions. Carrol and Field \cite{Carrol} have examined some observational consequences of propagating torsion in a wide class of models involving a scalar field. They conclude that for reasonable models the torsion could be detected experimentally. Recently a new interesting model with propagating torsion was proposed by Saa \cite{Saa1}-\cite{Saa5}. This model involves a non-minimally coupled scalar field as a potential of the torsion of space-time. As one can see Saa's model is very close to the dilaton gravity. In the present article we investigate both analytically and numerically a neutron star in the Saa's model and compare obtained results with these in the general relativity. We also discuss new predictions of the theory under consideration. The paper is organized as follows. In section 2 we consider briefly Saa's model. In section 3 we give the necessary information for the vacuum solutions of the field equations. The equations determining static equilibrium solutions for a neutron star are discussed in section 4. Numerical results for the neutron star are discussed in section 5. The stability of the neutron star is discussed via catastrophe theory in section 6. The inconsistency of the Saa's model with Roll-Krotkov-Dicke and Braginsky-Panov experiments is discussed in section 7. | We have solved the system of equations (\ref{NF}) coupled with the state equations (\ref{EOS1}) and (\ref{EOS2}) numerically using the method due to Runge-Kutta-Merson with automatic error control. The results are shown in the corresponding figures.\\ Hereafter all masses are measured in units $M_{\bigodot}$. \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz01a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz01b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{a) $M -log(\varepsilon_{c})$ dependence. \hskip 1.8truecm b) $M_{T} - M_{R}$ dependence.\hskip .7truecm} \vspace{.5truecm} \label{Fig1} \end{figure} First we concentrate our attention on the case of non-interacting neutron gas. In Fig. 1a) the dependence of the three masses $M_{T}$,$M_{K}$,$M_{R}$ on the central density $\varepsilon_{c}$ is shown. The appearance of a cusp in Fig. 1b) , where the dependence of the tensor mass $M_{T}$ on the rest mass $M_{R}$ is presented, shows that their maxima lie at the same point. Although, the maxima of the rest mass $M_{R}$ and Keplerian mass $M_{K}$ are too close in Fig. 1a) they don't lie at the same point, as it may be seen from Fig. 2a) which shows the dependence of $M_{K}$ on $M_{R}$. We see also that Keplerian mass is considerably greater than the tensor one -- about three times. In Fig. 3a) the $M_{T} - R$ dependence is represented. It's seen that the $M_{T} - R$ curve in our case is fairly similar to the one of general relativity, but there are significant differences, too. The maximum mass ${M_{T}}_{max} $ in our case is $ \approx 0.35M_{\bigodot}$, while in general relativity the Oppenheimer-Volkoff's mass is $M_{OV} = 0.7M_{\bigodot}$. The radius corresponding to the mass $M_{\bigodot}$ is $R = 4.2 km $, while in the case of general relativity $ R = 9.6 km $. If we look at Fig. 2b) where the dependence of $M_{T}$ on the central density ${\varepsilon}_{c}$ is shown, we note that ${M_{T}}_{max}$ lies at ${\varepsilon}_{c} \approx 4.5 * 10^{16} g/{cm}^3$, while $M_{OV}$ lies at ${\varepsilon}_{c} \approx 5 * 10^{15} g/{cm}^3$ in general relativity. The average density in our case is about $ 4 $ times greater than the one in general relativity. Hence, in the model under consideration the neutron star is more compact and has a mass about $1/2$ - times smaller than $M_{OV}$. In Fig. 3b) the dependence of Keplerian mass on the star radius is presented. It's seen that the Keplerian mass is about $1.5$ times greater than $M_{OV}$. \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz02a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz02b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{a) $M_{Kepler}-M_{R}$ dependence. \hskip 2truecm b) $M_{T} -log(\varepsilon_{c})$ dependence.\hskip .5truecm} \vspace{.5truecm} \label{Fig2} \end{figure} \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz03a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz03b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{a) $M_{T}-R$ dependence. \hskip 2.4truecm b) $M_{Kepler}-R$ dependence.\hskip 1.1truecm } \vspace{.5truecm} \label{Fig3} \end{figure} \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz04a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz04b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{a) $k- r$ dependence. \hskip 2truecm b) $K -log(\varepsilon_{c})$ dependence.\hskip .5truecm} \vspace{.5truecm} \label{Fig2_} \end{figure} In the Fig. 4a) the dependence $k(r)$ is shown inside the star (for central density $7.5*10^{15}g/{cm}^3 $). In accordance to the general considerations $k$ increases from the center of the star to the surface, where $k$ takes a value $K=k(R) \approx 0.45 - 0.46$, which is close to $0.5$. The dependence $K(\varepsilon_{c})$ of $K$ on the star central density $\varepsilon_{c}$ is shown in Fig. 4b). It's seen that $K$ decreases when density increases, which is similar to the previous case. So, the ratio of the torsion force to the gravitational one takes its minimum value at the center of the star and is the greatest at the surface. \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz05a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz05b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{a) $K-M_{T}$ dependence. \hskip 2.4truecm b) $K-R$ dependence.\hskip 1.1truecm } \vspace{.5truecm} \label{Fig3_} \end{figure} As it may be seen from Fig. 5a) expressing the dependence $K(M_{T})$, the torsion-urged effects are relatively strongest in the case of small masses -- with increasing of the star mass (up to the point where the star loses its stability) $K$ decreases. It's seen from Fig. 5b), where the dependence of $K$ on the star radius $R$ is shown, that in the area of stability $K$ decreases when $R$ decreases too -- the more compact stars are, the smaller $K$ they have. Fig. 6a) presents the dependencies $\theta(r)$ and $\nu(r)$ inside the star. One may see that ${1 \over 2}\nu - 3\Theta < 0$ everywhere. The dependencies $m_{T}(r)$,$m_{Kepler}(r)$ and $m_{\theta}(r)$ are shown in Fig. 6b) for central density $7.5*10^{15}$. As it has already mentioned all masses increase with $r$. \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz06a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz06b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{a) $\Theta, \nu -r$ dependence. \hskip 1.9truecm b)$m_{T},m_{Kepler},m_{\theta} -r$ dependence.\hskip .7truecm } \vspace{.5truecm} \label{Fig4} \end{figure} The following figures illustrate the case of Tsuruta-Cameron equation of state (TCES). \vskip 2truecm \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz07a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz07b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{TCES. a) $M -log(\varepsilon_{c})$ dependence . \hskip 1.2truecm b) $M_{T} -M_{R}$ dependence.\hskip 1.2truecm } \vspace{.5truecm} \label{Fig5} \end{figure} We see from the figures that the maximum tensor mass in this case is about $2M_{\bigodot}$ and the corresponding radius is about $7.5km$ - the same quantities in general relativity are correspondingly $\approx 1.6M_{\bigodot}$ and $\approx 11.5km$. Hence, the interaction between the nucleons leads to an increase in the maximum mass, as in general relativity. Note the differences between the Fig. 5b) and Fig. 6b) (for the case of non-interacting neutron gas), and the corresponding Fig. 11b) and Fig. 12b) (for the case of Tsuruta-Cameron equation of state). There one can see the strong dependence of some results in the Saa's model of gravity with propagating torsion on the equation of state of star's matter. \vskip 2truecm \begin{figure}[htbp] \vspace{3.5truecm} \special{psfile=tfiz08a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz08b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{TCES. a)$M_{Kepler}-M_{R}$ dependence. \hskip .4truecm b) $M_{T}-log(\varepsilon_{c}) $ dependence.\hskip 1truecm} \vspace{.5truecm} \label{Fig6} \end{figure} \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz09a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz09b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{TCES. a) $M_{T}-R$ dependence. \hskip 1.6truecm b) $M_{Kepler}-R$ dependence.\hskip 1.7truecm} \vspace{.5truecm} \label{Fig7} \end{figure} \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz10a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz10b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{TCES. a) $k-r$ dependence. \hskip 1.3truecm b) $K -log(\varepsilon_{c}) $ dependence.\hskip 3truecm} \vspace{.5truecm} \label{Fig8} \end{figure} \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz11a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz11b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{TCES. a) $K-M_{T}$ dependence. \hskip 1.3truecm b) $K-R $ dependence.\hskip 3truecm} \vspace{.5truecm} \label{Fig8_} \end{figure} \begin{figure}[htbp] \vspace{4truecm} \special{psfile=tfiz12a.eps hoffset=-15 voffset=-160 hscale=30 vscale=30} \special{psfile=tfiz12b.eps hoffset=200 voffset=-160 hscale=30 vscale=30} \vskip 0.5truecm \caption{TCES. a) $\Theta, \nu-r$ dependence. \hskip 1.3truecm b) $m_{T},m_{Kepler}, m_{\theta}-r$ dependence.\hskip 3truecm} \vspace{.5truecm} \label{Fig8+} \end{figure} We have also examined the Harrison-Wheeler's equation of state \cite{HTWW}. As in general relativity the numerical results are very close to these for the noninteracting neutron gas. For example the maximum tensor and Keplerian mass is correspondingly $\approx 0.35M_{\bigodot}$ and $\approx 1M_{\bigodot}$, and the corresponding radius is $3.8 km$. Other equations of state (of politropic type) have been examined, too. The corresponding maximum tensor mass of a neutron star reaches a value about $2.5-2.6M_{\bigodot}$, while the corresponding Keplerian mass is about $6-6.5M_{\bigodot}$. As it is seen from numerical calculations the tensor mass and Keplerian mass differ very significantly from each other. This behaviour of the masses is qualitatively the same as in the case of a boson star in Brans-Dicke theory with $\omega=-1$ \cite{Whinnett}. Saa's model corresponds to the value of the Brans-Dicke parameter $\omega=-{4\over 3}$ which is close to $-1$. That's why the observed qualitative agreement is natural. Hence, including a scalar field with an approximately the same $\omega$ in physically different kinds of stars we find similar departures of the corresponding predictions of general relativity. This conclusion holds also in the case of large $\omega$. For example, in \cite{TLS} and \cite{TSL} Brans-Dicke boson star has been considered with $\omega = 400$. There the results are quite similar to corresponding ones in general relativity which is just the limit $\omega \to \infty$. The star is a bit lighter but has a higher density in future. In this article we have examined the basic spherically symmetric stationary state of stars in the Saa's model of gravity with propagating torsion. In the model under investigation there is no need to consider unknown charges creating the torsion-dilaton field. Its source is the very spinless matter. The whole geometry of the space-time (including metric and torsion) is determined by the familiar properties of this matter. The parameters of the vacuum solution are determined only by the spinless matter without adopting an existence of new properties, too. In contrast to the corresponding models in the general relativity here we have two parameters $K$ and $a$ of the vacuum solutions. The values of these parameters depend on the mass distribution in the star which is related with the equation of state of the star's matter. For a fixed equation of state both parameters become functions only of the star mass, but these functions are not the same for the different equations of state. The first parameter $K$ being the ratio of the magnitude of torsion-dilaton force and of the magnitude of gravitational force for realistic equations of matter state takes values in the interval $[{1\over 3},{1\over 2}]$, depending on the star's mass. The second one -- $a$ is analogous to the gravitational radius in general relativity and takes positive values depending on the value of the parameter $K$ and on the value of the star's mass. To be specific in the present article we restrict our attention to the model of neutron stars where the effects of nonlinearity are essential as in general relativity. Numerical results and analytical considerations show that the space-time torsion may have a significant role in their structure. The new torsion force decrease in some extent the role of the gravity in the star configuration and may lead to an increasing or decreasing of the maximum neutron star mass depending on the equation of state. The complete investigation of the consistence of the whole Saa's model of gravity with propagating torsion (including all type of physical fields) with the reality is still an open problem. The results of the present article may have not only independent value, but are necessary for reaching the solution of this critical problem. For example, after the first version of the present article was send for publication a new results based on it which show the inconsistency of Saa's model with solar system gravitational experiments were found and published independently \cite{FY}. Saa's model is a simple model based on pure geometrical reasons which allows to overcome a basic difficulties of the old models of gravity with torsion: the inconsistency of the application of the minimal coupling principle in action principle and directly in the equations of motion \cite{Saa1}. Moreover, it leads to propagating torsion which is another important physical property of this model. Unfortunately, it is not compatible with basic physical experiments. This means that one has to modify this model preserving its important new physical properties in a proper way to comply with real physics. A new investigations in this direction are in progress (see for example \cite{F2}, \cite{F3}). Nevertheless the model under consideration contradicts to basic experiments its detailed investigation is quite instructive because Saa's model is a special case of scalar-tensor theory of gravity with a nonminimal coupling of the scalar torsion-dilaton field with matter. A similar, but more general coupling of string-dilaton can be expected in the string theory and it leads different physical effects. For example, in the article by Damour and Polyakov \cite{DP} as a consequences of nonminimal coupling between matter and dilaton was derived a violation of the equivalence principle, as far as some string-dilaton effects in the early universe \cite{DP}. Unfortunately, the present days string theory is not able to predict definitely the form of the coupling between string-dilaton and the real matter. In contrast, Saa's model is the only one we know with complete determined interaction of the (torsion) dilaton with all kinds of matter. This interaction is similar to the one expected in other models. In the present article it is shown at first that the nonminimal coupling of the dilaton field with the matter will emerge in extreme conditions in a neutron star and that it leads to clear new physical phenomena of different kind, some details of which may depend on the equation of matter state. Most probably similar effects will appear in other possible modifications of the theory of torsion-dilaton, as far as in the other types of theories of dilaton. For example the neutron stars may give us a way to a real physics in string theories if the string-dilaton interactions with real matter will be established. Hence, the most important conclusion of present article is that looking for physical manifestation of the dilaton one has to investigate in details its influence on the neutron star structure. \bigskip \bigskip \bigskip \noindent{\Large\bf Acknowledgments} \bigskip \bigskip \bigskip We are deeply grateful to the unknown referees who suggested to consider different types of masses of the neutron star in Saa's model and stability analysis of the neutron star via catastrophe theory, and to add to the present paper results about the consistence of this model with physical reality, as far as for pointing out the references \cite{DT}, \cite{TLS}, \cite{CS}, \cite{TSL}, \cite{HRR}, \cite{SG}, \cite{Whinnett}, \cite{KMS1} and \cite{KMS2}. The work on this article has been partially supported by the Sofia University Foundation for Scientific Researches, Contracts~No.No.~245/97,~257/97, and by the Bulgarian National Foundation for Scientific Researches, Contract~No.~F610/97. One of us (PF) is grateful to the leadership of the Bogoliubov Laboratory of Theoretical Physics, JINR, Dubna, Russia for hospitality and working conditions during his stay there in the summer of 1998 when a part of this investigation has been completed. \bigskip \bigskip \bigskip | 98 | 3 | gr-qc9803084_arXiv.txt |
9803 | astro-ph9803185_arXiv.txt | We report the results of a systematic study of the vacuum- ultraviolet ($\lambda \simeq$ 1150 to 2000 \AA) spectra of a sample of 45 starburst and related galaxies observed with the IUE satellite. These span broad ranges in metallicity (from 0.02 to 3 times solar), bolometric luminosity ($\sim 10^{7}$ to $4 \times 10^{11} L_{\odot}$), and galaxy properties (e.g. including low-mass dwarf galaxies, normal disk galaxies, and massive galactic mergers). The projected size of the IUE spectroscopic aperture is typically one to several kpc and therefore usually encompasses the entire starburst and is similar to the aperture-sizes used for spectroscopy of high-redshift galaxies. Our principal conclusion is that local starbursts occupy a very small fractional volume in the multi- dimensional manifold defined by such fundamental parameters as the extinction, metallicity, and vacuum-UV line strengths (both stellar and interstellar) of the starburst and the rotation speed (mass) and absolute magnitude of the starburst's `host' galaxy. More metal-rich starbursts are redder and more heavily extinguished in the UV, more luminous, have stronger vacuum- UV lines, and occur in more massive and optically-brighter host galaxies. We advocate using these local starbursts as a `training set' to learn how to better interpret the rest-frame UV spectra of star-forming galaxies at high-redshift, and stress that the degree of similarity between local starbursts and high-redshift galaxies in this multi-dimensional parameter space can already be tested empirically. The results on local starbursts suggest that the high- redshift `Lyman Drop-Out' galaxies are typically highly reddened and extinguished by dust (average factor of 5 to 10 in the UV), may have moderately high metallicities (0.1 to 1 times solar?), are probably building galaxies with stellar surface-mass-densities similar to present-day ellipticals, and may be suffering substantial losses of metal-enriched gas that can `pollute' the inter-galactic medium. | Starbursts are sites of intense star-formation that occur in the `circum-nuclear' (kpc-scale) regions of galaxies, and dominate the integrated emission from the `host' galaxy (cf. Leitherer et al 1991). The implied star-formation rates are so high that the existing gas supply may sustain the starburst for only a small fraction of a Hubble time (in agreement with detailed models of the observed properties of starbursts, which imply typical burst ages of-order 10$^{8}$ years). Both optical objective prism searches and the IRAS survey have shown that starbursts are major components of the local universe (cf. Huchra 1977; Gallego et al et al 1995; Soifer et al 1987). Indeed, integrated over the local universe, the total rate of (high-mass) star-formation in circumnuclear starbursts is comparable to the rate in the disks of spiral galaxies (Heckman 1997; Gallego et al 1995; Tresse \& Maddux 1998). Thus, starbursts deserve to be understood in their own right. Starbursts are even more important when placed in the broader context of contemporary stellar and extragalactic astrophysics. The cosmological relevance of starbursts has been dramatically underscored by one of the most spectacular discoveries in years: the existence of a population of high-redshift (z $>$ 2) star-forming field galaxies (cf. Steidel et al 1996; Lowenthal et al 1997). The sheer number density of these galaxies implies that they almost certainly represent precursors of typical present-day galaxies in an early actively-star-forming phase. This discovery therefore moves the study of the star-forming history of the universe into the arena of direct observations (Madau et al 1996), and gives added impetus to the quest to understand local starbursts. Observations in the vacuum-UV spectral regime are crucial for both understanding local starbursts, and for relating them to galaxies at high-redshift. This is the spectral regime where we can most clearly observe the direct spectroscopic signatures of the hot stars that provide most of the bolometric luminosity of starbursts (e.g. Sekiguchi \& Anderson 1987; Fanelli, O'Connell, \& Thuan 1988; Leitherer, Robert, \& Heckman 1995). Moreover, the vacuum-UV contains a wealth of spectral features, including the resonance transitions of most cosmically-abundant ionic species (cf. Morton 1991). These give UV spectroscopy a unique capability for diagnosing the (hot) stellar population and the physical, chemical, and dynamical state of gas in starbursts. Since ground-based optical observations of galaxies at high-redshifts sample the vacuum-UV portion of their rest-frame spectrum, we cannot understand how galaxies evolved without documenting the vacuum-UV properties of galaxies in the present epoch. In particular, a thorough understanding of how to exploit the diagnostic power of the rest-frame UV spectral properties of local starbursts will give astronomers powerful tools with which to study star-formation and galaxy-evolution in the early universe. Accordingly, we have undertaken an analysis of the vacuum-UV spectroscopic properties of a sample of 45 starburst and related galaxies in the local universe using the IUE archives. In section 2, we describe our analysis of these data. In section 3 we use the data to document the empirical vacuum-UV spectroscopic properties of local starbursts and to assess the dependence of these UV properties on crucial parameters like the dust content, metallicity, and luminosity of the starburst, and the mass and luminosity of the `host galaxy'. We will also describe how published analyses of HST UV spectra allow us to better understand the empirical results from the IUE data, and in particular to ascertain the origin of the absorption features seen in the IUE data (stellar vs. interstellar). In section 4, we will interpret the results from section 3, use these results to elucidate some of the properties of starbursts at high redshift, and point-out some potentially far-reaching implications. | The principal conclusion of this paper is that local starbursts occupy a very small fractional volume in the multi-dimensional manifold defined by such fundamental parameters as the extinction, metallicity, and vacuum-UV line strengths (both stellar and interstellar) of the starburst and the rotation speed (mass) and absolute magnitude of the starburst's `host' galaxy. More metal-rich starbursts are redder, more heavily extinguished (IR-dominated), have stronger vacuum-UV lines, occur in more massive and brighter host galaxies, and can be more luminous. Starbursts with solar or higher metallicity are heavily extinguished indeed: only 1 to 10\% of the vacuum-UV radiation escapes and the resulting UV spectral energy distribution is nearly flat in F$_{\lambda}$ (Fig. 2). Only the most metal-poor starbursts ($<$ 10\% solar) are relatively unaffected by dust (Fig. 2). These correlations are not surprising, since we expect the dust-to-gas ratio in the starburst's ISM to scale roughly with metallicity. In principle, relatively red vacuum-UV spectral energy distributions could be due to the effects of the age of the stellar population or the form of the initial mass function (IMF) rather than dust. That is, the redder spectra might correspond to older starbursts, starbursts whose UV spectrum was more heavily contaminated by an older underlying stellar population, or starbursts with a steeper (less O-star-enriched) IMF. However, the strong correlation between the UV color ($\beta$) and the strengths of the CIV and SiIV (predominantly) stellar wind lines (Fig. 5b) shows that this cannot be the case: an age or IMF effect would require the stellar wind features due to O stars to be {\it weaker} in the redder starbursts - exactly the opposite of what we observe. The metal-poor starbursts have systematically weaker vacuum-UV absorption-lines (Fig. 1). In the case of the stellar wind lines (Fig. 3), this result agrees with expectations based on both theory and observations of metal-poor, hot, high-mass stars. In the case of the interstellar lines (Figure 4), this probably reflects both higher average ionic column densities and larger velocity dispersions in the more metal-rich starbursts. That is, even though the strong interstellar lines are optically-thick (and their strength will thus be set to first-order by the velocity dispersion), there will still be a weak dependence of equivalent width on column density. Moreover, the more metal-rich starbursts can also be more powerful (Fig. 8) and are situated in more massive galaxies (Fig. 11). We expect them to have higher average ISM velocity dispersions, since both gravity and the mechanical `stirring' provided by supernovae and stellar winds will contribute to the gas dynamics. The trend for the more metal-rich starbursts to occur in brighter and more massive host galaxies (Fig. 11 and 12) presumably reflects the well-known mass-metallicity relation defined by normal galaxies (e.g. Zaritsky, Kennicutt, \& Huchra 1994). That is, the ISM in massive galaxies will already be pre-enriched to relatively high metallicity before the starburst is even initiated. During the starburst itself, the fraction of the newly-synthesized metals that are retained (rather than blown out) should also be a strong function of the depth of the gravitational potential well (e.g. Dekel \& Silk 1986). The correlation between the metallicity and the luminosity of the starburst (even though it is due in part to selection effects - see section 3.4) may reflect other fundamental trends: more massive galaxies can host more powerful starbursts (Fig. 11) and have more metal-rich interstellar gas with which to fuel the starburst. The luminosity {\it vs.} rotation-speed connection in starbursts has been previously discussed by Heckman (1993), Lehnert \& Heckman (1996b), and Meurer et al (1997): basic considerations of causality in a self-gravitating system demand that the maximum star-formation rate will scale as the cube of the rotation speed. The above results are not only very illuminating regarding the nature of the starburst phenomenon, they also have a variety of interesting implications for the interpretation of the rest-frame-UV properties of galaxies at high-redshift. First and foremost, starbursts in the present universe emit most of their light in the far-infrared, not in the ultraviolet (Fig. 9). Thus, an ultraviolet census of the local universe would significantly underestimate the true high-mass star-formation-rate and would systematically under-represent the most powerful, most metal- rich starbursts (Fig. 2), occuring in the most massive galaxies (Fig. 11). This {\it may} also be true at high-redshift, where the current estimates of star-formation rely almost exclusively on data pertaining to the rest-frame vacuum-UV (e.g. Madau et al 1996, 1998). If there {\it is} a dependence of dust extinction on luminosity at high-redshift, this will affect the apparent shape of the luminosity function. For example, current samples at high-redshift might under-represent young/forming massive elliptical galaxies. These would have high metallicities (and thus a high dust-to-gas ratio) coupled with very large average gas surface mass densities ($\sim 10^{22} - 10^{23}$ cm$^{-2}$), and would hence have large dust opacities. Using the strong correlation between the vacuum-UV color of local starbursts ($\beta$) and the ratio of far-IR to vacuum-UV light emitted by local starbursts, Meurer et al (1998) estimate that an average vacuum-UV-selected galaxy at high-redshift (e.g. Steidel et al 1996; Lowenthal et al 1997) suffers about 2 to 3 magnitudes of UV extinction (in agreement with modeling of the rest-frame UV-through-visible spectral energy distributions by Sawicki \& Yee 1998). The `correct' prescription for de-extincting the high- redshift galaxies in order to correctly obtain the bolometric luminosity and star-formation rate is a matter of on-going debate (see Pettini 1997 and Madau et al 1996, 1998 for other viewpoints). For example, one uncertainty at high-redshift is the unknown relative importance of the age or the IMF of the stellar population {\it vs.} dust in producing the observed UV spectral energy distribution. As we have argued above, the trend for the {\it redder} starbursts to have {\it stronger} absorption features due to massive stars (Fig. 5b) rules out age or the IMF as the primary determinant of UV color in our sample. Thus, it will be extremely interesting to see if the same correlation holds in the high-redshift galaxies. In any case, it seems fair to conclude that the history of high-mass star-formation in the universe at early times will remain uncertain until the effects of dust extinction are better understood. The extinction corrections advocated by Meurer et al (1997) imply large bolometric luminosities for the high-redshift galaxies ($\sim$ 10$^{11}$ to 10$^{13}$ L$_{\odot}$ for H$_{0}$ = 75 km s$^{-1}$ Mpc$^{-1}$ and q$_{0}$ = 0.1). Interestingly, the bolometric surface-brightnesses of the extinction-corrected high-redshift galaxies are then very similar to the values seen in local starbursts: $\sim$ 10$^{10}$ to 10$^{11}$ L$_{\odot}$ kpc$^{-2}$ (Meurer et al 1997). The high-redshift galaxies thus appear to be `scaled-up' (larger and more luminous) versions of the local starbursts. The physics behind this `characteristic' surface-brightness is unclear (cf. Meurer et al 1997; Lehnert \& Heckman 1996a). However, it is intruiging that the implied average surface-mass-density of the stars within the half- light radius ($\sim$ 10$^{2}$ to 10$^{3}$ M$_{\odot}$ pc$^{-2}$) is quite similar to the values in present-day elliptical galaxies. Are we witnessing the formation of elliptical galaxies and/or bulges, or only the formation of precursor fragments thereof (e.g. Sawicki \& Yee 1998; Trager et al 1997)? If the high-redshift galaxies are indeed scaled-up versions of local starbursts, we might expect that they would follow the same trends summarized above. For example, do the high- redshift galaxies should show the same strong correlation between the strength of the UV absorption-lines (stellar and interstellar) and $\beta$ (Fig. 5), whereby the more metal-rich local starbursts are both redder and stronger-lined? As more near-IR (rest-frame visible) data become available, will the high-redshift galaxies obey the starburst correlations (Fig. 11 and 12) of both vacuum-UV color and absorption-lines strengths with M$_{B}$ and rotation speed (do they follow a mass-metallicity relation?). When the high-redshift galaxies are extinction-corrected, are the more luminous systems redder and stronger-lined in the UV (as in the case of the local starbursts - Fig. 9 and 10 respectively)? Finally, do the strengths of the low-ionization (primarily interstellar) and high-ionization (significant stellar wind contribution) absorption-lines correlate well with one-another, as in the local starbursts (Fig. 7)? If the answers to these questions are `yes' (that is, if the high-redshift galaxies and local starbursts do occupy the same parts of the multi-dimensional parameter space summarized above), we might be able to use the UV properties of the high- redshift galaxies to `guesstimate' their metallicities. For example, the strong correlation between vacuum-UV color ($\beta$) and metallicity in local starbursts (Fig. 2) - if applied naively to high- redshift galaxies - would suggest a broad range in metallicity from substantially subsolar to solar or higher and a median value of perhaps 0.3 solar. This is somewhat higher than the mean metallicity in the damped Ly$\alpha$ systems (Pettini et al 1997; Vladilo 1997; Lu et al 1997), but this may be due to selection effects: the UV-selected galaxies are the most actively star-forming regions of galaxies, while the damped Ly$\alpha$ systems tend to sample the outer, less-chemically-enriched parts of galaxies (e.g. Prochaska \& Wolfe 1997) or perhaps proto-galactic fragments or dwarf galaxies (e.g. Vladilo 1997). Our `guesstimated' metallicities of the high-redshift galaxies are much higher than those suggested by Trager et al (1997), who argue that we are witnessing the formation of Population II spheroidal systems. It is also important to sound a cautionary note: most of the strong absorption-lines in the spectra of local starbursts are of {\it interstellar} rather than {\it stellar} origin. This is probably generically true for the low-ionization resonance lines (although B stars may make a non-negligible contribution). Even the high-ionization resonance lines like CIV and SiIV can contain strong interstellar contributions. In fact, in extreme cases (e.g. NGC 1705) even these lines are dominated by interstellar gas. This highlights the difficulty in using the UV spectra of galaxies to deduce their stellar content: data of relatively high spectral resolution and signal-to-noise are needed to reliably isolate the stellar and interstellar components. Simply measuring the equivalent widths of the lines is not enough: it is the {\it profile shapes} that contain the key information. The stellar photospheric absorption-lines identified in Fig. 1 may be the most straightforward features for the detection of stars in high-redshift galaxies, since they are not resonance lines, and so are uncontaminated by interstellar absorption. Unfortunately, most of these lines are weak. The relatively strong CIII$\lambda$1175 line is lost in the Ly$\alpha$ forest at high-redshift, and the CIII$\lambda$1892, FeIII$\lambda$1925, and FeIII$\lambda$1960 lines are often `clobbered' by the OH night-sky lines in galaxies at z $>$ 2.6. Finally - based on local starbursts - it seems likely that the gas kinematics that are measured in the high-redshift galaxies using the interstellar absorption-lines are telling us a great deal about the hydrodynamical consequences of high-mass star-formation on the interstellar medium, and thus it will be difficult to use them to straightforwardly glean information about the gravitational potential or mass of the galaxy. Even the widths of the optical nebular emission-lines in local starbursts are not always reliable tracers of the galaxy potential well (cf. Lehnert \& Heckman 1996a; but see Melnick, Terlevich, \& Moles 1988). This means that it may be tricky to determine masses for the high-redshift galaxies without measuring real rotation-curves via spatially-resolved spectroscopy. On the brighter side, there is now rather direct observational evidence that the kinematics of the interstellar absorption-lines in the high-redshift galaxies reflect an outflowing metal-enriched gas, thereby allowing us to directly study the `pollution' of the inter-galactic medium with metals in the early universe. The signature of these outflows is interstellar absorption lines that are systematically blueshifted by several hundred km s$^{-1}$ relative to the Ly$\alpha$ emission-line (Franx et al 1997; Heckman 1997; Pettini 1997). In local starbursts, this is due to outflowing gas that both produces the blue-shifted absorption-lines and absorbs-away the blue side of the Ly$\alpha$ emission-line (Lequeux et al 1995; Gonzalez-Delgado et al 1998a,b; Kunth et al 1998). The case for an outflow in the high-redshift galaxies could be strengthened by using the weak stellar photospheric lines (Fig. 1) to unambiguously determine the galaxy systemic velocity - cf. Heckman \& Leitherer (1997) and Sahu \& Blades (1997). If these outflows escape the galactic potential well, we should see their cumulative effect in the form of a metal-enriched inter-galactic medium (IGM). The most dramatic such evidence is the existence of an IGM in rich clusters of galaxies whose metal content probably exceeds that of the cluster galaxies themselves. Recent ASCA X-ray spectra of hot gas in clusters of galaxies show approximately solar abundance ratios for the $\alpha$-process elements like O, Ne, and Si relative to Fe (Mushotzky et al 1996; Ishimaru \& Arimoto 1997). This implicates `core-collapse' supernovae (the end product of high-mass stars) and - by inference - starburst-driven superwinds as the source of at least 90\% of the metals in the cluster IGM (Renzini 1997; Gibson, Loewenstein, \& Mushotzky 1997). More generally, the presence of metals in the `Ly$\alpha$ forest' at high-redshift (the cloudy component of the early IGM) is certainly suggestive of the dispersal of chemically- enriched material by early superwinds (cf. Cowie et al 1995; Madau \& Shull 1996). The Ly$\alpha$ forest clouds appear to have a high ratio of Si/C, suggestive of core-collapse supernovae as the source (Cowie et al 1995; Giroux \& Shull 1997). In conclusion, vacuum-ultraviolet observations give us a rich array of diagnostic probes of the stellar and interstellar components of starbursts in the local universe. The analysis described in this paper strongly suggests that these starbursts obey well-defined relationships between such key physical parameters as extinction, metallicity, starburst luminosity and gas dynamics, and host galaxy mass. These objects can therefore serve as local laboratories in which we can study the processes that were important in the galaxy- building events that can now be directly observed at high-redshift. The results of this synergistic study of the most actively-star- forming local and distant galaxies suggests that the latter have rather high UV extinctions (average factor 5 to 10), may have moderately high metallicities (typically 0.1 to 1 times solar), are probably building galaxies having surface mass-densities comparable to present-day ellipticals, and are suffering substantial losses of metal-enriched outflowing gas that can explain the observed `pollution' of the inter-galactic and intra-cluster medium. {\bf Acknowledgments} We would like to thank Anne Kinney and Ralph Bohlin for helping us access and understand the IUE spectra that form the basis of this paper. We also thank Gerhardt Meurer and Daniela Calzetti for many enlightening discussions about the properties of local starbursts, and ditto Mark Dickinson regarding the high-redshift galaxies. This research was primarily supported by NASA LTSA grant NAGW-3138. DG acknowledges the support of the NASA LTSA grant NAG5-6416. This research has made use of the NASA/IPAC extragalactic database (NED), which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the National Aeronautics and Space Administration. | 98 | 3 | astro-ph9803185_arXiv.txt |
9803 | astro-ph9803070_arXiv.txt | A phase transition in the nature of matter in the core of a neutron star, such as quark deconfinement or Bose condensation, can cause the spontaneous spin-up of a solitary millisecond pulsar. The spin-up epoch for our model lasts for $2\times 10^7$ years or 1/50 of the spin-down time (Glendenning, Pei and Weber 1997). | Neutron stars have a high enough interior density as to make phase transitions in the nature of nuclear matter a distinct possibility. According to the QCD property of asymptotic freedom, the most plausible is the quark deconfinement transition. According to lattice QCD simulations, this phase transition is expected to occur in very hot ($T\sim 200$ MeV) or cold but dense matter. Since neutron stars are born with almost the highest density that they will have in their lifetime, being very little deformed by centrifugal forces, they will possess cores of the high density phase essentially from birth if the critical density falls in the range of neutron stars. However the global properties such as mass or size of a slowly rotating neutron star are little effected by whether or not it has a more compressible phase in the core. In principle, cooling rates should depend on interior composition, but cooling calculations are beset by many uncertainties and competing assumptions about composition can yield similar cooling rates depending on other assumptions about superconductivity and the cooling processes. Moreover, for those stars for which a rate has been measured, not a single mass is known. It is unlikely that these measurements will yield conclusive evidence in the present state of uncertainty. Nevertheless, it may be possible to observe the phase transition in millisecond pulsars by the easiest of measurements---the sign of $\dot{\Omega}$. The sign should be negative corresponding to loss of angular momentum by radiation. However, as we will show, a phase transition, either first or second order, that occurs near or at the limiting mass star, can cause spin-up during a substantial era compared to the spin-down time of millisecond pulsars. We sketch the conventional evolutionary history of millisecond pulsars with the addition of the supposition that the critical density for quark deconfinement falls in the density range spanned by neutron stars. As already remarked, with this supposition, the star has a quark core from birth but its properties are so little effected that this fact cannot be discerned in members of the canonical pulsar population. However, by some mechanism, usually assumed to be accretion from a companion during the radio silent epoch following its $10^7$ year spin down as a canonical pulsar, the neutron star may be spun up. As it spins up, it becomes increasingly centrifugally deformed and its interior density falls. Consequently, the radius at which the critical phase transition density occurs moves toward the center of the star---quarks recombine to form hadrons. When accretion ceases, and if the neutron star has been spun up to a state in which the combination of reduced field strength and increased frequency turn the dipole radiation on again, the pulsar recommences spin-down as a visible millisecond pulsar. During spin-down the central density increases with decreasing centrifugal force. First at the center of the star, and then in an expanding region, the highly compressible quark matter will replace the less compressible nuclear matter. The quark core, weighed down by the overlaying layers of nuclear matter is compressed to high density, and the increased central concentration of mass acts on the overlaying nuclear matter, compressing it further. The resulting decrease in the moment of inertia causes the star to spin up to conserve angular momentum not carried off by radiation. The phenomenon is analogous to that of ``backbending'' predicted for rotating nuclei (Mottelson and Valatin 1960) and discovered in the 1970's (Johnson et.al. 1972, Stephens and Simon 1972) (see Fig.\ \ref{nucleus}) In nuclei, it was established that the change in phase is from a particle spin-aligned state at high nuclear angular momentum to a superfluid state at low angular momentum. The phenomenon is also analogous to an ice skater who commences a spin with arms outstretched. Initially spin decreases because of friction and air resistance, but a period of spin-up is achieved by pulling the arms in. Friction then reestablishes spin-down. In all three examples, spin up is a consequence of a decrease in moment of inertia beyond what would occur in response to decreasing angular velocity. | 98 | 3 | astro-ph9803070_arXiv.txt |
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9803 | astro-ph9803246_arXiv.txt | The observed Be and B relationships with metallicity clearly support the idea that both elements have a primary origin and that are produced by the same class of objects. Spallation by particles accelerated during gravitational supernova events (SNII, SNIb/c) seems to be a likely origin. We show, in the context of a model of chemical evolution, that it is possible to solve the Li, Be and B abundance puzzle with the yields recently proposed by Ramaty et al. (1997) provided that SNII are unable to significantly accelerate helium nuclei and that different mechanisms are allowed to act simultaneously. | The origin of the light elements is still an open question in Astrophysics. It is widely accepted that standard Big Bang nucleosynthesis can only produce $^7$Li. The primordial abundance produced of this isotope is one order of magnitude below Solar System values and roughly coincides with that observed by Spite \& Spite (1982) in the hot halo dwarfs ({\it the lithium plateau}). For this reason, it is generally accepted that such an abundance is representative of the primordial nucleosynthesis and that, since the LiBeB isotopes do not have a cosmological origin, they must be created by galactic activity. It is well known that the interaction of cosmic rays (CR) with the interstellar medium (ISM) can play an important role in the production of the LiBeB isotopes (Reeves, Fowler \& Hoyle 1970). In fact, the bulk of the $^6$Li, $^9$Be and $^{10}$B Solar System abundances can be naturally accounted with the standard model of galactic cosmic ray (GCR) nucleosynthesis. In this model, high energy protons and alpha particles collide with heavier nuclei (CNONe) present in the ISM to produce the light element isotopes (Meneguzzi, Audouze \& Reeves 1971). This model fails, however, to explain the present $^7$Li abundance and the Solar System $^7$Li/$^6$Li and $^{11}$B/$^{10}$B ratios. Furthermore, during the last decade, many observational studies have shown a linear relationship between Be and B abundances in metal-poor stars and the metallicity ([Fe/H]) (Rebolo et al. 1988; Gilmore et al. 1992; Boesgaard 1995; Molaro et al. 1997; Duncan et al. 1997). This {\it primary} behaviour cannot be easily explained by the standard GCR model: e.g Prantzos et al. (1993) used an ad-hoc hypothesis concerning the time evolution of the CR's escape-length, Abia et al. (1995) introduced artificial time dependences of the CR flux with the metallicity and the star formation rate and Casuso \& Beckman (1997) considered differential astration [see also Tayler (1995) and Yoshii, Kajino \& Ryan (1997)]. The reason for using these hypotheses is that in the standard GCR model the LiBeB production rates are proportional to the global metallicity of the ISM and to the CR flux. Since the latter is assumed to be proportional to the supernova rate it is, in consequence, also proportional to the production rate of metals in the galaxy. Such a dependence would predict a slope of about two in the B and Be relationships with [Fe/H] rather than a slope of one as the observations show. Furthermore, Duncan et al. (1997) and Garc\'\i a-L\'opez et al. (1998) have recently shown a nearly constant B/Be ratio of $\sim 20$ in dwarf stars for a wide metallicity range, although the uncertainties in this ratio are important\footnote{This B/Be value is obtained taking into account N-LTE effects in the derivation of B and Be abundances}. This value is still compatible with the idea of a spallative origin of Be and B (e.g Fields, Olive \& Schramm 1994). Since this ratio is similar to that observed in the Solar System (Anders \& Grevesse 1989) it cannot have experienced large variations during the galactic evolution, which strongly supports the idea of a similar origin for Be and B. A straightforward interpretation of these results is that the net production rate of Be and B does not depend on the metallic abundance in the ISM, i.e. the light-element production is not dominated by protons and alpha particles colliding with CNO nuclei but by these nuclei colliding with ambient protons and alpha-particles, probably in regions of massive star formation heavily enriched in these nuclei. The $\gamma$-ray observations from the Orion nebulae (Bloemen et al. 1994) provide additional support to this point of view since they are consistent with line emission from $^{12}$C$^*$ and $^{16}$O$^*$ produced by a large flux of low-energy ($<100$ MeV/nucleon) nuclei enriched in C and O. This might represent the first evidence of the existence of a considerable low-energy component in the spectrum of CRs (at least locally), as was suggested by different authors (Meneguzzi, Audouze \& Reeves 1971; Canal, Isern \& Sanahuja 1980). Since the ejecta of gravitational supernovae (type II, Ib/c) naturally match the above conditions, these objects have been proposed as preferential sites for LiBeB spallation production (Gilmore et al. 1992). The ejecta in these explosions are indeed heavily enriched in CO nuclei and, under some conditions (type Ib/c supernovae), the concentration of these nuclei might even exceed that of H and/or He. Because the yield of CO nuclei in supernovae is almost independent of the metallicity of the progenitor star, the LiBeB produced by spallation during supernova outcomes would have a primary character as is observed for Be and B. In fact, within the framework of a galatic evolutionary model, Vangioni-Flam et al. (1996) studied the production of Be and B, assuming a low-energy spectrum of the form $\rm{q(E)\sim E^{-n}}$, with $n=9$, and constant for $\rm{E\leq 30~MeV/n}$ in the CRs associated with Orion-like regions. They were able to explain the observed behaviour of Be and B vs. [Fe/H], and they obtained upper and lower limits for the contribution of this mechanism to the galactic evolution of Be and B abundances. Their results show that this mechanism might contribute up to $70\%$ of the observed Be and B abundances and that standard GCR nucleosynthesis is not the main source of $^9$Be and B in the galaxy. Recently, Ramaty et al. (1997; hereafter RKLR) studied the influence of different spectra and chemical compositions in the CRs produced by supernova on the LiBeB production. They showed from energetic arguments that due to the amount of Be necessary to account for its linear behaviour with [Fe/H], the CNO-rich, He-poor and H-poor CR source compositions are favoured. The observed Be/Fe ratio requires the investment of about $3\times 10^{49}$ to $2\times 10^{50}$ erg per gravitational supernova in these metallic CR, depending on whether or not H and He are accelerated with metals. Similar arguments led them to conclude that these CRs should have a hard-energy spectrum extending up to at least 50 MeV/n. From the constancy of the observed Be/Fe ratio and metallicity, they also derived the necessary Be yield per supernova that is needed. Assuming a $^{56}$Fe yield per SNII (Woosley \& Weaver 1995) of $\sim 0.11$ M$_\odot$, they obtained a Be yield of $2.8\times 10^{-8}$ M$_\odot$ (within a factor of two of uncertainty from the observed Be/Fe ratio), almost irrespective of the progenitor star metallicity. In this paper we have assumed that the Be yield from RKLR is representative of the Be produced per gravitational supernova. Since this yield automatically sets the corresponding Li and B yields for a given CR spectrum and chemical composition, we use this to study the impact of such objects on the galactic evolution of the light element abundances. We discuss the results in the framework of different scenarios for the progenitors of type II and Ib/c supernovae and possible mechanisms for the CR acceleration in such objects. | 1) Type Ib/c supernovae seem to be key objects for the production of Be and B by spallation. However, due to their low rates in the galaxy, an additional contribution to the Be and B production by spallation due to type II supernovae is needed. Using a reasonable shock spectrum for the CR source, it is possible to explain the observed linear relation of Be and B abundances with metallicity although a neutrino contribution to B is necessary to account for the observed B/Be$\sim 20$. In this situation, the standard GCR model would play a minor role in the early evolution of the Be and B abundances. Additional sources of $^7$Li are still necessary to account for the evolution of its abundance. Objects with longer lifetime than supernova like AGB stars (and/or novae) seem to be necessary. A multi-source nature of this element is thus obvious. 2) This scenario has a serious problem with the acceleration of CO-rich and HHe-poor matter in SNII ejecta to sufficient energies for spallation reactions. It is not clear how a huge Li production (due to $\alpha+\alpha$ reactions) in SNII relative to that of Be and B can be avoided if the acceletared particles are CO-poor. If that is the case a production ratio Li/Be$\geq 100$ is predicted for a wide variety of CR spectra. This high Li production would be incompatible with the lithium plateau observed at low metallicity as shown in Figure 5 (short dash- dotted line). 3) The CR source spectrum shape associated with supernova explosions is another crucial point. It must extend to high kinetic energies (E$\geq 50$ MeV/n) in order for this scenario to be energetically plausible and to avoid the overproduction of Li with respect to Be and B. Studies of LiBeB production using other CR spectra, not strict power laws, deserve special attention. 4) The existence of an additional source of $^{11}$B seems obligatory in order to explain the Solar System's $^{11}$B/$^{10}$B$=4.05$ value. Production of $^{11}$B by neutrino spallation seems to be the best candidate. In this case, a $^{11}$B/$^{10}$B ratio higher than 4 at low metallicities is predicted, which might be tested by using very high resolution spectroscopy on very large telescopes (Rebull et al. 1997). This work has been partially supported by CICYT grants PB96-1428 and ESP95-00091 and by a CIRIT grant. | 98 | 3 | astro-ph9803246_arXiv.txt |
9803 | astro-ph9803160_arXiv.txt | s{We present a method of subtracting the foreground contamination for the measurement of CMB polarization. We calculate the resultant errors on CMB polarization and temperature-polarization cross correlation power spectra for the high frequency instrument (HFI) aboard Planck Surveyor, and estimate the corresponding errors on cosmological parameters} | The upcoming satellite CMB experiment Planck surveyor offers an unprecedented opportunity to measure CMB polarization. A major hurdle in extracting the primary CMB signal from data, apart from noise, is galactic and extragalactic foregrounds. However, as the foregrounds differ from CMB in both frequency dependence and spatial distribution, one can hope to reduce their level in a multi-frequency CMB experiment. A multi-frequency Wiener filtering method to implement this scheme was developed \cite{bouchet,teg} and applied to cleaning the simulated CMB temperature map by Bouchet {\em et al.} \cite{bouchet}. They showed that the residual contamination after cleaning the map is much smaller than the CMB primary signal, and therefore the foregrounds may not be a major obstacle in the extraction of CMB temperature angular power spectrum. However, as the CMB polarization signal is expected to be one to two orders of magnitudes below the temperature signal, it is likely to be comparable to both the experimental noise and the level of foregrounds. Sethi {\em et al.} \cite{set} modelled and estimated the level of dust polarized emission at high galactic latitudes. Dust polarization is expected to be the dominant contaminant for measuring CMB polarization using Planck HFI. In this paper, we extend the multi-frequency Wiener filtering method to include the polarization and temperature-polarization cross-correlation. The aim of this exercise is to quantify errors in estimating various power spectra and consequently the errors in cosmological parameters. Our results are relevant for Planck HFI. | 98 | 3 | astro-ph9803160_arXiv.txt |
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9803 | astro-ph9803026_arXiv.txt | When Baade \& Zwicky proposed in 1934 that collapsed stars composed of neutrons could be formed in supernova explosions, they had little notion of how such creatures would manifest themselves observationally. The surprise discovery of pulsars, and in particular, of the Crab and Vela pulsars in their respective supernova remnants, heralded the first visual image of the isolated neutron star: that of a compact, highly magnetized (surface field $\sim 10^{12}$~G) star, spinning down slowly due to magnetic braking, emitting a collimated beacon, and exciting its surroundings via the injection of ultra-relativistic particles. However, even the simplest follow-up questions, such as whether all neutron stars form in supernovae, what fraction of supernovae produce neutron stars, and in particular, whether all young pulsars are born with properties like those of the Crab and Vela pulsars remain to this day naggingly unanswered. Here I summarize observations of neutron stars plausibly associated with supernova remnants. The last similar review was published by Helfand \& Becker (1984). Their interesting synthesis of the observational data is beyond the scope of this paper. More recent reviews including only radio pulsar/supernova remnant associations can be found elsewhere (Kaspi 1996; Kaspi 1998). | The most striking conclusion following examination of the associations between neutron stars and supernova remnants is that the paradigm that every young neutron star is like the Crab pulsar requires reconsideration. Neutron stars, it seems, come in many flavors, perhaps some yet to be discovered. The continued analysis of valuable archival high-energy data, the upcoming launch of the {\it Advanced X-ray Astrophysics Facility}, as well as the ongoing Parkes multi-beam and Nan\c{c}ay Galactic plane survey should ensure significant progress, and even perhaps more surprises in this field in the near future. | 98 | 3 | astro-ph9803026_arXiv.txt |
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9803 | astro-ph9803210_arXiv.txt | We try to explain the unusually high far-infrared emission seen by IRAS in the double-lobed radio-loud quasars 3C\,47, 3C\,207 and 3C\,334. High resolution cm--mm observations were carried out to determine their radio core spectra, which are subsequently extrapolated to the far-infrared in order to determine the strength of the synchrotron far-infrared emission. The extrapolated flux densities being considerably lower than the observed values, a significant nonthermal far-infrared component is unlikely in the case of 3C\,47 and 3C\,334. However, this component could be responsible for the far-infrared brightness of 3C\,207. Our analysis demonstrates that nonthermal emission cannot readily account for the difference between quasars and radio galaxies in the amount of their far-infrared luminosity. On the other hand, a significant role for this mechanism is likely; full sampling of the mm-submm spectral energy distributions is needed to address the issue quantitatively. | \label{intro} Thermal emission from cool circumnuclear dust is widely accepted as an important mechanism for producing far-infrared radiation in radio-quiet and radio-loud AGN (Sanders et al. 1986, Chini et al. 1989, Antonucci et al. 1990). This emission -- provided isotropic -- gives the opportunity to test the orientation unification model of radio-loud quasars (QSRs) and powerful radio galaxies (RGs) (Barthel 1989, Urry \& Padovani 1995), since matched samples should yield similar far-infrared output for both classes. However, from IRAS observations it appears that QSRs are stronger 60$\mu$m emitters than Fanaroff \& Riley class~II RGs (Heckman et al 1992, 1994, Hes, Barthel \& Hoekstra 1995). This would imply that the simple unification model does not hold or that the assumption of isotropic 60$\mu$m radiation is incorrect. In support of the latter, Pier \& Krolik (1992) drew attention to moderate anisotropy effects in the thermal far-infrared radiation from optically thick tori surrounding AGN. Furthermore, Hoekstra, Barthel \& Hes (1997) demonstrated that beamed nonthermal far-infrared radiation is not insignificant in radio sources having prominent nuclei, and that the 60$\mu$m differences might be attributed to a stronger, beamed, nonthermal component in the QSR class. Within the unified model framework, QSRs would -- at decreasing angle to the line of sight -- naturally lead into the radio-loud blazar class, objects in which the dominance of the beamed component is without doubt. Indeed, both blazars and core-dominated QSRs display a smooth, single component spectral energy distribution, extending from radio to X-rays, indicating that all continuum radiation is of synchrotron origin (Landau et al. 1986). This causes blazars and core-dominated QSRs to be the most luminous IRAS AGN (Impey \& Neugebauer 1988). IRAS detected three $z\sim0.5$, lobe-dominated 3CR QSRs at 60$\mu$m, which corresponds to $\approx 40\mu$m emitted wavelength. Comparable radio galaxies were not detected -- contrary to the expectation within the simple unification model. The QSRs, 3C\,47, 3C\,207, and 3C\,334, have relatively bright radio cores and large double-lobed structures. Two of these, 3C\,47 and 3C\,334, had been observed to display superluminal motion -- a clear sign of beamed emission (e.g., Zensus \& Pearson 1987) -- thus raising the suspicion that their far-infrared brightness could be (partly) due to a beamed non-thermal component. The goal of the present research is to investigate to what extent (beamed) nonthermal radiation can account for the infrared emission in lobe-dominated QSRs. We have determined the cm--mm core spectra of 3C\,47, 3C\,207, and 3C\,334 in order to assess their nonthermal 60$\mu$m radiation by means of extrapolation. Given the possibility of flaring submm components (observed in blazars -- e.g., Brown et al. 1989) this research is a first attempt to isolate nonthermal and thermal FIR radiation. Observations over a wider frequency range are forthcoming. | \label{conclusions} We cannot readily explain the extraordinary high FIR flux density observed in the 3C\,47, 3C\,207, and 3C\,334. While relativistic beaming of a single nonthermal component cannot account for the total FIR radiation, the radiation of relativistically beamed multiple core components is likely to contribute significantly in the far-infrared. Thermal mechanisms could play a role, in combination with optical thickness and aspect effects, but we need more data in the submm and infrared to address these in detail. We are currently engaged in measurements of matched pairs of radio galaxies and quasars, comparing their thermal FIR output by combining cm, mm, submm and FIR data from VLA, JCMT, and ISO observations. In the meantime, we have little doubt that the importance of nonthermal FIR and submm radiation in double-lobed radio sources has been underestimated. \bigskip \noindent {\bf Acknowledgments} \noindent We acknowledge initial involvement in this work of R.~Hes and H.~Hoekstra, OVRO support by N.~Scoville and A.~Sargent, and the use of archival VLA data from R.~Ekers, R.~Perley, and J.~Wardle. The NRAO VLA is operated by Associated Universities, Inc. under contract with the National Science Foundation. OVRO research is supported in part by NSF Grant AST~93-14079. | 98 | 3 | astro-ph9803210_arXiv.txt |
9803 | astro-ph9803338_arXiv.txt | Preliminary results of a search for distant clusters of galaxies using the recently released I-band data obtained by the ESO Imaging Survey are presented. In this first installment of the survey, data covering about 3 square degrees in I-band are being used. The matched filter algorithm is applied to two sets of frames that cover the whole patch contiguously and these independent realizations are used to assess the performance of the algorithm and to establish, from the data itself, a robust detection threshold. A preliminary catalog of distant clusters is presented, containing 39 cluster candidates with estimated redshifts $0.3 \leq z \leq 1.3$ over an area of 2.5 square degrees. | \label{sec:intro} One of the primary goals for undertaking the ESO Imaging Survey (EIS; Renzini \& da Costa 1997) has been the preparation of a sample of optically-selected clusters of galaxies over an extended redshift baseline for follow-up observations with the VLT. High-redshift clusters are, of course, a primary target for 8-m class telescopes. A large and well defined sample of clusters can be used for many different studies, ranging from the evolution of the galaxy population, to the search for arcs and lensed high redshift galaxies, to the evolution of the abundance of galaxy clusters, a powerful discriminant of cosmological models. In addition, individual clusters may be used for weak lensing studies and as natural candidates for follow-up observations at X-ray and mm wavelengths, which would provide complementary information about the mass of the systems. For some of these applications it suffices to find a large number of clusters, while for others it is vital to obtain a full understanding of the selection effects, to generate suitable statistical samples. Until recently, only a handful of clusters were known at redshifts $z \gsim 0.5$; visual searches for high redshift clusters were conducted by Gunn \etal (1986) and Couch \etal (1991), but their samples are severely incomplete beyond $z \sim 0.5$; at higher redshifts targeted observations in fields containing known radio-galaxies and QSOs have produced a handful of cluster identifications (\eg Dickinson 1995; Francis \etal 1996; Pascarelle \etal 1996; Deltorn \etal 1997). The first objective search for distant clusters was conducted by Postman \etal (1996; hereafter P96) using the 4-Shooter camera at the 5-m telescope of the Palomar Observatory. In their survey 10 out of the 79 cluster candidates have estimated redshift $\gsim 0.7$. Further evidence for the existence of clusters at high redshift has been obtained from X-ray (\eg Gioia \& Luppino 1994; Henry \etal 1997; Rosati \etal 1998), optical (\eg Connolly \etal 1996; Zaritsky et al. 1997) and infrared (Stanford \etal 1997) searches. However, the existing samples are small, and their selection effects largely unknown. Recently, observations of the first patch of the EIS, covering about 3 square degrees, have been completed, and the data made available to the community (Nonino \etal 1998; hereafter Paper I). Although the data are still in a preliminary form, much can already be learned regarding the characteristics of the sample of candidate clusters that can be detected using the EIS data. In this paper preliminary catalogs of objects detected on single 150 sec. I band frames are used (see Sects. \ref{sec:obs_and_data} and \ref{sec:gal_cat}) mainly to assess the capability of the EIS to detect clusters of galaxies at $z \gsim 0.5$. A discussion of a full cluster sample based on the galaxy catalog extracted from the coadded EIS images, is postponed to a future paper (Scodeggio \etal 1998). The reason for using here the single-frame catalogs is that they provide two independent datasets for the same area of the sky. The comparison between the cluster detections obtained using the two catalogs separately, can be used to quantify the reliability of the cluster detection procedure. When the handling of catalogs extracted from the coadded images is fully implemented in the EIS data reduction pipeline, the cluster search will be carried out using those catalogs, instead of the single-frame ones, to benefit from the deeper limiting magnitude of the coadded images. In the meanwhile a better quantification of the detection limits for distant clusters of galaxies within the EIS data could be obtained by comparing the results presented here with those obtained using independent cluster search methods. In Sects. ~\ref{sec:obs_and_data} and ~\ref{sec:gal_cat} the observations, data reduction and the object catalogs, that are used for the cluster search, are briefly discussed. The cluster finding procedure, based on the matched-filter algorithm proposed by P96, is described in Sect. \ref{sec:cluster_finding}. In Sect. ~\ref{sec:results} the preliminary cluster catalog is presented, and the properties of the detected candidates are discussed. In Sect. ~\ref{sec:future} conclusions of this work are summarized, and its possible extensions to the search for clusters using the coadded EIS images discussed. | 98 | 3 | astro-ph9803338_arXiv.txt |
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9803 | astro-ph9803048_arXiv.txt | Preliminary trigonometric parallaxes and BVI photometry are presented for two dwarf carbon stars, LP765$-$18 (= LHS1075) and LP328$-$57 (= CLS96). The data are combined with the literature values for a third dwarf carbon star, G77$-$61 (= LHS1555). All three stars have very similar luminosities (9.6$\,<\,$M$_{\rm V}<\,$10.0) and very similar broadband colors across the entire visual-to-near IR (BVIJHK) wavelength range. Their visual (BVI) colors differ from all known red dwarfs, subdwarfs, and white dwarfs. In the M$_{\rm V}$ versus V$-$I color--magnitude diagram they are approximately 2 magnitudes subluminous compared with normal disk dwarfs with solar-like metallicities, occupying a region also populated by O-rich subdwarfs with $-1.5<$[m/H]$<-1.0$. The kinematics indicate that they are members of the Galactic spheroid population. The subluminosity of all three stars is due to an as-yet-unknown combination of (undoubtedly low) metallicity, possibly enhanced helium abundance, and unusual line-blanketing in the bandpasses considered. The properties of the stars are compared with models for the production of dwarf carbon stars. | The discovery of the first dwarf carbon star (\cite{dah77}) sparked considerable research on the evolution of stars to and through the dwarf carbon star (dC) stage. It is now accepted that dC stars are formed through the accretion of carbon-rich material from a close companion star that was evolving on the asymptotic giant branch at the time of mass transfer. They are now recognized as members of a larger family of stars that have undergone mass-transfer binary evolution, a family that includes the halo CH giant stars, the disk Ba giants, the Ba dwarf or CH subgiant stars, and the extrinsic S stars (\cite{gre97}). De Kool \& Green (1995) have modeled dC star formation by following the evolution of a large variety of binary systems constructed from the observed distributions of component masses, orbital separations, and metallicities for unevolved binaries. Although these simulations involve a large number of poorly constrained parameters, the results indicate that the formation of a dC star is strongly favored by low initial metallicity, and that virtually no dCs are produced in systems with metallicities above half solar. Furthermore, the mass distribution is found to peak strongly in the 0.4 to 0.9 solar mass range. Despite the studies to date, our understanding of dC stars remains highly speculative, based as much on plausibility arguments as on objective facts. Distances of these stars are among the most essential missing data, and must be known in order to determine each star's luminosity, space velocity, mass, and evolutionary status. Presently, only one dC star, G77$-$61, has a reliable trigonometric distance determination and it indicates a high space velocity and -- by implication -- low metallicity. This implied low metallicity is consistent with the detailed atmospheric analysis of G77$-$61 by Gass et al. (1988) who derived the extremely low metallicity of [Fe/H]=$-5.6$, a value which certainly implies very early epoch halo formation. A larger sample of dC stars with properties well constrained by fundamental observational data is clearly needed to understand these stars as a class. Unfortunately, no known dC stars are bright enough to have had their parallaxes measured with high accuracy by the Hipparcos satellite. In this paper we report preliminary trigonometric parallaxes for two additional dC stars and discuss some of their properties. | Results interpreted from the parallax data are shown in Table 2. In calculating the formal uncertainties in M$_{\rm V}$ and M$_{\rm K}$ we have adopted the parallax uncertainties given in Table 1, along with $\pm\,$0.02 mag for the uncertainties in the V magnitudes and $\pm\,$0.03 mag for the uncertainties in the K magnitudes. The tabulated radial velocities (V$_{\rm r}$) are from Bothun et al. (1991) for LP765$-$18 and LP328$-$57 and from Dearborn et al. (1986) for G77$-$61. The space velocity components (U,V,W) in Table 1 have been corrected for a solar motion of (10, 15, 7) km s$^{-1}$ to the local standard of rest; U is in the direction of the galactic center. Inspection of Tables 1 and 2 reveals that these three dC stars are remarkably similar in terms of (1) overall kinematic properties, (2) infrared (JHK) colors, (3) optical (BVI) colors, and (4) luminosities (M$_{\rm V}$ and M$_{\rm K}$). Regarding the kinematics specifically, the large overall space velocities and the large negative V galactic velocity components are indicative of membership in the Galactic spheroid population. Regarding the colors, Green et al. (1992) have previously discussed the location of dC stars in the J$-$H versus H$-$K diagram and their apparent separation from the giant and subgiant carbon stars, presumably due (in part) to the higher surface gravities for the dC stars. In Figure 1 we show the location of these dC stars in the B$-$V versus V$-$I diagram. Because all three were originally identified as stars with high proper motion, we include for comparison a selection of other field stars commonly identified in proper motion surveys: later-type degenerates (open circles), disk dwarfs (filled circles), and metal-poor subdwarfs (crosses). The dC stars occupy a unique region in this diagram. The lone exception to a clear-cut separation for the dC stars is LP701$-$29, the only known late-type degenerate with strong CaI absorption which blankets the B bandpass (\cite{dah78}). Figure 1 suggests that even broadband photometric surveys of faint, high proper motion stars might succeed in isolating additional dC candidates. Their location to the left of the most extreme metal-poor field subdwarfs known at present suggests metallicities below [m/H]$\sim-$2. However, quantitative conclusions about metallicity are not possible because the field subdwarfs have O-rich rather than C-rich atmospheres; with the exception of the analysis of G77$-$61 by Gass et al. (1988), no complete studies based on model atmospheres have been presented for metal-poor, C-rich stars with dwarf-like gravities. Figure 2 shows the location of these dC stars in the M$_{\rm V}$ versus V$-$I color magnitude diagram. The solid line represents the observed mean disk main sequence as defined by USNO parallax stars (cf. \cite{mon92}). The dashed lines show the metal-poor main sequences modeled by Baraffe et al. (1997) for metallicities (scaled from solar) of [m/H]= $-1.0, -1.5$ and $-2.0.$ These authors demonstrated that their models successfully reproduce the main sequence of several globular clusters over this range of metallicities. However, these models are for O-rich atmospheres -- not the C/O$>$1 compositions which are appropriate for the dC stars. Furthermore, the position of dC stars in color-magnitude diagrams cannot necessarily be interpreted simply in terms of metallicity because of the possibility of enhanced helium abundance in their atmospheres, produced by the mass transfer event(s) that increased their carbon abundances. The complicated atmospheric situation and the difficulty in uniquely determining both log$\,g$ and the He abundance has been discussed by Gass (1988) and Gass et al. (1988). Therefore, until more detailed spectroscopic investigations are undertaken to derive the helium abundance independently, we are left with the implication from the kinematic information that low metallicity and nearly normal helium abundances are most likely. Figure 3 shows the location of the dC stars in the schematic M$_{\rm K}$ versus I$-$K diagram. Here the solid line represents the ``young disk'' stars defined by Leggett (1992) while the dashed lines are again the metal-poor main sequences models from Baraffe et al. (1997). In their analysis of G77$-$61, Gass et al. (1988) pointed out that the J and K fluxes should not depend strongly on the carbon abundance whereas the fluxes in the B,V,I and H bandpasses will depend strongly on composition. The three dC stars are also significantly subluminous with respect to the disk main sequence in this diagram. However, it is also clear that going to a redder color index still does not allow a quantitative interpretation in terms of metallicity because the location of G77$-$61 above the [m/H]=$-1.0$ curve is clearly at odds with the value of [Fe/H]=$-5.6$ derived by Gass et al. (1988). Figure 4 shows the location of G77$-$61 in the M$_{\rm V}$ versus T$_{\rm eff}$ plane using the value of T$_{\rm eff}\,=\,$4200$^{\circ}$K derived by Gass et al. (1988). Once again the dashed lines are from Baraffe et al. (1997). Here we find at least qualitative agreement in that this extremely metal-poor object lies below the [m/H]=$-2.0$ curve. This emphasizes the need for higher resolution spectrophotometric observations and C-rich model atmosphere analyses for LP765$-$18 and LP328$-$57 in order to establish (at least) T$_{\rm eff}$ values for them. The models of Baraffe et al. (1997) indicate that the derived masses are primarily sensitive to the absolute magnitudes (or luminosities), and only weakly dependent on the colors. Acknowledging the dangers of interpreting the dC stars with O-rich models, the masses so estimated are 0.39 M\solar\ for LP765$-$18, 0.37 M\solar\ for LP328$-$57, and 0.30 M\solar\ for G77$-$61. Models for dC star formation predict that a range of masses from 0.2 to 1.0~M{\solar} should exist (\cite{dek95}). However, the frequency distribution of these masses depends critically on the adopted Initial Mass Ratio Distribution (IMRD) for the unevolved binary which is quite uncertain. For a flat IMRD (i.e., dN$\,\propto\,$dq, where q is the mass ratio), the distribution for spheroid dC stars peaks rather sharply around 0.7 M\solar\ at a space density of dN/dlog$\,$M = 4$\,{\rm x}\,10^{-7}$ pc$^{-3}$, falling to roughly dN/dlog$\,$M = 2$\,{\rm x}\,10^{-7}$ pc$^{-3}$ for stars in the mass range inferred above. On the other hand, for uncorrelated component masses, the models predict a broader peak spanning 0.2 to 0.8 M\solar\ with a space density of dN/dlog$\,$M = 4$\,{\rm x}\,10^{-7}$ pc$^{-3}$ for spheroid stars. As noted by de Kool \& Green (1995), the observed absolute magnitudes seem to support an IMRD with uncorrelated component masses. However, the apparent similarity of these three dC stars is undoubtedly influenced to some extent by selection effects. Hotter, more massive dC stars are predicted by the models to be common, but hotter stars will have weaker carbon bands and may have spectra and broad-band colors that are not as obviously different from stars with oxygen-rich atmospheres as are these three dC stars. Another factor that makes the model predictions uncertain is the amount of dilution of the carbon-rich material being transferred from the asymptotic-giant-branch star onto its main-sequence companion: a more massive main sequence star (0.5--1.0~M{\solar}, for example), has a less massive convective envelope than the stars studied here (reducing the dilution), but the carbon-rich material may be mixed into the radiative zone as well (\cite{pro89}), thus increasing the dilution and making hotter dC stars more difficult to identify. Furthermore, these three dC stars have been chosen for parallax observations based partly on their unusually-high proper motions. This kinematic selection partially accounts for their all being members of the Galactic spheroid. Disk dC stars may be common --- the ratio of disk/spheroid dC stars is a quantity that will help constrain the models of dC-star formation. Parallax data for a sample of dC stars with smaller proper motions could help to address this issue. Several such dC stars are probably close enough to get significant parallax measurements -- however, none are presently being observed in the Naval Observatory parallax program pending time becoming available in the program at their respective right ascensions. We conclude that the similar properties of these three dC stars and their membership in the Galactic spheroid may result in part from the way these stars were selected. Accurate distances for a larger sample of dC stars are needed, and identification of dC stars without kinematic bias is essential in order to fully understand how and where dC stars are produced. | 98 | 3 | astro-ph9803048_arXiv.txt |
9803 | astro-ph9803081_arXiv.txt | Oscillations observed in the light curve of Nova V1974 Cygni 1992 since summer 1994 have been interpreted as permanent superhumps. From simple calculations based on the Tidal-Disk Instability model of Osaki, and assuming that the accretion disc is the dominant optical source in the binary system, we predict that the nova will evolve to become an SU UMa system as its brightness declines from its present luminosity by another 2-3 magnitudes. Linear extrapolation of its current rate of fading (in magnitude units) puts the time of this phase transition within the next 2-4 years. Alternatively, the brightness decline will stop before the nova reaches that level, and the system will continue to show permanent superhumps in its light curve. It will then be similar to two other old novae, V603 Aql and CP Pup, that still display the permanent superhumps phenomenon 79 and 55 years, respectively, after their eruptions. We suggest that non-magnetic novae with short orbital periods could be progenitors of permanent superhump systems. | \subsection{The (permanent) superhump phenomenon} Regular superhumps are quasi-periodic oscillations that appear in the light curve (LC) of the SU UMa subclass of dwarf novae, superimposed on the superoutburst LC of these systems. The superhump period is a few percent longer than the orbital period of the system in which it is observed (laDous 1993). Permanent superhumps (PSH), on the other hand, appear permanently or most of the time in systems where they prevail, and do not demand a superoutburst for their emergence. Superhumps are detected in the LC of cataclysmic variables (CVs) with short orbital periods ($P_{orbital}$=17-200 min. - Ritter \& Kolb 1998). Stolz \& Schoembs (1984) found a linear relation between the relative excess of the superhump period over the orbital one, and the orbital period. The Tidal-Disc Instability model (for a review see Osaki 1996) is the commonly accepted explanation of the phenomenon. The superhump periodicity is the beat of the orbital period of the binary system with the period of the precession of the accretion disc around the white dwarf (WD) in the binary co-rotating frame. The difference between regular superhumps and PSH is understood as a result of the much larger mass transfer rate in the system showing PSH than in SU UMa systems. In fact, Osaki's schematic diagram (Fig. 1) states that the values of only two parameters, the orbital period and the mass transfer rate, determine the basic differences among the four major classes of CVs, namely, U Gem, SU UMa, PSH and Nova-like variables. \begin{figure} \centerline{\epsfxsize=3.0in\epsfbox{fig1nneps.eps}} \caption{A schematic theoretical diagram based on Osaki (1996), showing the location on the (Orbital Period, $\dot{M}$) plane of four major groups of CVs: UG=U Geminorum, SU UMa, PSH=permanent superhumps and NL=nova like variables. The two dashed vertical lines represent the two ends of the well known period gap in the period distribution of CVs. Systems on the right hand side, the UG and NL groups, have accretion discs that are tidally stable. Systems on the left, PSH and SU UMa, are tidally unstable. Systems above the dotted tilted line, PSH and NL, have thermally stable discs. Systems below that line, UG and SU UMa, have thermally unstable discs. The upper dot represents the present location of V1974 Cyg in this plane. The arrow indicates the expected decrease in the $\dot{M}$ value in the future. The lower point is the location of value of $\dot{M}$, corresponding to the pre-nova magnitude according to Skiff (1997). If the fading of the nova, in magnitude units, continues linearly, the arrow will intersect the tilted line around the year 2000.} \end{figure} \subsection{Superhumps in Classical Nova Systems} Photometric observations of V1974 Cyg (N. Cyg 1992) revealed two periodic oscillations in the LC of the star (Retter, Ofek \& Leibowitz 1995). While there is an overall agreement that the shorter is the orbital period of the binary system, two interpretations for the nature of the second period, which is about 5\% larger than the orbital one, have been suggested. Semeniuk et al. (1995) and Olech et al. (1996), argued that it is the spin period of a rotating WD. Retter, Leibowitz \& Ofek (1997) and Skillman et al. (1997) proposed, on the other hand, that the longer period is that of PSH oscillations. We think that the increase of this period during 1995 and 1996, as well as other features of the optical LC of the nova discussed in Retter et al. (1997), argue against the magnetic interpretation for the longer period, and for the PSH one, which we adopt in this letter. Two other old novae, V603 Aql 1918 (Patterson \& Richman 1991; Patterson et al. 1993; Thomas 1993 and Patterson et al. 1997) and CP Pup 1942 (White \& Honeycutt 1992; White, Honeycutt \& Horne 1993 and Thomas 1993), also show PSH in their LCs. While the presence of PSH in the LC of V603 Aql is rather well established, there is still some controversy in the case of CP Pup. White et al. (1993) and Balman, Orio \& Ogelman (1995) proposed a WD spin interpretation for the second periodicity in the LC of this nova, which was initially thought to be 11\% longer than the orbital period. However, an extensive photometric study by Thomas (1993) revealed that the period excess is only about 2\%, and that the two periods obey the well-known relation of Stolz \& Schoembs (1984) between orbital and superhumps periods in SU UMa systems. In addition, the spin periods in all but one intermediate polars are shorter than their orbital periods (Patterson 1994). Even in the one exceptional case (RX J19402-1025), the period excess is very minute - $(P_{spin} - P_{orbital})/ P_{orbital} \approx$ 0.3\% (Patterson et al. 1995), much smaller than in a typical superhump system. We believe that the observed photometric features of CP Pup favour the superhump interpretation for this system and we adopt it in this work. There are a few more reports on short periods in other novae. RW UMi 1956, with a possible orbital period of $\sim$2~hr (Szkody et al. 1989), is a natural permanent superhump candidate, but intensive continuous photometry of the nova during 25 nights in 1995 and 1997, carried out at the Wise Obs., suggests that the variation is irregular. Two periods of $\sim$3.3~hr, which differ from each other by less than 2\%, were found in V1500 Cyg 1975. These periods, however, don't obey the Stolz \& Schoembs relation. The variation in the polarization, found in this system, is also a very strong evidence for the magnetic nature of the WD. (Stockman, Schmidt \& Lamb 1988). There are also some evidence that the 85-m period found in GQ Mus 1983 is caused by the rotation of the magnetic pole of the primary (Diaz \& Steiner 1994), thus cannot be identified with a superhump variation. | \subsection{Consistency} The value of the parameters of the three nova systems that we used in the previous section were derived by various researchers, in general independently of the superhumps phenomenon. Our calculations show that according to Osaki's theory, the very presence of superhumps in the LCs of these systems is indeed expected from these values. Further support for the validity of the calculations presented in this work, particularly for the bolometric correction that we used in Section 2.1 (equation 5), comes from the agreement of the mass transfer rate that we obtain for V603 Aql with estimates by other authors using other methods. We derive in Section 2.3 for this parameter the value $8\pm4 \times 10^{17}$ gr/sec. It is in good agreement with $4.8 \times 10^{17}$ gr/sec (Krautter et al. 1981), $7.6 \times 10^{17}$ gr/sec (Wade 1988), and $2.6 \times 10^{17}$ gr/sec (Duerbeck 1992). \subsection{The principal light source} Our calculations in Section 2 were made with the assumption that the visual light of all three novae emanates from the accretion discs in these systems. This assumption is supported by the optical spectrum of the three stars. Shai Kaspi kindly took for us two spectrograms of V1974 Cyg with the FOSC camera at the Wise Obs., one in 1995 July and one in 1996 August. The two spectra are dominated by nebular emission lines and by a continuum that seems to be free of stellar absorption features. Similarly, spectra of V603 Aql (Williams 1983) and of CP Pup (O'Donoghue et al. 1989, White et al. 1993), which were taken 69 years after outburst in the V603 Aql case, and 43, 45 and 46 years after outburst in the CP Pup case, do not show any obvious stellar absorption features. Thus, any contribution of the secondary to the light of these systems is indeed negligible. Leibowitz (1993) has drawn attention to the appearance of a kink in the visual LC of many classical novae a few tens of days after maximum light. He suggested that the abrupt change in the slope of the decaying LC signifies a transition of the main light source in the system from the envelope of the contracting nova to the accretion disc in the system. In the two old novae, V603 Aql and CP Pup, the photometry capable of detecting superhumps was performed many years after the outburst, long after the appearance of the kink in their LCs. In V1974 Cyg, the superhumps were also detected only after the kink in the LC of this nova, as shown in Fig.~2. This figure is a comprehensive visual light curve of the nova, from outburst in 1992 February to 1997 May. The data were taken from various amateur groups. The arrow in the figure indicates the time when superhumps were first observed. Since superhumps are a pure disc phenomenon and consistently with Leibowitz's hypothesis we may conclude that presently in V1974 Cyg, the accretion disc is indeed the major source of the optical continuum. \subsection{The future of the permanent superhumps in the two old novae} We showed in Section 2 that in the three novae discussed in this letter, the present visual magnitude is brighter than the critical value for transition from the PSH to the SU UMa state. The future of the superhumps in their LCs is correlated with the future run of their visual LCs. The question of what is the long term behaviour of the LC of classical novae, many years after outburst, is a wide open one. It is difficult to answer it observationally because of the scarcity of these objects on one hand, and the early age, of less than a century of modern observations in novae, on the other. This time duration is probably still shorter than the characteristic time of the returning of classical novae to their real quiescence state. In two pioneering works, Vogt (1990) and Duerbeck (1992) attempted to systematically investigate the long term photometric behaviour of classical novae. Duerbeck's sample of 21 old novae includes also the two old novae discussed in this letter. There are large systematic difficulties and uncertainties in the determination of the decline rate of old novae, as underlined by Duerbeck himself. He, nevertheless, suggested for V603 Aql an average decline rate of 10.7 mmag/year, and for CP Pup 36.9 mmag/year (Duerbeck 1992 Table 1). Taking these numbers at face value and using our results of Section 2.3, we obtain that V603 Aql will cross the critical line in Fig. 1 not before some 75 years in the future from today. CP Pup will undergo this phase transition not before some 55 years from today. In view of the large uncertainties in the value of the critical visual magnitude of these two novae, as well as in their average decline rate, these time intervals are rough lower limits at best. \subsection{The future of the permanent superhumps in Nova Cygni 1992} V1974 Cyg is still in a phase of relatively steep decline from its recent outburst. In its visual LC shown in Fig. 2, three sections of nearly linear decline, at three different slopes, are clearly discernible. The decline rate in the last section, lasting now for more than two years, is about 0.7 mag/year. If the nova keeps this decline rate for another few years, it will reach the critical visual magnitude, indicated by the dot in the figure, sometime around the year 2000. The vertical error bars around the dot denote the formal uncertainty in the critical visual magnitude value (see Section 2.2). The horizontal bars represent the uncertainty in the time of crossing the line of phase transition, due to the uncertainty in the slope of the present average linear decline. \begin{figure} \centerline{\epsfxsize=3.0in\epsfbox{fig2neps.eps}} \caption{Five years light curve of Nova~Cyg~92. The data were taken from visual estimates of three groups of amateur astronomers - AFOEV (main source), VSNET and AAVSO. The solid line is a linear extrapolation based on the mean decline rate of the nova in the last two years. See text for further details.} \end{figure} \subsection{Sensitivity of the results to the adopted parameters of N. Cygni 1992} Table 1 presents values of system parameters that we used in our calculations. When more than one value is suggested in the literature, the table gives a range of possible values, where the range limits are two particular values that have been suggested for the corresponding parameter. The results of our calculations in Section 2.2 and 2.3 are accordingly given also as interval of possible critical values. For a few parameters of V1974 Cyg, however, there are in the literature more extreme estimates that put the parameter value outside the indicated interval. Here we show that even when using these more extreme values in our calculations, the results do not change very significantly. Balman et al. (1997) suggested an upper limit of 1.37 $M_{\odot}$ on the mass of the WD in V1974 Cyg. Paresce et al. (1995) proposed the low value of 0.75 $M_{\odot}$ for this parameter. The use of these two extreme values in our calculation widens the range of the predicted critical visual magnitude of V1974 Cyg by no more than 0.4 magnitude on each side. Similarly, the upper limit on the distance to this system of 3.2 kpc (Paresce et al. 1995) increases the critical range of magnitudes by about 1.1 mag. It seems, however, that the expansion velocities of the nebulae used by these authors are rather large. All these possible modifications, while changing the calculated time that is required before phase transition occurs, do not alter the qualitative scenario, presented in this work. Nova Cyg '92 is indeed expected to decline further in its optical brightness, as it is presently still nearly four magnitudes brighter than the system pre-nova magnitude $m_{V}=19.5$ (upper horizontal line in Fig. 2 - Annuk et al. 1993). It may have even larger room to decline, if the progenitor had the magnitude $m_{B}=21$ (Fig. 2 - lower horizontal line - Skiff 1997). Naturally, the possibility that V1974 Cyg changes its rate of decline, or even reverses it into brightening, before reaching the critical visual magnitude, cannot be ruled out. In this case PSH will continue to prevail in the LC of this system, much like they do in the LCs of V603 Aql and CP Pup. \subsection{A proposed evolution scenario} Based on the example of the three classical novae V603 Aql, CP Pup and V1974 Cyg, and on the calculations presented in this letter, we may speculate that short period novae with non-magnetic WDs may be the progenitors of PSH systems. This quasi-stable stage in nova life might take a few decades or centuries, before the system returns to its quiescent state, and then becomes a regular SU UMa star. Naturally, such a scenario must be checked quantitatively by a proper statistical analysis of the population of the involved classes of stars. | 98 | 3 | astro-ph9803081_arXiv.txt |
9803 | hep-ph9803394_arXiv.txt | We study the conditions under which an overdamped regime can be attained in the dynamic evolution of a quantum field configuration. Using a real-time formulation of finite temperature field theory, we compute the effective evolution equation of a scalar field configuration, quadratically interacting with a given set of other scalar fields. We then show that, in the overdamped regime, the dissipative kernel in the field equation of motion is closely related to the shear viscosity coefficient, as computed in scalar field theory at finite temperature. The effective dynamics is equivalent to a time-dependent Ginzburg-Landau description of the approach to equilibrium in phenomenological theories of phase transitions. Applications of our results, including a recently proposed inflationary scenario called ``warm inflation'', are discussed. \vspace{0.34cm} \noindent PACS number(s): 98.80 Cq, 05.70.Ln, 11.10.Wx | Kinetic equations describe the time evolution of a certain chosen set of physical variables. The choice of physical variables in principle is arbitrary, but often in practice is governed by the measurement of interest. Typical examples are the order parameter of a complex system or the coordinate of a Brownian particle in a heat reservoir. The kinetic approach is usually implemented through a proper separation of the microscopic equations of motion of the chosen physical variables into regular and random parts. An averaging over the random part then generates the effective partition function for the regular part. This averaging is often referred to as a coarse-graining. One typical application of the kinetic approach is when the physical variables of interest possess energy in relative excess or deficiency to the rest of a large system. Kinetic theory then describes the approach to equilibrium of the chosen physical variables, as for example in the kinetics of phase transitions or in Brownian motion. In the former case, the system is able to release energy to the environment due to some change in its internal state. Provided the environment is disproportionately large relative to the system, the process is irreversible. For a continuous transition, the focus of the present work, this process of equilibration can be described by the monotonic change of an appropriate order parameter, which is the chosen physical variable. Many systems are known to relax in this manner. Phenomenologically, they are successfully described by the time-dependent Ginzburg-Landau theory \cite{gold}. Here we are interested in examining under which circumstances physical variables whose microscopic dynamics is second order in time, as for example the Higgs order parameter of spontaneous symmetry breaking, may have a dynamics which is effectively first order in time as in Ginzburg-Landau phenomenological models. Qualitatively it is not difficult to argue the plausibility of this standard picture for the Higgs symmetry breaking scenario. A single variable, the Higgs order parameter, is modeled to control the release of energy to all the modes that couple to it. By basic notions of equipartition, one anticipates that some portion of the order parameter's energy will flow irreversibly to any given mode. Provided the Higgs order parameter couples to a sufficient number of modes, the motion of the order parameter will be overdamped. In particle physics models, Higgs symmetry breaking is accompanied by mass generation. Thus the natural couplings for the Higgs field $\phi$ to bosonic fields $\chi_i$ is $\phi^2 \chi_i^2$, gauge fields $A^{\mu}_i$ is $\phi^2 A^{i\mu} A_{i \mu}$ and fermionic fields $\psi_i$ is $\phi {\bar \psi}_i \psi_i$. For a microscopic realization of time dependent Ginzburg-Landau theory for the Higgs scalar order parameter in a particle physics setting, these are the most obvious types of couplings to investigate. In this paper we will examine the case of purely bosonic couplings in the ``symmetry restored'' regime. That is, we will study the relaxation of an order parameter which is initially away from the only minimum of the free energy density describing the system. Much of the formalism required for this has already been done in \cite{GR} but we will extend that calculation to the overdamped regime. In an upcoming paper, we plan to study the symmetry broken case. To our knowledge, this paper is the first study of overdamping in quantum field theory with realistic couplings between system and environment, as inspired by particle physics. Overdamping has been studied in quantum mechanical reaction rate theory for a particle escaping from a metastable state (for a review please see \cite{hangii}). This is sometimes referred to as the Kramer's problem, with the overdamped limit also called the Smoluchowski limit. Quantum mechanical models describing this problem are commonly of the system-heat bath type. Microscopic quantum mechanical models have been constructed along these lines, in which the particle (system) is coupled to a set of otherwise free harmonic oscillators (heat bath). Such microscopic system-heat bath models are often referred to as Caldeira-Leggett (CL) models. In many cases they have been exactly solved \cite{fordkac}. The overdamped limit has been derived in these models for the case where the coupling is linear with respect to the oscillator variables but arbitrary with respect to the particle variable \cite{hangii,cl2}. A Caldeira-Leggett type model has also been formulated for the case where the system is a self interacting scalar quantum field coupled linearly to a set of otherwise free fields and the overdamped limit has been obtained \cite{ab54}. This model does provide a microscopic quantum mechanical realization of time-dependent Ginzburg-Landau dynamics in scalar quantum field theory. However, since the couplings between system and environment variables are linear, it should be considered as a first step toward more realistic treatments. More importantly, the calculational method used in \cite{ab54} cannot be extended to the case when the system variable couples quadratically to other fields. Although the analysis of overdamping in this paper has general applicability, it was motivated by the warm inflation scenario of the early universe \cite{ab54,wi}. In \cite{wi} it was realized that the standard Higgs symmetry breaking scenario, when put into a cosmological setting, provides suitable conditions for the universe to enter a de Sitter expansion phase and then smoothly exit into a radiation dominated phase. The overdamped motion of the order parameter in this scenario may sustain the vacuum energy sufficiently long for de Sitter expansion to solve the horizon and flatness problems. Simultaneously, the relaxational kinetics of the order parameter can maintain the temperature of the universe and permit a smooth exit from the de Sitter phase into the radiation dominated phase. Finally, the thermal fluctuations of the order parameter provide the initial seeds of density perturbations, which in addition could be scale free under specified conditions \cite{wi,bf2}. An elementary analysis of this scenario, based on Friedmann cosmology for general realizations of order parameter kinematics, indicated that if the universe's temperature does not fall too much during de Sitter expansion, then the cosmological expansion factor from the de Sitter phase should be of order the lower bound set by observation \cite{ab55}. Although this is not a tight constraint of this scenario, it is a natural one. An analysis of COBE data motivated by this expectation did indicate a slight preference for a small super-Hubble suppression scale, which could be interpreted as arising from a de Sitter expansion with duration near its lower bound \cite{bfh}. Furthermore, the overdamped limit required by warm inflation, when expressed in different terms, was noted \cite{ab54} to be an adiabatic limit, for which known methods from dissipative quantum field theory \cite{GR,hs1,morikawa} are presumed valid. These facts provide further motivation to seek a microscopic model of the scenario, which is the goal of the present work. The calculational methods used here, based on Schwinger's close-time path formalism, were developed in \cite{GR}. There are several other works in the literature that apply this formalism to a variety of different situations. [See, for example, the works of Refs. \cite{hu,boya1,morikawa,greiner,yoko}.] The new feature of the present paper is to shift focus to a kinematic regime dominated by strong dissipation, in order to establish under which conditions this regime leads to overdamped motion. This approach will allow us to have a unique understanding of the microphysical origin of such dynamical behavior, which is in general invoked phenomenologically in applications ranging from condensed matter physics to inflationary cosmology. The paper is organized as follows. In Sec. II our model of interacting bosons is presented and the effective action is computed perturbatively for a homogeneous time dependent background field configuration $\phi(t)$. In Sec. III the effective Langevin-like equation of motion is obtained for $\phi$ in the symmetry-restored phase. In Sec. IV the overdamped limit of this equation of motion is derived and regions of validity are given. In Sec. V the results of the previous sections, which are for Minkowski space, are extrapolated into a cosmological setting and a preliminary examination is made of the warm inflation scenario. In Sec. VI concluding remarks are given. Two Appendixes are included to clarify a few technical details, like the evaluation of the imaginary part of the self-energies and to stress the importance of taking fully-dressed field propagators to properly describe dissipation in the adiabatic approximation for the field configuration. | \label{sec6} In this paper a microscopic quantum field theory model has been presented, describing overdamped motion of a scalar field. Commonly, such behavior is treated phenomenologically by Ginzburg-Landau order parameter kinetics. Our model provides a first principle explanation of how kinetics equivalent to the Ginzburg-Landau type, which is first order in time, arise for inherently second order dynamical systems. The microscopic treatment of this problem, in principle, should be well controlled, due to its fundamental reliance on the adiabatic limit, and our model exemplifies this expectation. The calculational method for treating dissipation in this paper has one distinct difference from several other related works. In our calculation, we consider the effect of particle lifetimes in the effective equation of motion. To our knowledge, this effect has been discussed in only a few works in the past \cite{GR,hs1,morikawa}. Secs. II-IV presented a general, flat-space treatment, which offers a microscopic justification to the often used limit of diffusive Ginzburg-Landau scalar field dynamics. We have shown how it is possible to obtain an effective evolution for the scalar field which is first-order in time, due to its own thermal dissipation effects, interpreted microscopically as its decay into many quanta. In a sense, the field acts as its own brakes, the slowing of its dynamics being attributed to the highly viscous medium where it propagates, a densely populated sea of its own decay products. The application that we considered in Sec. V was in expanding spacetime, for the cosmological warm inflation scenario. Although we did not formally derive the extension of our flat-space model of Secs. II-IV to an expanding background, we did present heuristic arguments that validate this extension for the special needs of warm inflation. The results of the simple analysis in Sec. V are strongly dependent on initial conditions and may be difficult to implement for models of observational interest. Nevertheless, these results will provide useful guidance both for modifications of this model and for our next study of the symmetry broken regime. The direct significance of the present study to inflationary cosmology would be to the initial state problem \cite{muw} in scenarios during symmetry breaking. The initial conditions required for warm inflation in the symmetry broken case are similar to new inflation. The requirement is a thermalized inflaton field, which at the onset of the warm inflation regime is homogeneous with expectation value $\langle \phi \rangle_{\beta}=0$. Although we have made no detailed application of our results to this problem, some general features are evident from the analysis in Sec. V. In particular, both the suppression of large fluctuations and thermalization are mutually consistent with strong dissipative dynamics. Many of the difficulties that have been discussed \cite{muw,cl} in association with the initial state problem, are eliminated in the strong dissipative regime. In addition, the damping of fluctuations should simplify the formal problem of coupling this model to classical gravity. Thus, the strong dissipative regime appears to have the correct features both to carry the universe into an inflation-like phase and then to smoothly exit into a hot big-bang regime. | 98 | 3 | hep-ph9803394_arXiv.txt |
9803 | astro-ph9803204_arXiv.txt | We present a new well defined sample of BL\,Lac objects selected from the ROSAT All-Sky Survey (RASS). The sample consists of 39 objects with 35 forming a flux limited sample down to $\rm f_{X}(0.5 - 2.0\,keV) = 8\cdot 10^{-13}\,ergs\,cm^{-2}\,s^{-1}$, redshifts are known for 33 objects (and 31 of the complete sample). X-ray spectral properties were determined for each object individually with the RASS data. The luminosity function of RASS selected BL\,Lac objects is compatible with results provided by objects selected with the {\em Einstein} observatory, but the RASS selected sample contains objects with luminosities at least tenfold higher. Our analysis confirms the negative evolution for X-ray selected BL\,Lac objects found in a sample by the {\em Einstein} observatory, the parameterization provides similar results. A subdivision of the sample into halves according to the X-ray to optical flux ratio yielded unexpected results. The extremely X-ray dominated objects have higher redshifts and X-ray luminosities and only this subgroup shows clear signs of strong negative evolution. The evolutionary behaviour of objects with an intermediate spectral energy distribution between X-ray and radio dominated is compatible with no evolution at all. Consequences for unified schemes of X-ray and radio selected BL\,Lac objects are discussed. We suggest that the intermediate BL\,Lac objects are the basic BL\,Lac population. The distinction between the two subgroups can be explained if extreme X-ray dominated BL\,Lac objects are observed in a state of enhanced X-ray activity. | The most common view about BL\,Lac Objects is that we are looking into a highly relativistic jet (Blandford \& Rees, \cite{blandford}). The high variability, the polarization, and the spectral energy distribution can in principle all be explained with this model. But there are still unsolved problems for this model. Examples are the nature of the mechanism(s) that generates and collimates the jet. The evolution of physical parameters as the energy and momentum density along the jet is not yet clear, too. An important question is also whether the jets are composed of highly relativistic hadronic or leptonic plasma (Kollgaard, \cite{kollgaard}). Furthermore there exists the competing theory that at least some BL\,Lac objects originate from microlensed QSO (Ostriker \& Vietri, \cite{ostriker}). Models of BL\,Lac objects must explain why X-ray and radio selected BL\,Lac objects differ in some observational characteristics (degree of variability and polarization (Jannuzi et al., \cite{jannuzi})) so that it is not necessarily true that both object classes are of the same astrophysical origin. A sensitive test for the models is the cosmological evolutionary behaviour. To evaluate the luminosity function and evolution behaviour of an object class complete samples with known distances of each individual object are necessary. For BL\,Lac objects two problems exist which render the determination of their luminosity function difficult. First, on account of the inconspicuous optical spectral properties of BL\,Lacs (the absence of strong emission and absorption lines being one of the defining criteria) more effort is required in their selection than is the case with most other object classes. Although there are defining criteria suitable for optical selection (variability, polarisation) no optically selected sample of BL\,Lac objects with more than 10 objects exists (Kollgaard, \cite{kollgaard}). Second, their redshifts are difficult to determine for the same reason. As a consequence, flux limited samples of BL\,Lac objects with nearly complete redshift information are rare. Stickel et al. (\cite{stickel}) selected a sample of 35 objects in the radio wavelength region with radio fluxes above 1\,Jy. They found a flat redshift distribution between $0.1 < \rm{z} < 1.2$ and their analysis of the sample revealed positive evolution; radio selected BL\,Lac objects are more numerous in cosmological distances than in our local neighbourhood. At X-ray wavelengths a complete well defined sample was built up with the EMSS (Stocke et al., \cite{emss}). Morris et al. (\cite{morris}) presented a sample of 22 objects down to $f_{X}(\rm 0.3 - 3.5\,keV) = 5\cdot 10^{-13}\,ergs\,cm^{-2}\,s^{-1}$, which was later expanded to 30 objects and lower fluxes by Wolter et al. (\cite{wolter}). The EMSS objects exhibit a redshift distribution with strong concentration to low redshifts $\rm z < 0.3$ and the sample shows clear signs for negative evolution. The EMSS sample was constructed with serendipitously found BL\,Lac objects from pointed observations of the {\em Einstein} satellite. Because most these pointed observations were of short exposure, objects with low fluxes are less frequent and are rare in the sample. We contribute to this discussion with a new sample consisting of 39 X-ray selected BL\,Lac objects with redshifts available for more than $80\%$ of them. The selection is based on the ROSAT All-Sky Survey (RASS, Voges et al., \cite{rass}). With a flux limit of $f_{X}(\rm 0.5 - 2.0\,keV) = 8\cdot 10^{-13}\,ergs\,cm^{-2}\,s^{-1}$ and redshifts up to $\rm z \sim 0.8$ the sample is suitable to test the evolutionary behaviour of the sample and to determine the luminosity function of RASS selected BL\,Lac objects. In an earlier paper (Nass et al., \cite{nass}) we presented a larger sample to discuss selection processes of BL\,Lac objects within the RASS, but we could not deal with the evolutionary behaviour because the flux limit was not deep enough and many redshifts were unknown. The new sample consists of objects with a wide range of spectral energy distributions, from extremely X-ray dominated objects to objects with intermediate spectral energy distribution. Objects with a maximum of $\nu f_{\nu}$ in the radio or mm wavelength region, which are frequently found among the 1\,Jy sample, are not contained in our X-ray flux limited sample. We were able to subdivide the sample into two subgroups according to \alphaox\footnote{we define \alphaox\ as the power law index between 1\,keV and 4400\,\AA\ with $\rm f_{\nu} \propto \nu^{-\alphaox}$}. Several properties are different in these subgroups. The following section describes the observational characteristics of the sample in the radio-, optical- and X-ray region. Details of the redshift determination are discussed, bcause it is a delicate process for BL\,Lac objects. Section III analyses the properties of the sample, in particular the luminosity function and the evolutionary behaviour. In Sect. IV the new results are discussed in the context of unified schemes for BL\,Lac objects. Several new objects need comments to their redshift determination which are presented in an appendix. We use cosmological parameters $\rm H_{0} = 50\,km\,s^{-1}\,Mpc^{-1}$ and $q_{0} = 0$ in this paper. | We have shown that the unified models do not predict the distinction between the two subgroups in our sample and additional free parameters would be necessary to adjust the model with the observations. Currently the observational constraints are of low statistical significance. Two populations of BL\,Lac objects with different physical origin are also consistent with the results. The dividing point, $\alphaox = 0.91$, is the median of the sample and has no special physical significance. Our study has confirmed the negative evolution of X-ray selected BL\,Lac objects and we find that this behaviour could be restricted to the extremely X-ray dominated BL\,Lac objects. The cosmological evolutionary behaviour of the intermediate objects is consistent with no evolution. Since these two subgroups show a different redshift distribution the interesting evolution time of X-ray dominated BL\,Lac objects is shifted outside to $\rm 1.0 < z < 2.0$ (the $\rm z_{max}$ from the $V/V_{max}$ analysis). There is still no answer for this odd behaviour, which up to now has not been found in other AGN classes. Possibly the jets in these objects, which are especially powerful in the X-rays, need a minimal time for development. This speculation has gained some plausibility since the analysis of our sample has shown that the relevant times lie farther away. Nevertheless X-ray dominated BL\,Lac objects must have been considerably rarer between $\rm 2 < z < 3$, the time when QSO activity was at its highest point. The distinct properties in the two subgroups are already revealed in the bright part of the $\log N(> S) - \log S$ distribution. It is therefore implausible that selection effects, as described by Browne \& March\~{a} (\cite{browne}), are responsible for this distinction. Nevertheless, BL\,Lac objects are ``susceptible'' to selection effects and therefore, in future, galaxy clusters should be carefully analyzed in the selection process. The observations suggest the intermediate BL\,Lac objects as the basic population. They have the lowest luminosity in X-rays and radio wavelengths, and they have the highest space density. Both the extremely radio and X-ray dominated BL\,Lac objects have higher luminosities. This led us to speculate about a beaming scenario. The spectral energy distribution of extremely X-ray dominated BL\,Lac objects can be interpreted with a high energy cutoff of the synchrotron spectrum. The high X-ray luminosity of objects with $\alphaox < 0.91$ can be explained with large bulk Lorentz factor of relativistic electrons in the jets of these objects. This result is compatible with conclusions from multiwavelength observations of outbursts in Mrk\,421 and Mrk\,501. Therefore we suggest that extreme X-ray dominated objects are observed in a state of enhanced activity which would explain the anticorrelation between X-ray luminosity and \alphaox. Our investigation has shown the importance of the redshift parameter for the study of samples of BL\,Lac objects. Advances in observing techniques have made easier the determination of redshifts in optical spectra which are devoid of strong features. In order to have stronger observational constraints it is necessary to enlarge the sample and to determine redshifts for objects for which this has not yet been possible. The uncertainty of the $V/V_{max}$ statistics depends on $1/\sqrt{12N}$. In the last years several projects began to select BL\,Lac objects with combined X-ray - radio methods (e.g. Wolter et al., \cite{pilot}). They will certainly yield valuable results about intermediate objects and low flux BL\,Lac objects below the adopted flux limit in this paper. However, the extremely X-ray dominated BL\,Lac objects are rare and their negative evolution leads to a very flat $\log N - \log S$. Their distinct population properties are an important tool to understand the BL\,Lac phenomenon, and they can be selected most preferably with the full sky coverage of the RASS. | 98 | 3 | astro-ph9803204_arXiv.txt |
9803 | astro-ph9803032_arXiv.txt | We have undertaken toward 30 Mira or semi-regular variables and one OH/IR object highly sensitive observations of the $v=1, J=2 \rightarrow 1$ and $3 \rightarrow 2$ transitions of SiO simultaneously with observations of the $J = 1 \rightarrow 0$ transition of CO during three observing sessions in the period 1995 to 1996. As in our previous observations of 1994, we observe that for several stars the SiO profiles exhibit unusually broad wings which sometimes exceed the terminal velocity of the expanding circumstellar envelope traced by the thermal CO emission. We have discovered a clear dependence of the SiO wing emission on the optical phase. These wings are probably due to peculiar gas motions and varying physical conditions in relation with the stellar pulsation. However, we cannot exclude other mechanisms contributing to the observed wings. In particular, SiO turbulent motions for the semi-regular variables or the asymmetric mass loss mechanism may play a role. We conclude that the SiO wing emission is due to masing processes and that this emission very likely arises from the inner part of the circumstellar envelope. | The SiO molecule exhibits widespread maser emission from O-rich circumstellar envelopes (CEs) around Long Period Variables (LPVs). The variety of rotational transitions emitted from several vibrational levels and the strength of the emission make these lines very useful for a study of the innermost layers of CEs. Recently, during sensitive SiO observations of several O-rich late-type stars, we discovered unexpectedly broad wing SiO emission (Cernicharo et al. 1997, hereafter referred as CABG). We found that the $v=1, J=2 \rightarrow 1$ SiO emission reaches and sometimes exceeds the maximum velocity traced by the quasi-thermal emission of the CO molecule. Several mechanisms may be invoked to explain such wings: turbulent motions, rotation of dense SiO clumps, high velocity shocks produced during the pulsation of the star, or high velocity bipolar ejection of gas from the star. These mechanisms were discussed by CABG, but no firm conclusions were reached, although it was recognized that the pulsation of the star and asymmetric mass loss could play an important role. Obviously interferometric observations and a monitoring of the SiO wing emission are needed to shed light on the location and the physical origin of the high velocity SiO emission. The main purpose of this paper is to present and analyze new data gathered at different epochs on the SiO velocity wings in order to study the physical mechanisms responsible for this phenomenon, and, at the same time, to obtain deeper insight into the complex kinematics of the circumstellar shells around late-type stars. The latter question is clearly important since it is related to other crucial problems such as the expansion of the CE and the mass loss \begin{figure*} [ht] \begin{center} \epsfxsize=15.cm \epsfbox{figure1.epsf} \end{center} \caption []{Examples of SiO $v=1, J=2 \rightarrow 1$ spectra for {\em Mira variables} taken with $0.136$ {\kms} spectral resolution in June 1995 (continuous line), April 1996 (dashed line) and October 1996 (dotted line). For each star the top panel shows the full main beam brightness temperature scale and the bottom panel corresponds to an enlargement of the same data, with the CO line emission width ($\Delta V(CO)$ above the $2\sigma$ level) represented by a horizontal segment. For R Leo, no data are available for April 1996. To transform the intensities into Jy one must multiply by $4.4$ Jy/K.} \label{asio_spectra} \end{figure*} phenomenon, or the processes leading to the formation of dust. We have considered a rather small but homogeneous sample of stars as it includes Miras and a few semi-regulars for which all stellar characteristics (mass loss, temperature, spatial distribution \ldots) are uniformly represented. We have monitored the SiO maser line profiles in order to investigate the relation of broad (weak) emission in wings with the stellar light phase. In particular, we wish to test whether the shocks driven by the stellar pulsation could play a dominant role in the occurrence of SiO high velocity features. In Sects. 2 and 3 we give details of our observations and present our main results. In Sect. 4 we discuss the different hypotheses which could explain the SiO line wings and their location. \newpage | We have observed a rather large sample of late-type stars including Miras and semi-regulars with the IRAM 30-m radiotelescope at four epochs covering the period 1994 to 1996 in order to investigate the correlation between the SiO linewing activity and the stellar light phase. The SiO $v=1, J=2 \rightarrow 1$ and $J=3 \rightarrow 2$ lines were observed simultaneously with the CO $J=2 \rightarrow 1$ quasi-thermal emission line. Several high velocity wings have been detected in the red and blue edges of the SiO profile. The SiO wing emission could result from complex mechanisms combining stellar pulsation of the fundamental mode, and perhaps of other modes, with asymmetric mass loss (as for R Leo) and structure changes. We deduce from our observations that the time evolution of the SiO line wings is related to the stellar pulsation. However, it is difficult to specify how the pulsation induces local physical variations responsible for variations of the SiO wing emission. Pulsation, through shocks, may produce high velocity emissions which then vary according to the optical phase. For the semi-regular variables there is some indication on the importance of turbulent motions in the formation of the high velocity emission as well. Finally, we have shown that the SiO wing emission results from masing processes and that this emission very likely arises from the innermost gas layers of the circumstellar envelope. | 98 | 3 | astro-ph9803032_arXiv.txt |
9803 | astro-ph9803342_arXiv.txt | We present ultra-high resolution (0.32 km s$^{-1}$) spectra obtained with the 3.9m Anglo-Australian Telescope (AAT) and Ultra-High-Resolution Facility (UHRF), of interstellar Na~{\sc i} D$_1$, Na~{\sc i} D$_2$, Ca~{\sc ii} K, K~{\sc i} and CH absorption toward two high galactic latitude stars HD~141569 and HD~157841. We have compared our data with 21-cm observations obtained from the Leiden/Dwingeloo H~{\sc i} survey. We derive the velocity structure, column densities of the clouds represented by the various components and identify the clouds with ISM structures seen in the region at other wavelengths. We further derive abundances, linear depletions and H$_2$ fractional abundances for these clouds, wherever possible. Both stars are located in regions of IRAS 100$\mu$m emission associated with high galactic latitude molecular clouds (HLCs) : HD~141569 lies, in projection, close to MBM~37 and the Lynds dark cloud L~134N while HD~157841 is in the vicinity of the MBM~151. Toward HD~141569, we detect two components in our UHRF spectra : a weak, broad $b$ = 4.5 \kms\ component at -- 15 \kms\ , seen only in Ca~{\sc ii} K absorption and another component at 0 \kms\ , seen in Na~{\sc i} D$_1$, Na~{\sc i} D$_2$, Ca~{\sc ii} K, K~{\sc i} and CH absorption. The cloud represented by the -- 15 \kms\ component, is warm and may be located in a region close to the star. The cloud represented by the 0 \kms\ component has a Ca linear depletion $\delta$(Ca) = 1.4 $\times$ 10$^{-4}$ and shows evidence for the presence of dust, consistent with strong 100$\mu$m emission seen in this region. The H$_2$ fractional abundance $f$(H$_2$), derived for this cloud is 0.4, which is typically what is observed toward HLCs. We conclude that this 0 \kms\ cloud is associated with MBM~37 and L~134N based on the presence of dust and molecular gas (CH) and good velocity agreement with CO emission from these two clouds. This places HD~141569 beyond MBM~37 and L~134N, which are estimated to be at $\approx$ 110 pc. In the case of the HD~157841 sightline, a total of 6 components are seen on our UHRF spectra in Na~{\sc i} D$_1$, Na~{\sc i} D$_2$, Ca~{\sc ii} K, K~{\sc i} and CH absorption. 2 of these 6 components are seen only in a single species. The cloud represented by the components at 1.85 \kms\ has a Ca linear depletion $\delta$(Ca) = 2.8 $\times$ 10$^{-4}$, indicating the presence of dust. The $f$(H$_2$) derived for this cloud is 0.45 and there is good velocity agreement with CO emission from MBM~151. To the best of our knowledge, this 1.85 \kms\ component towards HD~157841 is the first one found to have relative line widths that are consistent with pure thermal broadening only. We associate the 1.85 \kms\ cloud seen in our UHRF spectra with MBM~151 and conclude that HD~157841 must lie beyond $\sim$ 200 pc, the estimated distance to MBM~151. | In this paper, we report the results of an ultra-high-resolution (0.32 \kms) absorption line study of two stars, HD~141569 and HD~157841. Both stars were chosen as background sources for detailed spectroscopic studies of high galactic latitude molecular clouds (HLCs), as well as gas in the nearby halo. HLCs show structure, with dense clumps as small as 0.03 pc, which appear to be virially unstable (Pound et al. 1990). Some of these clouds appear to be embedded in the local hot interstellar medium as suggested by the probable detection of X-ray shadowing from two MBM clouds (Burrows and Mendenhall, 1991). If this is indeed the case, it is not clear what physical process or combination of processes is causing such enhanced molecular abundances. This has led some authors (Blitz, 1990) to conclude that HLCs may represent a completely distinct population of clouds. More recent work by Gir et al. (1994) indicates that HLCs are almost always located along filamentary or looplike HI structures and appear to have condensed from atomic gas in situ rather than having been entrained in the HI. These clouds are an important factor both in the composition and the energy balance of the nearby interstellar medium. They have molecular masses about 50 M$_{\odot}$ and the molecular gas contribution of these high latitude clouds in the solar neighborhood up to 100 pc, is about 5000 M$_{\odot}$ (Magnani 1993). The relation and exchange of energy between the molecular gas in these clouds and the ambient highly ionized, moderately ionized and atomic gas is unclear although it has been suggested by Blitz (1990) that compression of these clouds could lead to a phase transition of atomic hydrogen into molecular hydrogen. Absorption line studies toward stars whose lines of sight intercept regions close to HLCs, probe the diffuse component of the ISM associated with these high latitude clouds. Typical values of the kinetic temperatures for the diffuse component of the ISM range between 45 to 130 K (Savage et al. 1977). To resolve absorption lines arising from clouds with cloud kinetic temperatures of \Tk\ = 50 K, a velocity resolution of 0.30 \kms\ corresponding to a resolving power R = 10$^6$ is necessary. The Ultra High Resolution Facility (UHRF) at the Anglo Australian Telescope (AAT) is currently the world's highest-resolution optical astronomical spectrograph (Diego et al., 1995, Barlow et al., 1995) and is ideal for studies of such cool, interstellar clouds. We have begun a study of the spatial distribution and properties of atomic and molecular interstellar species towards HLCs at ultra high resolution (velocity resolution = 0.32 \kms\ ), and have compiled a list of background stars with reliable photometric and spectroscopic data and lying in the line of sight to HLCs. We have already observed a small fraction of this sample, utilizing the UHRF (Blades et al. 1997), and in the following sections, we present UHRF observations toward HD~141569 and HD~157841. \subsection{Overview of the HD~141569 sightline} HD~141569 has been variously classified as an isolated Herbig Ae/Be star (Th$\acute{e}$ et al. 1994) and a B9.5V star by Penprase (1992) and an AOVe star by Dunkin et al. (1997). It is an emission-line star associated with an IRAS source, IRAS~15473-0346, which is not exactly coincident with it. The heliocentric radial velocity of the star is --6.4 \kms\ (Frisch, 1987). In the direction of HD~141569, which is located in galactic coordinates at ($l$, $b$) = (4.$^\circ$2, +36.$^\circ$9), this translates into a LSR radial velocity of --20.1 \kms\ . The star exhibits excess far-IR emission (Oudmaijer et al. 1992) believed to arise due to emission from dust grains present in a circumstellar disk, and Sylvester et al. (1996) include it in their list of Vega-like systems. A high galactic latitude dark cloud complex at ($l$, $b$) $\sim$ (4$^\circ$, 36$^\circ$) lies close to the star in projection. This cloud complex includes the Lynds dark clouds L~134, L~183 (which is also often referred to as L~134N; henceforth we will refer to this cloud as L~134N) and L~1780. The dense molecular CO core MBM~36 is located within L~134, while the CO core in L~134N is referred to as MBM~37 (Magnani et al. 1985). Of these two dark clouds, L~134N, centered at ($l$, $b$) = (6$^\circ$.0, 36$^\circ$.8) and 2$^\circ$ north of L~134, lies closest to HD~141569. HD~141569 lies well within the 100$\mu$m emission contours of the L~134N dark cloud (Laureijs et al. 1991). Penprase (1992) estimated a distance of 190$\pm$110 pc for HD~141569 based on spectroscopic and UBV and Str\"{o}mgren photometric data. There are various distance estimates to the L134, L134N and L1780 complex: 100$\pm$50 pc, based on optical extinction and surface brightness observations by Mattila (1979), 110$\pm$10 pc based on photometric data by Franco (1989) and 160 pc based on reddenings of a large number of stars in this direction by Snell (1981). The parallax of HD~141569 from the Hipparcos Catalogue is 10.10$\pm$0.83~mas, corresponding to a distance of 99$\pm$8~pc. This implies that the distance to the star is very similar to that of the L~134, L~134N and L~1780 complex. Penprase (1992) estimated an E(B--V) of 0.12 $\pm$ 0.09 for HD~141569. The Hipparcos Catalogue lists Tycho magnitudes which transform to give (B--V) = 0.078$\pm$0.007, which for an A0V star corresponds to an E(B--V) of 0.10$\pm$0.03 (the uncertainty of 0.03 is due to the uncertainty in the (B--V)$_{o}$ calibration for dwarf stars). The star falls just outside the outermost CO contours which trace MBM~37, the CO core in L~134N (Caillault et al. 1995). Based on available information, although HD~141569 appears to lie approximately at the distance of L134N and clearly does not lie within the dark cloud itself, it is not certain whether it lies in front of, or behind this dark cloud. This point is discussed further in $\S$ 3.1.3. and our UHRF data indicate that HD~141569 lies beyond L~134N and MBM~37. \subsection{Overview of the HD~157841 sightline} HD~157841, a B9 star (SIMBAD database) at ($l$, $b$) = (16$^\circ$.8, 15$^\circ$.7), lies in projection near the HLC, MBM~151, which is located at ($l$, $b$) = (21$^\circ$.5, 20$^\circ$.9). We have searched the SIMBAD database and related literature and have been unable to find any radial velocity measurement for this star.\\ Based on CO velocity agreement, Penprase (1993) suggested that MBM~57 located at ($l$, $b$) = (5$^\circ$.1, 30$^\circ$.8) is related to MBM~151, which is $\sim$ 19$^\circ$ away. Photometric data of stars in the region around MBM~57 (Franco, 1989), show an increase in extinction beyond 100 pc. Penprase (1992) also observed this increase in the reddening and estimated a distance of 150 to 210 pc for the near edge of MBM~57. Assuming that MBM~57 and MBM~151 are related, the distance to MBM~151 is $\sim$ 200 pc. The CO (J= 1$\rightarrow$0) LSR velocity of MBM~151 is --0.8 \kms\ (Magnani et al., 1985). The region near HD~157841 has not been mapped in CO and it is not known if there is any CO emission present at the position of the star or in its vicinity. The Hipparcos Catalogue lists a parallax for HD~157841 of 5.70$\pm$1.16~mas corresponding to a distance of 175$\pm$$\stackrel{45}{_{30}}$~pc. It also lists Tycho magnitudes which transform to yield (B--V) = 0.213$\pm$0.008. For a B9V star, this corresponds to an E(B--V) of 0.28$\pm$0.03. In $\S$ 3.2.5., we identify the various components seen in the UHRF spectra toward HD~157841. | We have made an ultra-high resolution study of the Na~{\sc i}, Ca~{\sc ii} K, K~{\sc i} and CH interstellar absorption lines toward two stars, HD~141569 and HD~157841. These absorption spectra have been compared to 21-cm data obtained from the Leiden/Dwingeloo H~{\sc i} survey. Both stars probe the gas in regions close to high galactic latitude molecular clouds : the HD~141569 sightline intercepts a region close to MBM~37 and L~134N while the HD~157841 sightline probes a region close to MBM~151. The results of our investigation are as follows. \noindent 1) Toward HD~141569, two components are seen in our UHRF spectra : one at --15 \kms\ and another at 0 \kms\ . The --15 \kms\ component, seen only in Ca~{\sc ii} K, is weak, broad and shows low linear depletion, $\delta$(Ca) $>$ 0.12. In the absence of turbulence, the kinetic temperature derived is \Tk\ $\le$ 47,900~K and the cloud represented by this component may be located in a region close to HD~141569, whose stellar radial LSR velocity is --20.1 \kms\ . The $\sim$ 0 \kms\ component is seen in Na~{\sc i}, Ca~{\sc ii} K, K~{\sc i} and CH absorption. The cloud represented by these components has a linear Ca depletion of $\delta$(Ca) = 1.4 $\times$ 10$^{-4}$, implying the presence of dust, consistent with strong 100 $\mu$m emission from this region, and H$_2$ fractional abundance $f$(H$_2$) = 0.4. We conclude that this 0 \kms\ cloud is associated with MBM~37 and L~134N because of the presence of dust and molecular (CH) gas at velocities close to CO velocities derived for the two HLCs. The Hipparcos distance of 99$\pm$8~pc for HD~141569 is comparable to the distance of 100 -- 110~pc that has been estimated for the MBM~37 and L~134N clouds. However, our absorption line data imply that these clouds must lie a little closer than HD~141569. \noindent 2) A total of 6 components are seen in our UHRF spectra in Na~{\sc i}, Ca~{\sc ii} K, K~{\sc i} and CH absorption toward HD~157841. The cloud represented by the 1.85 \kms\ absorption components has a Ca linear depletion $\delta$(Ca) = 2.8 $\times$ 10$^{-4}$ and the derived H$_2$ fractional abundance is $f$(H$_2$) = 0.45. The relative $b$-values determined for the 1.85 \kms\ Na~{\sc i}, Ca~{\sc ii} and K~{\sc i} absorption lines are found to be consistent with thermal broadening (with no significant turbulent component), corresponding to a kinetic temperature of \Tk\ = 5700$\pm$500~K. To our knowledge, this is the first velocity component found whose line widths can unambiguously be attributed to pure thermal broadening. We conclude that the 1.85 \kms\ cloud is associated with MBM~151 because of the presence of dust and molecular (CH) gas at velocities close to CO velocities derived for MBM~151. HD~141569 must therefore lie beyond MBM~151, estimated to be at $\sim$ 200 pc and its Hipparcos distance of 175$\pm$$\stackrel{45}{_{30}}$~pc is consistent with this conclusion. | 98 | 3 | astro-ph9803342_arXiv.txt |
9803 | astro-ph9803174_arXiv.txt | The cluster Abell~370 is a distant ($<z>=0.374$~\cite{M88}), rich, massive cluster, with two central giant galaxies dominating its optical image. The X-ray luminosity of the cluster is\footnote{We assume a Hubble constant $H_0=75$~km~s$^{-1}$~Mpc$^{-1}$ and a deceleration parameter $q_0=1/2$ throughout this paper, giving a luminosity-distance to the cluster of 1613~Mpc.} $3.7 \times 10^{44}$ erg~s$^{-1}$ \cite{Henry}, which is as expected from the high value of the velocity dispersion of this cluster, 1340~km~s$^{-1}$\cite{M88}. A370 was shown by \cite{Bautz}, \cite{BO84} and \cite{M88} to contain an anomalous fraction of blue objects, as compared to nearby clusters (50~\% of the cluster members in the central region show evidence for star-formation activity, as compared to a mere 3~\% in Coma). Based on optical and near-IR photometry, it has been shown (\cite{Aragon}, \cite{Stanford}, \cite{McLean}) that most galaxies in the cluster are reasonably well fitted by passive evolution models, while the existence of galaxies in a post-starburst phase is controversial (\cite{Aragon} vs. \cite{Stanford}). The spectral energy distribution of the cluster galaxies must be known over a wider wavelength range before their nature can be firmly established, and in this context ISO observations are very important. {\psfig{file=a370_opt_c.ps,width=15.cm,angle=0}} {\bf Fig.1.} I-band optical image of A370 (from \cite{Kneib}); the original image has been transformed to match the ISOCAM image field-of-view. The coordinates are arcsec. The strange feature in the upper part of the figure is straylight from a very bright star off the image. \medskip A370 contains the first spectroscopically confirmed gravitational arc, A0 at $\mbox{z=0.724}$~\cite{Soucail}. Other arcs have since been detected in the cluster (\cite{Kneib} and references therein). Based on {\em Hubble Space Telescope} images, \cite{Smail} found evidence for a faint spiral structure and an apparent bulge. The spiral nature of the A0 arc was also indicated by its optical and near-IR colours and it was detected by its CO line in emission by \cite{Casoli}. In this paper we report on the mid-IR observations of A370 done with ISOCAM onboard ESA's {\em Infrared Space Observatory} (ISO) satellite. A370 and another three clusters were selected for observation in the context of an ISO guaranteed-time project aimed at generating mid-IR images of known giant arcs and providing their mid-IR fluxes. In \S~2 we give a description of our observations and data-reductions; in \S~3 we present the ISOCAM image, and compare it with an I-band image in order to derive a colour map. We give our conclusions in \S~4. | In this paper we have presented the results of observations of A370 with ISOCAM onboard ISO. We have detected the giant arc at 7 and 10 $\mu$m; this is the first detection ever of a gravitational arc in a galaxy cluster at these wavelengths. Moreover, we have also detected many other galaxies, in the cluster field. The comparison of an I-band image and the ISOCAM 7 $\mu$m image in the LW2 filter has allowed us to build a colour-map which seems to indicate a starbursting phase for the gravitationally lensed galaxy. Most cluster galaxies however appear to have a ratio of 7 $\mu$m to 0.9 $\mu$m flux density as predicted from spectral energy distribution models of normal ellipticals \cite{Mazzei1}. While our data reduction methodology is still being developed and refined, and while a full exposition of quantitative results remains to be achieved, the present results already demonstrate the capability of ISO in detecting faint mid-IR sources at large distances. With our ongoing program we hope to detect other gravitational arcs in clusters, and so to constrain the evolutionary status of these high redshift sources. We will widen the wavelength coverage of the spectral energy distributions for cluster galaxies, thus providing a better leverage for constraining models of galaxy formation and evolution. An additional exciting perspective of our future observations is the serendipitous discovery of optically undetected "IR-only" arcs or arclets. | 98 | 3 | astro-ph9803174_arXiv.txt |
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9803 | astro-ph9803319_arXiv.txt | The dispersion in the peak luminosities of high redshift type Ia supernovae will change with redshift due to gravitational lensing. This lensing is investigated with an emphasis on the prospects of measuring it and separating it from other possible sources of redshift dependent dispersion. Measuring the lensing induced dispersion would directly constrain the power spectrum of density fluctuations on smaller length scales than are easily probed in any other way. The skew of the magnification distribution is related to the bispectrum of density fluctuations. Using cold dark matter models it is found that the amount and quality of data needed is attainable in a few years. A parameterization of the signal as a power law of the angular size distance to the supernovae is motivated by these models. This information can be used in detecting lensing, detecting other systematic changes in supernovae and calculating the uncertainties in cosmological parameter estimates. | There is presently a large effort underway to predict and detect the weak gravitational lensing caused by Large Scale Structure (LSS) or the ``cosmic shear'' (see \cite{valdes83,mould94,vill96,kais92,kais96}). Such a measurement would constitute a direct probe of the mass density fluctuations on large scales irrespective of how light and baryons are distributed. This would make it possible to measure one of cosmology's least well understood processes, how light traces mass and whether this is a function of scale. Gravitational lensing causes images to be both magnified and demagnified as well as stretched asymmetrically (shear). Most of the methods proposed for detecting LSS lensing are based on measuring the shear in high redshift galaxy images. Although the distortions in the ellipticities of individual galaxies are expected to be small they are distorted in coherent ways. Indications of lensing are sought in the alignment of the galaxy images with the assumption that they are not intrinsically aligned. This technique has already been used with great success on galaxy cluster lensing and is presently being applied to random fields in an attempt to detect the lensing effects of large scale structure. In this paper a method of detecting lensing directly through its magnification rather than shear is proposed. The study of high redshift supernovae (SNe) is another area of cosmology and astrophysics that is seeing a lot of activity. Type Ia SNe, the brightest type of supernova (SN), are believed to be caused by the thermonuclear explosion of an accreting white dwarf. It has been found empirically that the peak magnitude of type Ia's have a dispersion of only 0.2 - 0.3 mag in B band. It has further been found that the peak magnitude is related to the width of the SN's light-curve which can then be used to reduce the dispersion to about 0.17 mag if one color is used \cite{hamuy96} and 0.12 mag if multiple colors are used \cite{RPK96}. In addition to the light-curve width, the SN's color and spectral features are related to the peak luminosity \cite{BNF97,nugent95}. It may be possible to reduce this dispersion in the future by incorporating additional observables into the correction procedure. The uniformity in type Ia SNe combined with their high luminosity makes them an excellent tool for doing cosmology. Using them to measure the redshift-luminosity distance relation has recently resulted in tightened constraints on the cosmological parameters $\Omega$ and $\Omega_\Lambda$ (\cite{Perl97,Perl98,garn98,Riess98}). There are now systematic searches for Ia SNe at high redshift which can reliably discover on the order of ten SNe in a night's observing and do spectroscopic followup (\cite{Perl97}). In addition there are several ongoing searches for low redshift SNe. To date there have been more then 100 type Ia SNe discovered with redshifts between $z=0.4$ and $0.97$ and many additional ones at lower redshift. In this paper I concentrate on the lensing produced by dark matter composed of microscopic particles. Matter in macroscopic compact objects can cause microlensing. However, the known populations of stars will not cause a significant number of microlensing events. The microlensing of SNe has been discussed by \cite{SW87} and \cite{LSW88}. A SN could also be lensed by one dominant galaxy cluster or individual galaxy. There is the possibility of getting multiple observable images and high magnifications (strong lensing) in this case. It has been suggested that observing SNe behind galaxy clusters would be a way of lifting the mass sheet degeneracy that exists in shear measurements of the gravitational lensing \cite{kolatt97}. The likelihood of a SN at $z=1$ being strongly lensed by a galaxy or cluster is small unless they are specifically sought out. In general a SN will be lensed by many galaxies and larger structures, each have a weak contribution to the total magnification. This will increase the dispersion of high redshift SNe and decrease the precision of cosmological parameter determinations. This decrease in precision has been investigated by \cite{Frieman97} using analytic methods and by \cite{wamb97} using N-body simulations. \cite{HW98} calculates the lensing of point sources using numerical simulations which assume that all matter is in unclustered galaxy halos. \cite{Kant98} (and \cite{Kant95}) does analytic calculations of this effect under the assumption that all matter is in compact objects and none of them are close enough to the line of sight to a SN to cause significant lensing. In this paper the problem will be turned around and we will ask how well the lensing itself can be determined from the SN data. There are several important differences between the lensing of SNe and the lensing of galaxies. For SNe the signal to noise ratio in each measurement can in fact be greater. Lensing is estimated to contribute about $5$ to $10\%$ to the observed root-mean-squared ellipticity of a $z=1$ galaxy. For SNe the variance in the lensing contribution is comparable to the intrinsic variance in the peak magnitude at this redshift. In addition the unlensed ellipticity distribution of galaxies at high redshift is not known and can not be easily extrapolated from zero redshift galaxies. For this reason lensing must be inferred by correlations between galaxies, either between lensed galaxies or between lensed galaxies and foreground galaxies. The result is that the lensing of galaxies measures the shear averaged over a finite area on the sky. This average shear drops rapidly with increasing area and the signal to noise is reduced to something more like $1\%$ per galaxy on the $1\mbox{ deg}^2$ scale. This is made up for by the large number of evaluable galaxies ($\sim 10^5\mbox{ deg}^{-1}$) to the extent that fluctuations in the projected mass density at $1-100$~arcmin scales are expected to be detectable in the near future. In contrast, the dispersion in the absolute magnitude of type Ia SNe is presumably independent of redshift. The dispersion can then be measured in a low redshift population where lensing is not important. The lensing of galaxies does have the advantage of sources that are generally at higher redshifts where the lensing is stronger. On the other hand, the redshift distribution of faint galaxies is not strongly constrained which adds systematic uncertainty. The redshifts of SNe are individually measured. It is clear that the lensing of SNe and the lensing of galaxies probe different scales of density fluctuations for several reasons. Because the lensing of galaxies can only be detected through correlations in their shear, or positions and shear, lensing structures that are smaller than roughly the separation between galaxies are not detectable. Since each SN is an independent measure of the lensing at a point, not an average over a region on the sky, they will be sensitive to smaller scale structures. Also the lensing of SNe is a direct measurement of the magnification which, in the weak lensing limit, is directly related to the mass density along the line of sight. The shear is dependent on the mass outside of the ``beam'' and thus a shear map is in a way a smoothed version of a surface density map. In the thin lens approximation the shear and the magnification are related to each other through a differential operator \cite{KS93} which makes the shear insensitive to a uniform offset in the surface density - the ``mass sheet degeneracy''. Although the lensing of galaxies has limitations on small scales it does have the potential of measuring lensing over a large range of scales, arcminutes to degrees. With the possible exception of SNe viewed through galaxy clusters, it will be difficult to find enough SNe that are close enough together to measure correlations in their lensing. In the next section I describe first how the lensing of SNe is related to the density fluctuations and then I describe how the lensing signal could be identified in the data. In section~\ref{models} a specific model of structure formation is used to estimate the level of signal and to motivate some parameterizations. The last section contains conclusions and remarks about possible complications. | It has been shown that measuring the gravitational lensing of type Ia SNe is feasible if the noise can be reasonably constrained. It would be best to solve for the best fit cosmological parameters (ie. $\Omega$, $\Omega_\Lambda$), lensing strength (ie. $\eta_o$) and intrinsic noise ($\sigma_M$) simultaneously using SNe at all redshifts. The photometric uncertainties should be comparatively well determined for each SN. One could then marginalize over the intrinsic variance, $\sigma_M$. The greatest worry is of course that the type Ia SNe properties or their galactic environments are changing with redshift. Observations of spectral features and colors (\cite{Perl97}) suggest that this is not the case, but the possibility remains. It is possible that a systematic change in the metallicity of progenitors could change both the average peak luminosity and its dispersion. Another worry is that extinction corrections change with redshift. This could systematically reduce $\overline{m}(z)$, increase its dispersion and make its distribution non-Gaussian. Extinction should be accompanied by reddening, which can be detected, but the extinction law is not certain. These changes would affect cosmological parameter parameter estimates as well as lensing estimates. The methods discussed here can be directly applied to detecting any redshift dependent change in the dispersion. In testing for possible redshift evolution lensing must be incorporated. There are difficulties in searching for SNe at higher redshifts. The region of the spectrum that is used to do the light curve correction passes out of the visible at $z\gtrsim 1$. To go to significantly higher redshift may require switching to the IR. There is also difficulties with the K-correction and the subtraction of atmospheric lines. But with this in mind it seems that gravitational lensing of SNe could be detectable in the next few years when hundreds of high redshift SNe are observed and systematic effects are better understood. CCD cameras with fields of view approaching a square degree and very small pixel sizes are being built now. They will be used for weak lensing measurements using galaxy shear. SNe searches could be incorporated into these surveys with the benefits of improved cosmological parameter estimation and complimentary weak lensing measurements. Combined with the limits on the cosmological parameters the lensing of SNe can constrain the power spectrum of the true mass density, unbiased by the light distribution, on the scale of galaxy halos. At present the mass distribution on these scales is not well known with the exception of within galaxy clusters which are certainly atypical regions. In this paper it has been assumed that the the large majority of matter in the universe is in the form of WIMPS or some other small particle. If the dark matter is in compact objects like MACHOs the distribution of magnifications will be different. In this way the lensing of SNe also provides information on the composition of dark matter. In addition to the variance of the magnification distribution the skewness would provide an important constraint on the nature of structure in the nonlinear regime. A future paper will treat this subject in more detail and relate the magnification distribution to the nature of dark matter and the structure of galaxy halos. | 98 | 3 | astro-ph9803319_arXiv.txt |
9803 | astro-ph9803296_arXiv.txt | We present new observations of the isolated young stars \hd\ and \cd. Pointed \rosat\ observations show that their X-ray properties, including X-ray luminosity and variability, are consistent with those of \pms\ (PMS) stars. These observations do not reveal any additional PMS candidates in 40\arcmin\ fields centered on \hd\ and \cd. Hipparcos observations of TW Hya (Wichmann \etal\ 1998\mc{w98}) and \hd\ (Soderblom \etal\ 1998\mc{s98}) show that both stars are roughly 50 pc away and are PMS with ages of $\sim 10^7$ yr. We searched the Hipparcos catalog (complete down to $\sim$ 2--3 $L_\odot$ at this distance) for other PMS stars in the same area. In a 10-pc radius volume of space centered on the previously known PMS stars, we find one additional candidate PMS star (CD $-$36$\arcdeg$7429) with a low space velocity, X-ray emission comparable to that of \hd, and Li absorption. There are eight other stars in this area that have dwarf spectral types and lie above the main sequence, but based on their weak X-ray emission, high space velocities, and lack of Li in low-resolution spectra (i.e.\ EW(Li) $< 0.1$ \AA), these are probably mis-classified subgiants or giants. The current positions and proper motions of TW Hya, \hd, and CD $-$36$\arcdeg$7429 are inconsistent with them having formed as a group. | Recent observations have revealed a small population of stars that bear many of the hallmarks of low-mass pre--main-sequence stars but lie far from any obvious region of recent star formation (as revealed by substantial dark clouds of molecular gas and dust). Gregorio-Hetem \etal\ (1992\mc{gh92}) identified 33 candidate T Tauri stars based on spectroscopy of stars in the {\it IRAS\/} Point Source Catalog and a few additional emission-line stars. One of the more interesting findings of their work was the identification of four stars (HD 98800, \cd, \cdb, and \hen) within 10\arcdeg\ of TW Hya, earlier identified as a possible isolated T Tauri star by Rucinski \& Krautter (1983\mc{rk83}). The proximity of these five systems to each other suggested a possible loose cluster of young stars, possibly formed by a small molecular cloud that has since dissipated (e.g., Feigelson 1996\mc{f96}). In the absence of association with any known cloud complex, the distances to these stars were very uncertain, and thus their pre--main-sequence status was in question. To investigate whether two of these stars, \hd\ and \cd, are in fact young and whether they are part of a larger group of young stars, we observed them with the \rosat\ X-ray satellite. One distinguishing feature of low-mass pre--main-sequence stars is strong X-ray emission, presumed to arise from solar-like chromospheric activity (see, e.g., \rne\ 1997\mc{n97} for a recent review). Thus, X-ray observations can be used to search for young stars that may not have other hallmarks of youth such as infrared excesses or strong H$\alpha$ emission. Observations of Taurus-Auriga with the {\it Einstein} satellite led to the discovery of the naked T Tauri stars, stars that have no circumstellar material but that are nonetheless coeval with classical T Tauri stars (Walter 1986\mc{w86}). Much observational attention has been focused on the stars in the vicinity of TW Hya recently, and several other investigations have proceeded in parallel with ours. Kastner \etal\ (1997\mc{k97}) reported \rosat\ observations of some of the same stars we report on here. They found that the strength of X-ray emission from these stars is consistent with them being young stars and argued that the five young stars in this area make up a physical association. Hoff \etal\ (1996\mc{hph96}, 1998\mc{hhp98}) investigated TW Hya and \cdb\ and their surroundings with low spatial resolution using \rosat\ PSPC pointed observations. Soderblom \etal\ (1998\mc{s98}) and Wichmann \etal\ (1998\mc{w98}) report Hipparcos distances to HD 98800 and TW Hya, respectively; both systems have ages of $\sim$ 1--2 $\times\ 10^7$ yr and distances of $\sim 50$ pc. The work we report here is complementary to this other recent work. We analyze the X-ray data of our target stars in more detail than Kastner \etal\ (1997\mc{k97}) and explore the status of other X-ray sources in the surrounding fields. We also combine X-ray and Hipparcos data to search for other young stars in the vicinity in order to address the larger question of the origin of these isolated young stars. In \sec\ \ref{sec:observations} we report the details of our observations. In \sec\ \ref{sec:xproperties} we show that the X-ray properties of the two target stars are consistent with their being \pms\ (PMS) stars. We then investigate (\sec\ \ref{sec:othermembers}) whether there is any evidence that the young stars near TW Hya formed as a group, using our X-ray observations as well as data from the Hipparcos satellite. Finally, in \sec\ \ref{sec:discussion} we discuss the implications of our findings for our understanding of isolated young stars. | \label{sec:summary} We have presented pointed \rosat\ observations of the isolated young stars \hd\ and \cd. Their X-ray fluxes and X-ray variability are consistent with them being \pms\ stars. No other \pms\ stars were found among X-ray sources within 40\arcmin\ of these stars. Hipparcos observations of the area reveal the presence of nine additional stars that lie above the main sequence in a volume roughly 10 pc in diameter centered on the five previously known young stars. The X-ray properties and kinematics of these stars indicate that one of them (CD $-$36$\arcdeg$7429) is quite likely a PMS star with an age of $10^7$ yr, while the others are more likely post--main-sequence stars. Observations of Li abundances in these stars confirm this conclusion. Comparison with other fields at the same Galactic latitude shows that other fields selected in the same way show similar numbers of stars above the main sequence. However, the field around TW Hya has significantly more stars with low space velocities (less than 5 km s\per\ relative to the LSR) than the other fields. Thus, there is some indication of an excess of young stars in this area. Nonetheless, the proper motions of \hd, TW Hya, and CD $-$36$\arcdeg$7429, if extrapolated back in time, do not indicate a common place of origin. \noindent \centerline{\bf Note added in proof:} Our analysis of whether HD 98800, CD $-$36$\arcdeg$7429, and TW Hya could have formed as a group neglected an important source of uncertainty. The uncertainty in distance to these stars, while it does not affect the observed proper motions, does introduce uncertainty into the solar motion that must be subtracted in order to determine the stars' intrinsic proper motions with respect to the LSR. The 1$\sigma$ effect of this uncertainty is shown by the dashed lines in Figure 3. Including this uncertainty indicates that TW Hya and CD $-$36$\arcdeg$7429 could plausibly have formed together; only a $1.5\sigma$ error on their distances is required. Regarding HD 98800, bringing its current distance $2\sigma$ closer, to 36.8 pc, would allow its projected point of origin to overlap with those of the other two stars. However, this distance makes HD 98800 a main-sequence star, in conflict with its observed properties and, more importantly, of a very different age than the other two stars. We conclude that HD 98800 is unlikely to share a common origin with TW Hya or CD $-$36$\arcdeg$7429. | 98 | 3 | astro-ph9803296_arXiv.txt |
9803 | astro-ph9803225_arXiv.txt | Results obtained from 9 X-ray observations of 3C 273 performed by {\it ASCA} are presented (for a total exposure time of about 160 000 s). The analysis and interpretation of the results is complicated by the fact that 4 of these observations were used for on-board calibration of the CCDs spectral response. In particular, we had to pay special attention to the low energy band and 5--6 keV energy range where systematic effects could distort a correct interpretation of the data. The present standard analysis shows that, in agreement with official recommendations, a conservative systematic error (at low energies) of $\sim$ 2--3 $\times$ 10$^{20}$ cm$^{-2}$ must be assumed when analyzing {\it ASCA} SIS data. A soft-excess, with variable flux and/or shape, has been clearly detected as well as flux and spectral variability that confirm previous findings with other observatories. An anti-correlation is found between the spectral index and the flux in the 2-10 keV energy range. With the old response matrices, an iron emission line feature with EW $\sim$ 50--100 eV was initially detected at $\sim$ 5.6-5.7 keV ($\sim$ 6.5-6.6 keV in the quasar frame) in 6 observations and, in two occasions, the line was resolved ($\sigma \sim$ 0.2-0.6 keV). Comparison with the Crab spectrum indicates however that this feature was mostly due to remaining calibration uncertainties between 5--6 keV. Indeed, fitting the data with the latest publicly available calibration matrices, we find that the line remains unambiguously significant in (only) the two observations with lowest fluxes where it is weak (EW $\sim$ 20-30 eV), narrow and consistent with being produced by Fe K$_{\alpha}$ emission from neutral matter. Overall, the observations are qualitatively consistent with a variable, non-thermal X-ray continuum emission, i.e., a power law with $\Gamma$ $\sim$ 1.6 (possibly produced in the innermost regions of the radio-optical jet), plus underlying ``Seyfert-like'' features, i.e., a soft-excess and Fe K$_{\alpha}$ line emission. The data are consistent with some contribution (up to a few 10\% level in the {\it ASCA} energy band) from a ``Seyfert-like'' direct continuum emission, i.e. a power law with $\Gamma$ $\sim$ 1.9 plus a reflection component, as well. When the continuum (jet) emission is in a low state, the spectral features produced by the Seyfert-like spectrum (soft-excess, iron line and possibly a steep power law plus a reflection continuum) are more easily seen. | The remarkable discovery by EGRET on-board {\it CGRO} that blazars (i.e. BL Lacertae objects and flat-spectrum radio quasars) are strong $\gamma$-ray emitters has drawn in recent years the attention of the astronomical community to this class of objects. Observations indicate that the overall energy distribution of blazars shows the signature of two different types of emission mechanisms: beamed non-thermal jet radiation producing the overall broad band (from radio to $\gamma$-rays) continuum emission common to all blazars, and quasi-isotropic thermal radiation by an accretion-disk producing a ``UV Bump'' observed in a large number of quasars and Seyfert galaxies but absent in BL Lac objects (e.g. Sambruna, Maraschi \& Urry, 1996, Elvis et al. 1994). The non-thermal continuum consists of IR-optical and $\gamma$-ray peaks. The first peak is interpreted in terms of synchrotron emission and the second peak in terms of inverse Compton emission (see Urry \& Padovani 1995 for a review on the subject). From object to object, the observed different spectral characteristics may be due to the relative importance of one emission mechanism to the other which, in turn, is likely to be related to the amount of beaming in one object or the other (e.g. Dondi \& Ghisellini 1995). One of the most well-studied and characteristic example of blazars is the bright quasar 3C 273 (z$\simeq$0.158). It is a good example where both non-thermal and thermal emission mechanisms might be observed because its broad band energy distribution exhibits two large peaks, one peaking in the IR-optical and one peaking in the $\gamma$-rays with a `UV bump'' superimposed on it (Courvoisier et al. 1987, Lichti et al. 1995, von Montigny et al. 1997). The study of its X-ray properties may provide important clues to understand the origin of both emission mechanisms because a) soft-X-ray excess emission has been observed and interpreted as the high-energy tail of the UV bump (Turner et al. 1985, Courvoisier et al. 1987, Walter et al. 1994, Leach, Mc Hardy \& Papadakis 1995) and b) the 2-10 keV spectrum which is most likely associated to the $\gamma$-ray emission is known to be variable in time and shape (Turner et al. 1985), thus allowing for tests of different X- and $\gamma$-ray emission models. \pn To clarify the mechanism responsible for the X-ray emission, high quality data are first necessary to disentangle the contributions from the different spectral components, namely the jet and Seyfert components. In this paper, we report on observations with {\it ASCA} during the first year of the mission. The source spectral properties are shown with particular attention to calibration uncertainties, most relevant in this source because it was used for on-board calibration of the CCDs. We show evidence of complex spectral features (soft-excess, Fe K emission line, flux and spectral correlated variability) and discuss their possible interpretation. Throughout the analysis we use $H_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$ and $q_{0}$ = 0. | To date, {\it ASCA} has observed 3C 273 10 times. Results from the first 9 observations, all performed during the first year of the mission have been presented here. These confirm and expand the evidence that the X-ray emission of 3C 273 is complex, with different spectral components contributing to its X-ray emission. Because 4 of the 9 observations were used for the on-board calibration of the CCDs, great care had to be taken when interpreting the observational results for this source. As a rule, {\it absolute} values, in particular those obtained from the SIS, require a detailed estimate and knowledge of the instrumental systematic errors to be trusted. {\it Relative} measurements (like flux and/or spectral variability) obtained from comparing different observations are, however, more reliable since they should not be affected by calibration uncertainties. With this caveat in mind, it is found that: \pn \begin{enumerate} \item A conservative systematic error at low energies that corresponds to an extra-absorption column of $\sim$ 2--3 $\times$ 10$^{20}$ cm$^{-2}$ is found for the SIS response, consistent with the ASCA Team's official prescriptions. \item 2--10 keV flux variations by up to $\sim$ 60\% on a time-scale of $\sim$ 200 days and day-to-day variations as large as $\sim$ 20\% were observed. \item Extra soft X-ray emission is required by the data during the first observation, when the source was in its lowest flux level. \item Flux and spectral slope variations are clearly detected as well and, for the first time, there is a statistically significant evidence that the index and flux are anti-correlated. \item Iron line emission is detected in (only) the two observations with the lowest flux levels. The line is in both cases weak (EW$\sim$20-30 eV), but statistically significant at more than 99\% confidence level, narrow and consistent with Fe K$_{\alpha}$ emission from neutral matter. \end{enumerate} We then speculate that all the above observable properties of the X-ray spectrum of 3C 273 can be interpreted in terms of the sum of two emission mechanisms. These are a non-thermal emission from the innermost regions of the jet which dominates the 2-10 keV region and whose signatures are the spectral variability and a flat ($\Gamma \sim 1.6$) power law continuum (that extrapolates well into higher energies), plus a diluted Seyfert-like spectrum whose signatures are the soft-excess and iron line emission. The newly discovered index-flux anti-correlation may be interpreted either by intrinsic variations of the jet power law index or by some contribution of the Seyfert-like continuum spectrum (say, a power law with $\Gamma \sim 1.9$) as the jet component varies. The overall scenario predicts that when the (dominant) jet component is in a low flux state, the spectral features produced by the Seyfert-like spectrum should be more easily detected (owing for variability of the Seyfert-like spectrum itself). | 98 | 3 | astro-ph9803225_arXiv.txt |
9803 | astro-ph9803155_arXiv.txt | We report on the mid-infrared imaging at 5, 7, 10 \& 15 $\mu$m of the galaxy cluster Abell 2218 obtained with the ISOCAM instrument onboard ESA's Infrared Space Observatory (ISO), as part of an on-going program to image gravitational arcs and arclets in distant clusters. Several cluster galaxies as well as field galaxies are detected. We discuss their mid-IR flux properties. | Abell 2218 is a very rich galaxy cluster (richness class 4, according to Abell et al. 1989), characterized by a large velocity dispersion of the galaxy population ($\sigma_v=1370$~km/s, Le Borgne et al. 1992), and a high X-ray temperature and luminosity ($T_x=7.2$~keV, see e.g. Markevitch 1997, $L_x = 2.9 \times 10^{44}$ erg/s\footnote{In this paper we adopt $H_0=75$~km/s/Mpc and $q_0=1/2$} in the energy band 0.5--4.4 keV, see e.g. Kneib et al. 1996). These properties, coupled to a relatively small distance ($<z>=0.175$, Le Borgne et al. 1992), made this cluster an attractive target for studies of the Sunyaev-Zeldovich effect (Birkinshaw \& Hughes 1994), and one of the closest clusters where gravitational arcs are detected. The amazing concentration of gravitational arcs and arclets has stimulated a huge observational effort first to get sub-arcsec imaging (from the ground, see, e.g. Kneib et al. 1995, and from space with HST, Kneib et al. 1996), and then to get spectra of the arcs (Ebbels et al. 1997). The optical and near-IR observations have allowed a very detailed modelling of the mass distribution within the cluster (Kneib et al. 1996), confirmed later to a great level of accuracy, (Ebbels et al. 1997). It was found that the cluster mass distribution is bi-modal, with the main concentration centred on the cD galaxy. The X-ray surface brightness, as obtained via ROSAT observations, does not trace the gravitational potential as derived from the lensing analysis, a possible indication that the X-ray emitting gas is far from hydrostatic equilibrium (Kneib et al. 1996). This would also explain the discrepancy in the X-ray and lensing mass estimates (Markevitch 1997). In this paper we report on the mid-IR observations of Abell 2218 done with ISOCAM on-board ESA's {\em Infrared Space Observatory} (ISO) satellite. The high sensitivity of ISOCAM allows us to determine the photometric properties of lensed and cluster galaxies in the mid-IR band, thus widening the wavelength coverage of their spectra. Mid-IR observations are critical in understanding the intrinsic nature of these distant galaxies, since a starburst is known to emit a large fraction of its energy in the mid-IR, due to dust grain re-processing of the emitted radiation. | The LW1, LW2 and LW3 maps are presented in figures 1 \& 2a, as overlays on top of the HST image (courtesy J.-P. Kneib) of Abell 2218. \begin{figure} \plottwo{altierib1.eps}{altierib2.eps} \caption{Abell 2218 4.5$\mu$m (a) \& 7$\mu$m (b) contour maps} \label{fig-1} \end{figure} In the cluster core, no arclet is detected at a significant level at any wavelength in the mid-IR. This non-detection of arcs or arclets contrasts with our recent mid-IR imaging of Abell 370 (Metcalfe et al. 1997), where the A0 giant arc is clearly detected as the main feature and emitter in the cluster core at 15$\mu$m. But this giant arc is already a prominent feature in the optical, (first giant arc discovered), whereas arcs and arclets are fainter in Abell 2218. \\ The cD galaxy which is apparently centred on the cluster potential is clearly detected in the LW1 \& LW2 filter bands. At 4.5$\mu$m (rest-frame wavelength 3.8 $\mu$m) the cD emission is extended, covering part of the optical halo; at 7$\mu$m the emission is much more confined to the centre, but with an extension along the optical major axis, probably contaminated by the neighbouring merging dwarf galaxies. At 10$\mu$m and 15$\mu$m it lies just above the detection limit of ISOCAM and no statement can be made on its extension. The mid-IR spectral energy distribution up to 15$\mu$m seems to follow a simple Rayleigh-Jeans tail of the cold stellar component as found in optically-selected normal early-type galaxies (E, S0, S0a) in the Virgo cluster (Boselli et al. 1998). But due to the contamination by close-by galaxies it is very difficult to estimate the 4.5$\mu$m flux of the cD. \\ The brightest cluster member galaxies are also detected at 4.5 $\mu$m and 7 $\mu$m. At the redshift of the cluster the 4.5 $\mu$m emission is mostly coming again from cold stellar photospheres. The 7 $\mu$m (LW2 band) (rest-frame 5.7 $\mu$m) emission includes a small fraction of PAH emission but comes mainly again from the cold stellar photospheres. At 10 \& 15 microns most of cluster members have vanished (ie. are below or very close to our detection limit), as expected for normal early-type galaxies (ellipticals) where there is only the Rayleigh-Jeans tail mid-IR emission from cold stellar photospheres, in dust \& gas poor cluster core ellipticals.\\ Still one cluster member galaxy \#373 (numbering system from Le Borgne et al. 1992 in the following) is detected at 15$\mu$m, but also at 4.5 $\mu$m and 7 $\mu$m, it has a 'cart-wheel' like aspect, from the ring surrounding it and is probably a big face-on spiral galaxy. The 4.5$\mu$m and 7$\mu$m emission comes mostly from the nucleus whereas the 15$\mu$m originates apparently from the 'ring' structure. It could be PAH and small hot grain emission enhanced by star formation in tidal shocks in the 'ring'. \begin{figure} \plottwo{altierib3.eps}{altierib4.eps} \caption{Abell 2218 15$\mu$m contour maps (a) and spectral energy distributions of 4 objects} \label{fig-2} \end{figure} Object \#323 is detected at 15 $\mu$m. This object was suspected to be an arclet candidate in the first deep multi-filter ground-based imaging (Pello et al. 1992), since it appeared as a fine unresolved extended structure, together with another 32 arclets in total. It is classified as one of the 235 arclet candidates in the first HST observation of Abell 2218 (Kneib et al. 1996). However, as noted by Kneib et al., it was also suspected not to be a strongly lensed object when comparison is made with shear orientation. This was again confirmed by Ebbels et al. (1997), object \#323 being very elongated but at an angle of 45 degrees from the shear direction. Identification of several absorption features in its spectrum reveal it to be a cluster member at a redshift of z=0.179, a spiral seen edge-on. To our surprise the brightest source by far from 7$\mu$m to 15$\mu$m is object \#395, an apparently insignificant z=0.1032 (Le borgne 1992) foreground Sb galaxy, showing an ultraviolet excess and a strong H$\alpha$ emission line, both indicating strong star formation of massive stars. Most of this mid-IR emission could then originate from the reprocessing by the small hot dust heated by a strong UV radiation field. The second strongest emitter at long wavelength (from 10$\mu$m to 15$\mu$m), but much less at shorter wavelengths is object \#317, a lensed galaxy at z=0.474 (Ebbels et al. 1997), a rather red object from the visible to the mid-IR. First indications show that this type of lensed object is common on our larger field imaging towards lensing clusters (Metcalfe et al. 1998 in preparation) The spectral energy distributions from the visible (B,R,I) to the mid-IR of the 4 objects discussed above are shown in fig. 2b | 98 | 3 | astro-ph9803155_arXiv.txt |
9803 | astro-ph9803188_arXiv.txt | In a universe reionized in patches, the Doppler effect from Thomson scattering off free electrons generates secondary cosmic microwave background (CMB) anisotropies. For a simple model with small patches and late reionization, we analytically calculate the anisotropy power spectrum. Patchy reionization can, in principle, be the main source of anisotropies on arcminute scales. On larger angular scales, its contribution to the CMB power spectrum is a small fraction of the primary signal and is only barely detectable in the power spectrum with even an ideal, i.e.~cosmic variance limited, experiment and an extreme model of reionization. Consequently patchy reionization is unlikely to affect cosmological parameter estimation from the acoustic peaks in the CMB. Its detection on small angles would help determine the ionization history of the universe in particular the typical size of the ionized region and the duration of the reionization process. | It is widely believed that the cosmic microwave background (CMB) will become the premier laboratory for the study of the early universe and classical cosmology. This belief relies on the high precision of the upcoming MAP\footnote{{\tt http://map.gsfc.nasa.gov}} and Planck Surveyor\footnote{{\tt http://astro.estec.eas.nl/SA-general/Projects/Planck}} satellite missions and the high accuracy of theoretical predictions of CMB anisotropies given a definite model for structure formation (\cite{Hu95}\ 1995). To realize the potential of the CMB, aspects of structure formation affecting anisotropies at only the percent level in power must be taken into account. A great uncertainty in models for structure formation is the extent and nature of reionization. Fortunately, this uncertainty is largely not reflected in the CMB anisotropies due to the low optical depth to Thomson scattering at the low redshifts in question. Reionization is known to be essentially complete by $z \sim 5$ from the absence of the Gunn-Peterson effect in quasar absorption spectra (\cite{Gun65}\ 1965). Significant reionization before $z\sim 50$ will be ruled out once the tentative detections of the CMB acoustic peaks at present are confirmed (\cite{Sco95} 1995). It should be possible to deduce the reionization redshift $z_i$ through CMB polarization measurements (\cite{Zal97}\ 1997). Nevertheless, the duration of time spent in a partially ionized state will remain uncertain. Moreover as emphasized by \cite{Kai84} (1984) secondary anisotropies generated by the Doppler effect in linear perturbation theory are suppressed on small scales for geometric reasons (gravitational instability generates potential flows, leading to cancellations between positive and negative Doppler shifts). Higher order effects which are not generally included in the theoretical modeling of CMB anisotropies are likely to be the main source of secondary anisotropies from reionization below the degree scale. Such effects rely on modulating the Doppler effect with spatial variations in the optical depth. Incarnations of this general mechanism include the Vishniac effect from linear density variations (\cite{Vis87}\ 1987), the kinetic Sunyaev-Zel'dovich effect from clusters (\cite{SZ}\ 1970), and the effect considered here: the spatial variation of the ionization fraction. Reionization commences when the first baryonic objects form stars or quasars that convert a part of the nuclear or gravitational energy into UV photons. Each such source then blows out an ionization sphere around it. Before these regions overlap is a period when the universe is ionized in patches. The extent of this period and the time evolution of the size and number density of these patches depend on the nature of the ionizing engines in the first baryonic objects. Theories of reionization do not give robust constraints ({\it cf.} \cite{Teg94}\ 1994; \cite{Ree96}\ 1996; \cite{Agh96}\ 1996; \cite{Loe97}\ 1997; \cite{Hai97}\ 1997; \cite{Sil98}\ 1998; \cite{Hai98}\ 1998). We therefore take a phenomenological approach to studying the effects of patchy reionization on the CMB. We introduce a simple but illustrative three parameter model for the reionization process based on the redshift of its onset $z_i$, the duration before completion $\delta z$, and the typical comoving size of the patches $R$. It is then straightforward to calculate the CMB anisotropies generated by the patchiness of the ionization degree of the intergalactic medium. We find that only the most extreme models of reionization can produce degree scale anisotropies that are observable in the power spectrum given the cosmic variance limitations. A large signal on degree scales requires early ionization, $z_i\gtrsim 30$, long duration, $\delta z\sim z_i$, and ionization in very large patches, $R\gtrsim 30$Mpc. Thus the patchiness of reionization is unlikely to affect cosmological parameter estimation from the acoustic peaks in the CMB (\cite{Jun96}\ 1996; \cite{Zal97}\ 1997; \cite{Bon97}\ 1997). On the other hand, the patchy reionization signal on the sub-arcminute scale can, in principle, surpass both the primary and the secondary Vishniac signals. These may be detectable by the Planck Surveyor and upcoming radio interferometry measurements (\cite{Par97}\ 1997) if point sources can be removed at the $\Delta T/T \sim 10^{-6}$ level. An explicit expression for the CMB anisotropies power spectrum generated in a universe reionized in patches is given \S \ref{sec:explicit}. Simple order of magnitude estimate of the anisotropy from patches is given in \S \ref{sec:order}. In \S \ref{sec:power} we give a rigorous definition of our three-parameter reionization model, and calculate the patchy part of the power spectrum. We discuss illustrative examples in \S \ref{sec:discussion}. | The signal from patchy reionization in our model depends on four quantities: the rms peculiar velocity $\left< v^2 \right>^{1/2}$ today, the redshift of reionization $z_i$, its duration $\delta z$ and the characteristic comoving size of the patches $R$. The structure formation model specifies the power spectrum of fluctuations which in turn tells us the rms peculiar velocity. Let us now consider the patchy reionized signal in the context of a specific model for structure formation. For illustrative purposes, let us consider a cold dark matter model with $h=0.5$, $\Omega_b =0.1$, and a scale-invariant $n=1$ spectrum of initial fluctuations. Normalizing the spectrum to the COBE detection via the fitting formulae of \cite{Bun97} (1997) (their equations [17]-[20]) and employing the analytic fit to the transfer function of \cite{Eis98} (1998) (their equations [15]-[24]) we find an rms velocity of $\left<v^2\right>^{1/2}=3.9\times 10^{-3}$. With the present optical depth of $\tau_0 = 0.122 \Omega_b h = 0.0061$, we have a maximal anisotropy of \begin{equation} ({l^2C_l\over 2\pi })_{\rm max}=2.41\times 10^{-15}{R\over {\rm Mpc}}\delta z(1+z_i)^{3/2}, \end{equation} at \begin{equation} l_{\rm max}={16958\over R/{\rm Mpc}} [ 1 - (1+z_i)^{-1/2}]. \end{equation} The power spectrum of the model in principle also tells us the remaining parameters of the ionization: its redshift $z_i$, duration $\delta z$ and typical patch size $R$. Unfortunately, these quantities depend on details of the cooling and fragmenting of the first baryonic objects to form the ionizing engines. We therefore consider $5 \lesssim z_i \lesssim 50$ which spans the range of estimates in the literature (\cite{Teg94}\ 1994; \cite{Ree96}\ 1996; \cite{Agh96}\ 1996; \cite{Loe97}\ 1997; \cite{Hai97}\ 1997; \cite{Sil98}\ 1998; \cite{Hai98}\ 1998). Reionization, once it commences, is generally completed in a time short compared with the expansion time at that epoch $\delta z /(1 +z_i) < 1$ by the coalescence of patches that are small compared with the horizon at the time $R/\eta_i \ll 1$ at the time. Again the exact relations depend on the efficiency with which the first objects form and create ionizing radiation (see e.g. \cite{Teg94}\ 1994). \begin{figure}[htb] \psfig{figure=patchf1.eps,width=3.3in} \caption{CMB anisotropy power spectra in a CDM model with extreme patchiness. Shown here are the primary anisotropy and the patchy reionization anisotropy, eq.~(\protect\ref{eqn:Cl}) with $z_i=10$, $\delta z=3$, $R=20$Mpc. These signals are compared with the cosmic variance of the primary anisotropy and the noise of the MAP satellite (in logarithmic bins).} \label{fig:patch10} \end{figure} Let us consider an extreme example of $z_i=10$, $\delta z=3$, $R=20{\rm Mpc}$. Then the maximal power is $\approx 5.3\times 10^{-12}$ at $l\approx 590$, the primary signal at these scales is $\approx 3\times 10^{-10}$ -- the contribution of the patchy reionization is small in comparison (see Fig.~\ref{fig:patch10}). However, in light of the high precision measurements expected from the MAP and Planck satellites such a signal is not necessarily negligible. The ultimate limit of detectability through power spectrum measurements is provided by so-called cosmic variance. This arises since we can only measure $2\ell+1$ realizations of any given multipole such that power spectrum estimates will vary by \begin{equation} \delta C_\ell = \sqrt{2 \over 2\ell+1} C_\ell^{(\rm primary)}. \end{equation} Detection of a broad feature such as that from patchy reionization is assisted in that we may reduce the cosmic variance by averaging over many $\ell$'s. We show an example of this averaging in Fig.~\ref{fig:patch10} (lower left boxes). In this model, the patchy reionization signal can be detected at the several $\sigma$ level if cosmic variance were the main source of uncertainty. Of course a realistic experiment also has noise and systematic errors. We also show the noise error contributions expected from the MAP experiment in Fig.~\ref{fig:patch10}. An important additional source of uncertainty is provided by other unknown aspects of the model. Indeed it is hoped that the CMB power spectrum can be used to measure fundamental cosmological parameters to high precision. Excess variance from patchy reionization can in principle cause problems for cosmological parameter estimation from the CMB if not included in the model. It would remain undetected and produce parameter misestimates if its signal can be accurately mimicked by variations in the other parameters. Fortunately, the angular signature we find here -- $\ell^2$ white noise until some cut off due to the patch size -- does not resemble the signature of other cosmological parameters which alter the positions and amplitudes of the acoustic peaks (see \cite{Bon97}\ 1997; \cite{Zal97}\ 1997). Coupled with the small amplitude of the effect on the 10 arcminute to degree scale for even this extreme model, it is unlikely that patchy reionization will significantly affect parameter estimation through the CMB. We have called the ($z_i=10,\delta z=3,R=20$) model extreme, because of the size of patches; the reionization redshift and duration would be considered reasonable by a number of theories. For example the early quasar model of Haiman \& Loeb (1998) does predict $z_i\sim 10$ and $\delta z \sim 3$. However, their ``medium quasar'' emits only $\sim 10^{67}$ ionizing photons during its life time. These photons cannot ionize a bubble larger than $R\sim 1$Mpc comoving. Perhaps more interesting is the case where reionization takes place at a higher redshift with for example $z_i=30$, $\delta z=5$, $R=3{\rm Mpc}$. The reduction in the patch size causes the signature to move to smaller angles where the primary signal is negligible due to dissipational effects at recombination. The increase in the optical depth at this higher redshift is counterbalanced by the reduction in the rms fluctuation due to the number of patches along the line of site such that the amplitude of the signal increases only moderately. Here the maximal power is $\approx 6.2\times10^{-12}$ at $l\approx 4650$ (see Fig.~\ref{fig:patch30}). Patchy reionization effects exceed the Vishniac signal at these scales ($\approx 3\times 10^{-12}$) which is believed to be the leading other source of secondary anisotropies (\cite{Hu96}\ 1996). \begin{figure}[htb] \psfig{figure=patchf2.eps,width=3.3in} \caption{CMB anisotropy power spectra in a CDM model with early reionization. Shown here are the primary anisotropy suppressed by rescattering and the patchy reionization anisotropy, eq.~(\protect\ref{eqn:Cl}) with $z_i=30$, $\delta z=5$, $R=3$Mpc. These signals are compared with the cosmic variance of the primary anisotropy achievable by an ideal experiment in the absence of galactic and extragalactic foregrounds.} \label{fig:patch30} \end{figure} Although the morphology and amplitude of the patchy reionization and Vishniac signals are similar, the Vishniac effect is fully specified by the ionization redshift and the spectrum of initial fluctuations and hence may be removed once these are determined from parameter estimation at larger angular scales. Likewise, since the rms peculiar velocity $\left< v \right>$ and the ionization redshift $z_i$ will be specified by the large scale observations, the amplitude of the signal can be used to estimate the duration of reionization $\delta z$ and its angular location the typical comoving size of the bubbles $R$. In summary, the patchiness of reionization leaves a potentially observable imprint on the CMB power spectrum, but one that is unlikely to affect cosmological parameter estimation from the acoustic peaks in the CMB. We show how the signature scales with the gross properties of reionization -- its redshift, duration, and typical patch size. Observational detection of this signature would provide useful constraints on the presently highly uncertain reionization scenarios but will likely require experiments with angular resolution of an arcminute or better and foreground subtraction at better than the $\delta T/T \sim 10^{-6}$ level. | 98 | 3 | astro-ph9803188_arXiv.txt |
9803 | astro-ph9803231_arXiv.txt | We present radio continuum observations of the spiral-shaped ionized feature (Sgr A West) within the inner pc of the Galactic center at three epochs spanning 1986 to 1995. The VLA A-configuration was used at $\lambda$2cm (resolution of 0\dasec1$\times$0\dasec2). We detect proper motions of a number of features in the Northern and Eastern Arms of Sgr A West including the ionized gas associated with IRS 13 with V(RA)= 113$\pm$10, V(Dec)=150$\pm$15 km s$^{-1}$, IRS 2 with V(RA)= 122$\pm$11, V(Dec)=24$\pm$34 km s$^{-1}$ and the Norther Arm V(RA)= 126$\pm$30, V(Dec)=--207$\pm$58 km s$^{-1}$. We also report the detection of features having transverse velocities $>$1000 \kms\ including a head-tail radio structure, the ``Bullet'', $\approx$4$''$ northwest of Sgr A$^*$ with V(RA)= 722$\pm$156, V(Dec)=832$\pm$203 km s$^{-1}$, exceeding the escape velocity at the Galactic center. The proper motion measurements when combined with previous H92$\alpha$ radio recombination line data suggest an unambiguous direction of the flow of ionized gas orbiting the Galactic center. The measured velocity distribution suggests that the ionized gas in the Northern Arm is not bound to the Galactic center assuming a 2.5 million solar mass of dark matter residing at the Galactic center. This implies that the stellar and ionized gas systems are not dynamically coupled, thus, supporting a picture in which the gas features in the Northern Arm and its extensions are the result of an energetic phenomenon that has externally driven a cloud of gas cloud into the Galactic center. | The ionized gas known as Sagittarius A West (Sgr A West) appears as a three-Arm spiral-like structure (North, East, and West Arms) engulfing the inner pc of the Galaxy where Sgr A$^*$, the compact radio source at or near the dynamical center of the Galaxy lies (Ekers et al. 1983). These features are surrounded by neutral gas in the circumnuclear disk (CND) rotating with the velocity of about 100 \kms\ at the distance of 2 pc from the Galactic center (e.g. G\"usten et al. 1988). The kinematics of ionized gas surrounding Sgr A$^*$ show systematic velocities along various components of Sgr A West including Western Arc with a radial velocity structure which varies regularly between --100 and +100 \kms\ in the North-South direction (e.g. Serabyn et al. 1988; Herbst et al. 1993; Roberts \& Goss 1993). However, the velocity structure of the inner 10$''$ where there is a hole in the distribution of ionized gas, known as the ``mini-cavity'', becomes increasing more negative $\approx-$350 \kms\ approaching Sgr A$^*$ (Yusef-Zadeh, Morris \& Ekers 1989; Roberts, Yusef-Zadeh \& Goss 1996, hereafter RYG). Recent observations of stellar proper motions shows evidence of a 2.5$\times10^6$ \msol\ object lying close to the position of Sgr A$^*$ (Eckart and Genzel 1997). The stars orbit randomly around the Galactic center with increasing velocity dispersion around Sgr A$^*$, reflecting the gravitational potential of central mass. The ionized gas, on the other hand, is part of a coherent flow with systematic motion that is decoupled from the stellar orbits. Understanding the kinematics of the system of ionized gas is complicated by its incomplete view of the 3-dimensional geometry with respect to Sgr A$^*$ as well as by the interaction of orbiting gas with non-gravitational forces, such as the winds from the cluster of hot mass-losing stars near the Galactic center. To examine the gas kinematics and the true geometry of the ionized flow, in this {\it Letter} we present the results of proper motion measurements of ionized gas at $\lambda$2cm. | By combining the transverse and radial velocities, we are able to unambiguously determine the direction of ionized flow at the Galactic center. The predominant component of the motion in the plane of the sky is from east to west for most of the measured features with the exception of few places where the velocity of ionized gas is anomalously large. It appears that the flow of ionized gas in the Northern Arm (Box 12) originates in the northeast with red-shifted velocities in the orbital plane. The ionized gas then follows an orbital trajectory to the southwest as it crosses the plane of the sky and passes by Sgr A$^*$ before the ionized gas moves to the northwest. One idea that has been suggested to explain the origin of the mini-cavity and its unusual characteristics (e.g. Lutz et al. 1993; Melia et al. 1996) is the the collision of fast-moving Blobs with the orbiting ionized gas. The Blobs are hypothesized to be formed as a result of high velocity winds of IRS 16 cluster escaping from but focussed by the gravitational potential of Sgr A$^*$. This focusing mechanism allows the diffuse outflowing materials to collide with each other and form dense Blobs of ionized gas leaving the gravitational potential of Sgr A$^*$ (Wardle \& Yusef-Zadeh 1992). The anomalous high-velocity features seen in Boxes 1, 4, 8 and 9 are consistent with the outflow picture. In particular, the Bullet is clearly escaping from the gravitational potential of the Galactic center region even when the mass of the stellar cluster is included. In this model, however, it is not expected to see a tail produced behind the fast moving Blobs. The existence of a bow shock or X-ray emitting gas associated with the head of the Bullet would favor a model in which these fast-moving features are ejected by mass-losing stellar sources. Further observations of the Bullet are needed to understand its origin. The comparison between the measured total velocities and the upper limits to the escape velocities at projected distance (r) from the center of each box to Sgr A$^*$ (Table 1) suggests that the ionized flow is on an unbound orbit around the Galactic center. This is consistent with the model of RYG that the ionized gas in the Northern Arm is on a hyperbolic orbit. Like the compact and relatively isolated features with high transverse velocities discussed above (e.g. the Bullet and the Blobs), the rest of the orbiting ionized features are extended and follow a global velocity field which may also be unbound to the Galactic center. If we use the projected distance as cos$^{-1}$ (45$^\circ$) of the actual distance, almost all the measured velocities are greater than the escape velocities listed in Table 1. The mass within the inner 10-20$''$ is assumed to be dominated by a compact source centered on Sgr A$^*$ having a mass of 2.5$\times10^6$ \msol\ as measured recently from stellar proper motion measurements (Eckart \& Genzel 1997). The escape velocity estimates may not be applicable at large distances where the mass of the evolved stellar cluster becomes important. From the comparison of the three dimensional stellar and ionized gas motions, it appears that these two systems are not dynamically coupled in the inner 20$''$ of the Galactic center. The effect of a strong gravitational potential due to the large concentration of dark matter near Sgr A$^*$ is manifested as high velocity gradients of over 600 \kms\ pc$^{-1}$. However, the existence of ionized gas in an unbound orbit is inconsistent with the notion that the ionized gas in the Northern Arm is a tidally stretched infalling feature (e.g. Serabyn et al. 1988). Additionally, the present proper motion data do not support the interpretation that the Northern Arm is a segment of a one-armed spiral pattern induced as a result of an instability in the rotating disk (e.g. Lacy et al. 1991). A significant variation in the velocity distribution of ionized gas along Northern Arm is not consistent with a small variations expected from Keplerian motion, thus supporting that the Northern Arm is a material feature. We believe that the high velocity of ionized gas on an unbound orbit supports a scenario in which an energetic phenomenon outside the inner few parsecs of the Galactic center accelerated a cloud to pass by the Galactic center, which then collides with the CND and results in the loss of angular momentum of the material in the CND (Serabyn et al. 1988). There is evidence of disturbed neutral gas and shocked molecular gas based on OH 1720 MHz maser emission at the interface of the extension of the Northern and Eastern Arms of Sgr A West and a ``gap'' in the CND (Yusef-Zadeh et al. 1996; Jackson et al. 1993). In this picture, the Northern and Eastern Arms delineate the edges of the intruding cloud photoionized by the UV radiation field at the Galactic center (Jackson et al. 1993). The neutral gas in the ``gap'' of the CND is interpreted to be the site of collision with a cloudlet pushed into the Galactic center, possibly by the energetic explosion of Sgr A East. Future modeling of the three-dimensional motion of ionized gas should constrain the inclination of the orbital plane of the ionized gas with respect to the orbital plane of the CND and examine the possibility that the CND is origin of the Northern and Eastern Arms. | 98 | 3 | astro-ph9803231_arXiv.txt |
9803 | astro-ph9803007_arXiv.txt | Since their discovery in 1973, Gamma-Ray Bursts (GRBs) have remained for many years one of the most elusive mysteries in High Energy-Astrophysics. The main problem regarding the nature of GRBs has usually been the lack of knowledge of their distance scale. About 300 GRBs are detected annua\-lly by BATSE in the full sky, but only a few of them can be loca\-lized accurately to less than half a degree. For many years, follow-up observations by other satellites and ground-based telescopes were conducted, but no counterparts at other wavelengths were found. The breakthrough took place in 1997, thanks to the observation by {\it BeppoSAX} and {\it RossiXTE} of the fa\-ding X-ray emission that follows the more energetic gamma-ray photons once the GRB event has ended. This emission (the afterglow) extends at longer wavelengths, and the good accuracy in the position determination by {\it BeppoSAX} has led to the discovery of the first optical counterparts -for GRB 970228, GRB 970508, and GRB 971214-, greatly improving our understanding of these puzzling sources. Now it is widely accepted that most bursts originate at cosmological distances but the final solution of the GRB problem is still far away. \vspace {5pt} \\ Key~words: multiwavelength observations; gamma-ray bursts. | In 1967-73, the four {\it VELA} spacecraft (named after the spanish verb {\it velar}, to keep watch), that where originally designed for verifying whether the former Soviet Union abided by the Limited Nuclear Test Ban Treaty of 1963, observed 16 peculiarly strong events (Klebesadel, Olson and Strong 1973, Bonnell and Klebesadel 1996). On the basis of arrival time differences, it was determined that they were related neither to the Earth nor to the Sun, but they were of cosmic origin. Therefore they were named cosmic $\gamma$-ray Bursts (GRBs hereafter). GRBs appear as brief flashes of cosmic high energy photons, emitting the bulk of their energy above $\approx$ 0.1 MeV (Fig. 1). They are detected by instruments somewhat similar to those used by the particle physicists at their laboratories. The difference is that GRB detectors have to be placed onboard balloons, rockets or satellites. In spite of the abundance of new observations of GRBs, their energy source and emission mechanism remain highly speculative. The KONUS experiment on {\it Veneras 11} and {\it 12} gave the first indication that GRB sources were isotro\-pically distributed in the sky (Mazets et al. 1981, Atteia et al. 1987). Based on a much larger sample, this result was nicely confirmed by BATSE on board the {\it CGRO} (Meegan et al. 1992). About 300 GRBs occur annually in the full sky, but only few of them are localized accurately. The apparent isotropy was interpreted in terms of GRBs produced at cosmological distances, although the possibility of a small fraction of the sources lying nearby, within a galactic disc scale of few hundred pc, or in the halo of the Galaxy, could not be discarded. Another result was that the time profiles of the bursts are very diffe\-rent, with some GRBs lasting a few ms and others lasting for several minutes. In general, there was no evidence of periodicity in the time histories of GRBs. However there was indication of a bimodal distribu\-tion of burst durations, with $\sim$25\% of bursts ha\-ving durations around 0.2 s and $\sim$75\% with durations around 30 s. A extensive review of the observational character\-istics can be found in Fishman and Meegan (1995). A deficiency of weak events was also noticed, and all these observational data led many researchers to believe that GRBs are indeed at cosmological distances. In this case, the released energies could be as high as 10$^{53}$ erg and models could involve coalescence of neutron stars in double systems, neutron star-black hole systems, accretion induced collapse in white dwarfs or $^{\prime\prime}$failed$^{\prime\prime}$ Type I supernovae. It was also proposed the possibility of having a mixture of two popu\-lations: a cosmological plus a galactic one, but the latter, probably formed by accreting neutron stars, will account for only a very small fraction of the total population. See Nemiroff (1984) for a review of the di\-fferent theoretical models. It is well known that an important clue for resol\-ving the GRB puzzle is the detection of transient emission -at longer wavelengths- associated with the bursts. Here I review all the efforts in the search for GRB counterparts throughout the electromagnetic spectrum. I will first review the searches prior to 1997, and afterwards I will discuss the important discove\-ries achieved last year. Previous reviews can be seen in Schaefer (1994), Hartmann (1995), Vrba (1996), Greiner (1996a) and Hurley (1998). \begin{figure}[htp] \epsfxsize=70mm \epsfysize=60mm \centerline{\epsfbox{fig1.eps}} \caption{One of the GRBs detected by {\it BATSE}, lasting for about 150-s. Further details in Connaughton et al. (1997).} \end{figure} | The existence of an X-ray afterglow in {\it all} bursts seems to be confirmed. The first optical/IR counterparts have been found in 1997. GRB 970508 was also detected in radio and mm, but prompt searches at other wavelengths failed to detect GRB 970111, GRB 970402 and GRB 970828. {\it BeppoSAX} and {\it RossiXTE} have opened a new window in the GRB field and it is widely accepted now that most GRBs, if not all, lie at cosmological distances. It is expected that {\it BeppoSAX}, {\it RossiXTE} and {\it CGRO} will facilitate the discoveries of other counterparts and, together with the new high-energy observatories ({\it AXAF}, {\it SPECTRUM X/$\Gamma$}, {\it XMM}, {\it INTEGRAL}, {\it HETE 2}) and other satellites of the future {\it 4th Interplanetary Network}, will definitively solve the long-standing Gamma-Ray Burst mystery. | 98 | 3 | astro-ph9803007_arXiv.txt |
9803 | astro-ph9803282_arXiv.txt | Microlensing is increasingly gaining recognition as a powerful method for the detection and characterization of extra-solar planetary systems. Naively, one might expect that the probability of detecting the influence of more than one planet on any single microlensing light curve would be small. Recently, however, Griest \& Safizadeh (1998) have shown that, for a subset of events, those with minimum impact parameter $u_{min} \lsim 0.1$ (high magnification events), the detection probability is nearly 100\% for Jovian mass planets with projected separations in the range 0.6--1.6 of the primary Einstein ring radius $R_E$, and remains substantial outside this zone. In this Letter, we point out that this result implies that, regardless of orientation, {\it all} Jovian mass planets with separations near 0.6--1.6$R_E$ dramatically affect the central region of the magnification pattern, and thus have a significant probability of being detected (or ruled out) in high magnification events. The probability, averaged over all orbital phases and inclination angles, of two planets having projected separations within $0.6$--$1.6R_E$ is substantial: 1-15\% for two planets with the intrinsic orbital separations of Jupiter and Saturn orbiting around 0.3--1.0$M_\odot$ parent stars. We illustrate by example the complicated magnification patterns and light curves that can result when two planets are present, and discuss possible implications of our result on detection efficiencies and the ability to discriminate between multiple and single planets in high magnification events. | A planetary microlensing event occurs whenever the presence of a planet creates a perturbation to the standard microlensing event light curve. These perturbations typically have magnitudes of $\lsim 20\%$ and durations of a few days or less. First suggested by Mao \& Paczy\'nski (1991) as a method to detect extra-solar planetary systems, the possibility was explored further by Gould \& Loeb (1992), who found that roughly 15\% of microlensing light curves should show evidence of planetary deviations if all primary lenses have Jupiter-mass planets with orbital separations comparable to that of Jupiter. Although these probabilities are relatively high, the use of microlensing to discover planets was largely ignored since in order to detect the primary events the microlensing survey teams must monitor millions of stars in very crowded fields, resulting in temporal sampling that is too low ($\sim 1\, {\rm day}$) and photometric errors that are too high ($\gsim 5\%$) to detect most secondary planetary deviations (Alcock et al.\ 1997a). Recently, the situation has changed dramatically as the real-time reduction of the survey teams has enabled them to issue electronic ``alerts,'' notification of on-going events detected before the peak magnification (Udalski et al.\ 1994, Pratt et al.\ 1996), allowing other collaborations to perform special purpose observations of the alerted events. These additional observations include denser photometric sampling by the PLANET and GMAN collaborations (Albrow et al.\ 1996, 1997, 1998 and Pratt et al.\ 1996, Alcock et al. 1997b) as well as spectroscopic follow-up of particular events (Lennon et al.\ 1997). Over 60 events are currently alerted per year towards the Galactic Bulge. Since only a handful of these are on-going at any given time, monitoring teams can sample events very densely and with high photometric accuracy, enabling the detection of many second order effects, including --in principle-- planetary anomalies. No clear planetary detections have yet been made in this way, but preliminary estimates of detection efficiencies show that PLANET, over the next two observing seasons, should be sensitive to planetary anomalies caused by Jovian planets orbiting a few AU from their parent star (Albrow et al.\ 1998). Thus, if these kinds of planets are common, they should be detected soon. If not, microlensing will be able to place interesting upper limits on the frequency of such systems. These observational developments have been accompanied by an explosion of theoretical work, including further studies of detection probabilities and observing strategies incorporating a variety of new effects (Bolatto \& Falco 1994, Bennett \& Rhie 1996, Peale 1997, Sackett 1997, DiStefano \& Scalzo 1998a,b), demonstration of planetary microlensing light curves (Wambsganss 1997), explorations of the degeneracies in the fits of planetary events (Gaudi \& Gould 1997, Gaudi 1998), and a study of the relation between binary and planetary lenses (Dominik 1998). It would thus seem that the theoretical understanding of the detection and characterization of planetary systems using microlensing should be well in hand. Surprisingly, however, the field still has surprises to offer. Recently, Griest \& Safizadeh (1998, hereafter GS) came to a rather startling | In this {\it Letter\/}, we have demonstrated that: (1) the probability of two planets having projected separations that fall in the ``standard lensing zone'' ($0.6 < b < 1.6$) is quite high, $\sim 1-15\%$ for planets with true separations corresponding to Jupiter and Saturn orbiting stars of typical mass; (2) the influence of multiple planets in and somewhat beyond the standard lensing zone can be profound for high magnification events ($u_{min}<0.1$,) however (3) for some geometries, the magnification pattern and resulting light curves from multiple planets are qualitatively degenerate with those from single-planet lensing, and (4) for high magnification events, finite source effects are likely to suppress more substantially the amplitude of multiple planet deviations than single planet deviations. Given these results, it would appear that the effects of multiple planets on the detection and characterization of planetary systems warrant future study. All previous theoretical studies have calculated microlensing planet detection sensitivities either by ignoring multiple planets or by treating each planet independently. For high impact parameter events (low magnification), this is probably a fair assumption, but as the magnification maps in Fig.~2 illustrate detection probabilities will need to be revised for small impact parameters (large magnification). The sense of revision will likely depend on finite source effects. It is also likely that for some geometries serious degeneracies exist between light curves arising from multiple and single planet high magnification events; these degeneracies are above and beyond those present in the single planet case discussed by Griest \& Safizadeh (1998). This possible degeneracy is especially pertinent in light of the fact that the conditional probability of having two planets in the lensing zone is substantial. Thus, the interpretation of any given high magnification event may be difficult: the degeneracies should be characterized and their severity determined in order to have a clear understanding of the kinds of systems whose parameters can be unambiguously determined from the deviations. Finally, the calculation of planet detection efficiencies for high magnification events should consider multiple planets in order to be able to reliably convert the observed frequency of planetary deviations into a true frequency of planetary systems. | 98 | 3 | astro-ph9803282_arXiv.txt |
9803 | astro-ph9803018_arXiv.txt | In observations with the Rossi X-ray Timing Explorer we have discovered quasi-periodic oscillations (QPOs) near 1 kHz from 4U~1705-44, a low--mass X-ray binary with a neutron star classified as an atoll source. In six separate observations, we detect one QPO with a frequency ranging between 770 and 870 Hz and a 4\% rms fraction in the full detector energy band. There is evidence for a second QPO at 1073 Hz in one interval. The separation in frequency of the two QPOs is $298\pm11$ Hz. The QPOs are present only in observations where the mass accretion rate is inferred to be at an intermediate level, based on the atoll source phenomenology. At the highest accretion rates, the QPOs are not detected with upper limits to the rms fraction of about 2\%. At the lowest accretion rates the upper limits are about 4\%. The QPO frequency increases with inferred mass accretion rate. This is expected in models where the QPO frequency is generated by motion at an inner edge of the accretion disk. An increased mass accretion rate causes the disk edge to move in, increasing the orbital frequency. Five Type--I X-ray bursts are observed with no detectable oscillations. | Quasi-periodic oscillations (QPOs) with frequencies near 1 kilohertz were discovered from X-ray binaries with the Rossi X-Ray Timing Explorer (RXTE) almost immediately after its launch (for reviews and references see van der Klis 1997, Kaaret \& Ford 1997a, Ford 1997). These QPOs are likely produced very near the accreting neutron star where the dynamical time scale is near $10^{-3}$ seconds. Observations of such fast oscillations provides a new opportunity to measure fundamental properties of neutron stars and the effects of strong field gravity in low mass X-ray binaries. Kilohertz QPOs have been discovered in both major subclasses of LMXBs: the Z-sources and the atoll sources (see Hasinger \& van der Klis 1989). Certain trends are already apparent. For example, the kilohertz QPOs in Z-sources are much weaker than those in the lower luminosity atoll sources. In the sources with luminosities intermediate between the majority of Z and atoll sources, the persistently bright galactic bulge sources, QPOs are apparently absent or very weak (Wijnands, van der Klis, \& van Paradijs 1997a). To understand the physical mechanisms at work, further observations of a large sample of sources are needed with the unique capabilities of RXTE, likely to be unequaled in the next decade. Here we report the discovery of fast QPOs from the low mass X-ray binary 4U~1705-44 using RXTE. 4U~1705-44 has been classified as an atoll source (Hasinger \& van der Klis 1989). Its different source states are supposed to reflect changes in mass accretion rate. We observe a link between the source state and the presence and frequencies of fast QPOs. In Section~2 we summarize the observations and analysis techniques. Section~3 presents the detections of fast QPOs and X-ray burst properties. In Section~4 we identify the atoll source states during our observations. Section~5 is a discussion of the relation of the source states to the QPO properties and physical implications. | The 760--870 Hz QPO in 4U~1705-44 is probably a single feature and is likely the lower frequency of two QPO peaks. The second QPO, which we identify at 1074 Hz, then is the higher frequency peak. More observations in the lower banana state are needed to confirm the presence of this second QPO. Such phenomenology is very similar to other atoll sources. A single strong QPO near 800 Hz is present at times in several sources, e.g. 4U~1608-52 (Berger et al. 1996; Mendez et al. 1998), 4U~1820-30 (Smale, Zhang \& White 1997), 4U~1636-53 (Zhang et al. 1996; Wijnands et al. 1997b). In all of these sources a second QPO is sometimes detectable (see above references). The frequency separation of the two peaks is similar to that in 4U~1705-44. Such double QPOs, and also the QPOs observed in X-ray bursts (e.g. Strohmayer, Zhang \& Swank 1997), suggest a beat frequency mechanism (Alpar \& Shaham 1985; Miller, Lamb \& Psaltis 1998). If such an interpretations holds, then the frequency separation of the QPOs in 4U~1705-44 implies that the spin period of the neutron star in this system is $3.35\pm0.12$ msec. The appearance and frequency of the fast QPOs in 4U~1705-44 are correlated with the states of the source. The states in turn are thought to be related to the mass accretion rate, with the accretion rate increasing from the island to the banana. At the lower end of the banana branch a QPO appears at about 780 Hz. Further along the branch the frequency increases to 835 Hz and then 865 Hz. This behavior supports models where the QPO modulation is generated at an inner disk edge which shrinks as the mass accretion rate increases (e.g. Miller, Lamb \& Psaltis 1998). We detect the second QPO at 1073 Hz only in the the lowest part of the banana where the frequency of the stronger QPO is at its lowest. Farthest along the banana branch the QPO becomes undetectable with rms fraction upper limits of about 2\%. At the highest accretion rates, the reduced QPO amplitude is perhaps due to spreading of the accretion stream, alternatively a puffed--up inner disk may obscure the central source (Smale, Zhang \& White 1997). Similar behavior is observed in other X-ray binaries. Atoll source state identifications have been made simultaneous with fast QPO detection in 4U~1636-53 (Zhang et al. 1996; Wijnands et al. 1997b), 4U~1735-44 (Wijnands et al. 1997c), 4U~1820-30 (Smale, Zhang \& White 1997) and KS~1731-260 (Wijnands \& van der Klis 1997d). These sources were observed in the banana branch. Consistent with the present data, the QPO in all cases is detected only in the lower part of the banana and disappears as the source moves up along the banana branch. In 4U~0614+091 a similar effect appears as the QPO amplitude decreases as the count rate increases (Ford 1997). The situation is less clear at low mass accretion rates, i.e. in the island states. Our upper limits of about 4\% rms fraction are not very constraining. QPOs have been detected in island states in both 4U~0614+091 (Mendez et al. 1997) and 4U~1608-52 (Yu et al. 1997). At the lowest count rates in the island state of 4U~0614+091 the QPO is not detected (Mendez et al. 1997). This fits the general trend we observe here. The QPO frequency shows no correlation with count rate or energy flux in the 2--10 keV band, though we note that the range of count rates is small (1205 to 1236 c/s). An effect similar to that seen in 4U~0614+091 may be at work, where there is no unique correlation between rate and frequency in different observations separated by several months (Ford et al. 1997a). The present data indicate that the QPO frequency is better correlated with the atoll state of the source. Changes in the source state are reflected as changes in the x-ray colors, which in turn are simply changes in the energy spectrum of the source. Therefore the state--frequency correlation also manifests as a spectrum--frequency correlation. From other sources we know that the changes in QPO frequency are linked to changes in the energy spectrum. The frequencies of the QPOs in 4U~0614+091 are well correlated with the flux of a blackbody component in the energy spectrum (Ford et al. 1997b), and in that source and also 4U~1608-52 the QPO frequencies are correlated with the spectral index of the power law component in the energy spectrum (Kaaret et al. 1998). We thank Mariano Mendez for assistance in implementing the frequency shifting technique. We thank Rudy Wijnands, Rob Fender and the referee for helpful discussions and comments. ECF acknowledges support by the Netherlands Foundation for Research in Astronomy with financial aid from the Netherlands Organization for Scientific Research (NWO) under contract numbers 782-376-011 and 781-76-017. | 98 | 3 | astro-ph9803018_arXiv.txt |
9803 | astro-ph9803312_arXiv.txt | In this talk, I have discussed some issues of recent interest and activity in the field of neutrino astrophysics and cosmology. The topics are: (1)~The origin of high peculiar velocities of pulsars; (2)~Energization of the supernova shock wave; (3)~Ultra-high energy neutrino astronomy; (4)~Possible implications of the recent measurements of low deuterium abundance. | It was known, since the birth of modern astrophysics in the early part of the 20th century, that neutrinos play an important role in various processes that occur within a stellar core and which are responsible for energy generation in a star. Gradually, the importance of neutrinos were understood in stars outside the main sequence. And, since the discovery of the microwave blackbody radiation, it was taken for granted that there is a similar cosmic background of neutrinos, although experimentally this background has not been detected so far. Various constraints from neutrino properties have been deduced from this belief, some of which are much better than the corresponding constraints from earth-based experiments. For example, one can cite the mass bound on stable neutrinos which are derived from the energy density of the universe as a whole. This sets an upper bound of order of a few tens of eV, whereas the direct measurement of the mass of $\nu_\tau$ sets upper bounds in the range of a few tens of MeVs. If the neutrinos are unstable, then also there exists quite severe bounds on their lifetimes. Unfortunately, in this talk I cannot review all of these aspects. Rather, I will have to assume that the audience is familiar with these concepts. The reason is that, fortunately, there has been a lot of progress in the field of neutrino astrophysics in the last year and a half, and quite a few of them are remarkable. I have to concentrate on these recent developments. I cannot guarantee that I will cover even all of the interesting recent developments. Let me say I will cover what I know, with the restriction that I will leave out topics such as solar neutrinos and atmospheric neutrinos, which are covered by other speakers in this conference. | 98 | 3 | astro-ph9803312_arXiv.txt |
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9803 | astro-ph9803124_arXiv.txt | Echelle spectra have been obtained of the \caII\ H and K lines for a sample of metal-poor subdwarf stars as well as for a number of nearby Population~I dwarfs selected from among those included in the Mount Wilson HK survey. The main conclusion of this paper is that \caII\ H- and K-line emission does occur among subdwarfs. It is particularly notable among those subdwarfs with colours of $B-V \geq 0.75$; all such stars observed exhibit chromospheric emission, although emission is observed among some subdwarfs bluer than this colour. The \caII\ K emission profile in most subdwarfs exhibits an asymmetry of $V/R >1$, similar to that seen in the integrated light of the solar disk. Two quantitative indicators of the contrast between the peaks in the emission profile and the neighbouring photospheric line profile are introduced. Measurements of these indicators show that the level of \caII\ emission among the subdwarfs is similar to that among low-activity Population~I dwarfs. | Chromospheric activity among main-sequence stars is well known to be a decreasing function of age. The fluxes in the chromospheric emission features formed in the cores of the \caII\ H and K lines show a decrease with age among stars younger than 4.5 Gyr (Skumanich 1972). Within most main-sequence stars chromospheric heating is thought to be governed by the activity of an interior magnetic dynamo, the strength of which depends on the stellar rotation rate. As main-sequence stars age they spin down, with a consequent decrease in dynamo activity and chromospheric heating (Hartmann \& Noyes 1987). The dynamo model has met with considerable success in describing the evolution of chromospheric activity among main-sequence stars. However, much of the data published to date on chromospheric activity among dwarf stars of spectral types F-G-K pertains to stars younger than $\sim$ 5 Gyr. Such is true, for example, of the studies by Skumanich (1972) and Simon, Herbig, \& Boesgaard (1985). By comparison, there is much less spectroscopy available pertaining to chromospheric activity among older dwarfs, particularly halo subdwarfs. Consequently, little is known about the evolution of chromospheric activity among the oldest main-sequence stars in the Galaxy. To study the low levels of activity expected among such stars, we have carried out a spectroscopic study of the \caII\ H and K lines for a sample of low-metallicity subdwarfs. Metal-poor subdwarfs which are members of the Galactic halo are amongst the oldest stars in the Galaxy. They are useful for studying chromospheric activity at great ages. The observations of \caII\ H and K line emission among subdwarfs reported in this paper complement the HST GHRS observations of \mgII\ lines among metal-poor solar-like stars by Peterson \& Schrijver (1997). | The main conclusion of this paper is that Population II subdwarfs do have chromospheres, i.e., their outer atmospheres are heated by some form of non-radiative process. In terms of the contrast of the peaks in the K-line emission profile relative to the neigbouring absorption profile, and the preponderance of an asymmetry of $V/R > 1$, the subdwarfs appear similar to low-activity dwarfs sampled from the Mount Wilson survey. These observations are consistent with data on the \mgII\ lines of metal-poor solar-type stars reported by Peterson \& Schrijver (1997). The origin of an asymmetry of $V/R > 1$ among both dwarfs and subdwarfs is not well understood. As Linsky (1980) has pointed out such an asymmetry can be produced by downward motions in the region of formation of the central K$_3$ absorption feature, or by upwards motions in the K$_2$ formation region if there are no systematic motions in the region of K$_3$ formation (see Ayres \& Linsky 1975 for the definitions of this terminology). Presumably the reason for the $V/R >1$ asymmetry among the subdwarfs is the same as for the Population~I dwarfs, including the Sun. However, the origin of this asymmetry even for the best studied case of the Sun is still not completely clear. This subject has been reviewed by Linsky (1980), Cram (1983), and Rutten \& Uitenbroek (1991). The concensus seems to be that the asymmetry is associated with upwards propagating gravity or acoustic waves (see, e.g., Carlsson \& Stein 1992, 1997). In addition to the uncertainty in the origin of the $V/R$ asymmetry, the \caII\ K-line emission observations do not permit a determination of whether the chromospheric activity of the subdwarfs is being modulated by a magnetic dynamo, as for the Sun, or whether it is produced by a process which is active even in the absence of such a dynamo. Statistical analyses of the chromospheric emission line fluxes of late-type Population~I stars have been used to argue that the emission can be divided into two components (Schrijver 1987a,b, 1995): a ``magnetic dynamo'' component which correlates with stellar rotation rate and decreases with age, and a rotation-independent ``basal'' component. The magnetic dynamo component is important for stars like the Sun, and is particularly strong in late-type dwarfs younger than 1 Gyr. The basal component is thought to reflect the contribution of acoustic heating to the maintenance of a chromosphere (Schrijver 1987b; Cuntz, Rammacher, \& Ulmschneider 1994). Whether the chromospheres of subdwarfs are ``basal chromospheres'' or are maintained by a magnetic dynamo is not clear from the \caII\ K line observations. The finding by Baliunas et~al. (1995) that the subdwarf HD~103095 exhibits an activity cycle of period 7.3 years in the \caII\ H and K emission lines, suggests an analogy with the solar cycle, implying that the subdwarfs have chromospheres whose long-term behaviour is dominanted by a magnetic dynamo. The similarity in \caII\ K-line emission strength among the subdwarfs and the lowest-activity stars sampled from the Mount Wilson survey seems consistent with the spectroscopy of main-sequence stars in the old open clusters NGC 188 and M67 by Barry, Cromwell, \& Hegge (1984) which suggests that the rate of decline of the fluxes in the Ca II H and K line cores flattens out after $\sim$ 3 Gyr. These observations, and those noted in the previous paragraph, would be consistent with a scenario in which the magnetic dynamo ceases to decline significantly in activity once main-sequence stars reach an age of around 5 Gyr. There could be two reasons for this. Either the dynamo reaches a true minimum level, as in the turbulent magnetic field scenario of Durney, De Young, \& Roxburgh (1993), or else the rate of spin down of dwarfs and subdwarfs older than 5 Gyr becomes very slow. The alternative to these dynamo scenario is that chromospheres among stars older than 5 Gyr are largely maintained by an age-invariant basal heating mechanism, and that activity produced by an age-dependent magnetic dynamo has dropped to a low level by comparison. In sufficiently old main-sequence subdwarf stars the dynamo might be relatively inactive and acoustic heating may become the dominant mechanism for maintaining a ``basal chromosphere.'' This is the scenario proposed by Peterson \& Schrijver (1997) on the basis of similarities observed between the \mgII\ emission line profiles of metal-poor solar-type stars and solar quiet regions. In this respect it is perhaps relevant to note that the most extreme $V/R >1$ \caII\ K line asymmetry in the integrated spectrum of the solar disk occurs around the times of solar minimum (White \& Livingston 1981). The type of observation that may differentiate between these scenarios is a search for soft x-ray emission among subdwarfs. If subdwarfs are found to have solar-like coronae, with soft x-ray properties consistent with temperatures of 1-2$\times 10^6$ K, then the case for dynamo-maintained chromospheres and coronae would be strengthened. We thank the referee for a number of useful comments and for drawing our attention to a number of relevant references in the literature. \onecolumn | 98 | 3 | astro-ph9803124_arXiv.txt |
9803 | astro-ph9803254_arXiv.txt | We have imaged a sample of 45 face-on spiral galaxies in the K-band, to determine the morphology of the old stellar population, which dominates the mass in the disk. The K-band images of the spiral galaxies have been used to calculate different characteristics of the underlying density perturbation such as arm strengths, profiles and cross-sections, and spiral pitch angles. Contrary to expectations, no correlation was found between arm pitch angle and Hubble type, and combined with previous results this leads us to conclude that the morphology of the old stellar population bears little resemblance to the optical morphology used to classify galaxies. The arm properties of our galaxies seem inconsistent with predictions from the simplest density wave theories, and some observations, such as variations in pitch angle within galaxies, seem hard to reconcile even with more complex modal theories. Bars have no detectable effect on arm strengths for the present sample. We have also obtained B-band images of three of the galaxies. For these galaxies we have measured arm cross-sections and strengths, to investigate the effects of disk density perturbations on star formation in spiral disks. We find that B-band arms lead K-band arms and are narrower than K-band arms, apparently supporting predictions made by the large scale shock scenario, although the effects of dust on B-band images may contribute towards these results. | Ever since the pioneering work of Lin \& Shu (1964, 1966) many theoretical models have been proposed to explain the existence of spiral structure in disk galaxies. Many are based upon the Lin-Shu Hypothesis (e.g. Lin, Yuan \& Shu 1969; Roberts 1969; Roberts, Roberts \& Shu 1975) which describes spiral patterns as density waves, which cause compression of the gas component as it flows through the arms leading to subsequent star formation. Modal theories (Bertin et al. 1989a, 1989b; Bertin \& Lin 1996) are more complex density wave models, which describe disks of galaxies as resonant cavities within which density waves of different modes can co-exist and interfere to produce a range of observed phenomena. Tidal models (Toomre \& Toomre 1972; Kormendy \& Norman 1979) set up spiral structure through transient density waves, which are caused by the tidal field of a nearby neighbour. Bars can also potentially drive spiral structure (Sanders \& Huntley 1976). In this case the formation of arms is driven by the effect of the bar potential on the dissipative interstellar medium. Finally Stochastic Self-propagating Star Formation (SSPSF) can describe flocculent structure. Here, short irregular arms are produced due to shearing of a part of the disk that has recently formed stars (Gerola \& Seiden 1978). SSPSF alone cannot describe Grand-Design structure, but some models (e.g. Sleath \& Alexander 1995, 1996) use SSPSF with a weak imposed density wave to describe such global spiral modes. Many of the recent advances in this area have arisen from a study of the atomic and molecular gas component in spiral galaxies, through studies of HI and CO line emission. This has permitted the detailed mapping of both the distribution and velocity field of the gas in the spiral arms of nearby galaxies, e.g. M81 (Visser 1980) M51 (Tilanus \& Allen 1991, 1993; Rand 1993; Nakai et al. 1994), M100 (Knapen 1993; Rand 1995), NGC~3627 (Reuter et al. 1996) and NGC~6946 (Regan \& Vogel 1995). These studies find streaming velocities of gas through the spiral arms of order a few 10s of km~s$^{-1}$, and also find offsets between the peaks of the gas density and the old stellar population in the arms and the star formation as revealed by H$\alpha$ emission. All of these findings are in general agreement with the predictions of density wave theories, with the gas being compressed and shocked as it flows into the spiral arm, and subsequently forming stars downstream from the density peak of the spiral arm. However, these studies are only possible in the strongest arms of the nearest spiral galaxies, and hence yield little information on the global importance of density waves in the general population of disk galaxies. Here we undertake the complementary approach of calculating the strength of density waves as mapped out by the old stellar population of a large sample of galaxies, using near-IR K-band ($2.2\mu m$) imaging. In this waveband the old stellar population is observed (Rix \& Rieke 1993), making this the best observational tracer of the stellar mass distribution. Also, K--band measurements are less affected by extinction than measurements in the optical. In a previous paper (Seigar \& James 1998 - hereafter Paper I) we described the properties of the bulges, disks and bars of a sample of 45 galaxies. In this paper we describe how spiral arm structures can be extracted for these galaxies and determine quantitative parameters for them. In section 2 we describe the observations; section 3 describes the data reduction and extraction of all of the spiral arm parameters used in the analysis; section 4 then draws on these parameters in discussing, in turn, each of the theoretical models of spiral structure; in section 5 we describe tests of models of star formation in spiral galaxy disks and in section 6 we summarise our findings. | We now summarise the implications of these observations for the main models of the formation and stability of spiral structure. No one theory fits all of the observations, but a combination can explain most or all of the features we have observed. This section summarises these conclusions. Simple density wave theories (e.g. Lin \& Shu 1964, 1966) are undermined in three ways. Firstly, we do not get the correlation between the fraction of light in the disk and pitch angle which is predicted by Lin \& Shu (1964) and Roberts et al. (1975). Secondly, Lin \& Shu (1964) predict that pitch angle should be the same for arms at the same radius within a given galaxy. We find that this is not the case. Finally, we have found that the FWHM of arm cross-sections are somewhat narrower than predicted by simple density wave theories. More complicated density wave theories, e.g. Modal theory (Bertin et al. 1989a, 1989b; Bertin \& Lin 1996), are supported by modulation effects that we have found in arm strength as a function of radius, and which appear to be fairly common. Bar driven models of spiral structure (Sellwood \& Sparke 1988; Sanders \& Huntley 1976) are not supported by our observations, as we find no correlation between arm strength and bar strength, even in the inner part of the disk. Tidal effects from near neighbours seem to have some effect. We find that 4 out of 7 galaxies with arm EA greater than $35^{\circ}$ have near neighbours whereas only 12 out of 45 galaxies in whole sample have near neighbours. We also find that in the Fourier analysis of galaxies with near neighbours, the even low-order modes are dominant, especially the $m=2$ mode. We have also discussed factors affecting star formation in the disks of spiral galaxies. We have found evidence both for and against the Large Scale Shock Scenario (Roberts 1969). Firstly, we find a correlation between arm contrast and the normalised star formation rate, in agreement with the predictions of the large scale shock scenario. Secondly, we find that B-band arms are more loosely wound and narrower than K-band arms, confirming the predictions made by density wave theories (Roberts 1969). However, this can also be interpreted as an extinction effect because dust lies on the trailing edges of arms. Finally, and possibly most significantly, we find no correlation between the arm properties of this sample of galaxies and their classified Hubble type. Combining this with the poor correlation between Hubble type and K-band B/D ratio (Paper I), it would appear impossible to allocate a Hubble type to these galaxies on the basis of these K-band observations, and the morphology of the old stellar population varies surprisingly little between Hubble types Sa and Sd. We speculate that the determining parameter for Hubble type is cold gas content, which controls the star formation rate and the distribution of the young stars, which dominate the optical appearance. | 98 | 3 | astro-ph9803254_arXiv.txt |
9803 | astro-ph9803062_arXiv.txt | We report the initial results of a survey for intracluster planetary nebulae in the Virgo Cluster. In two $16' \times 16'$ fields, we identify 69 and 16 intracluster planetary nebula candidates, respectively. In a third $16' \times 16'$ field near the central elliptical galaxy M87, we detect 75 planetary nebula candidates, of which a substantial fraction are intracluster in nature. By examining the number of the planetaries detected in each field and the shape of the planetary nebula luminosity function, we show that 1) the intracluster starlight of Virgo is distributed non-uniformly, and varies between subclumps A and B{}, 2) the Virgo Cluster core extends $\sim 3$~Mpc in front of M87, and thus is elongated along the line-of-sight, and 3) a minimum of 22\% of Virgo's stellar luminosity resides between the galaxies in our fields, and that the true number may be considerably larger. We also use our planetary nebula data to argue that the intracluster stars in Virgo are likely derived from a population that is of moderate age and metallicity. | The concept of intracluster starlight was first proposed by Zwicky (1951), when he claimed to detect excess light between the galaxies of the Coma cluster. Follow-up photographic searches for intracluster luminosity in Coma and other rich clusters (Welch \& Sastry 1971; Melnick, White, \& Hoessel 1977; Thuan \& Kormendy 1977) produced mixed results, and it was not until the advent of CCDs that more precise estimates of the amount of intracluster starlight were made (cf.~Guldheus 1989; Uson, Boughn, \& Kuhn 1991; V\'ichez-G\'omez, Pell\'o, \& Sanahuja 1994; Bernstein \etal 1995). All these studies suffer from a fundamental limitation: the extremely low surface brightness of the phenomenon. Typically, the surface brightness of intracluster light is less than 1\% that of the sky, and measurements of this luminosity must contend with the problems presented by scattered light from bright objects and the contribution of discrete sources below the detection limit. Consequently, obtaining detailed information on the distribution, metallicity, and kinematics of intracluster stars through these types of measurements is extremely difficult, if not impossible. An alternative method for probing intracluster starlight is through the direct detection and measurement of the stars themselves. Recent observations have shown this to be possible. In their radial velocity survey of 19 planetary nebulae (PN) in the halo of the Virgo Cluster galaxy NGC~4406 (M86), Arnaboldi \etal (1996) found 3 objects with $v > 1300$~\kms; these planetaries are undoubtably intracluster in origin. Similarly, Ferguson, Tanvir, \& von Hippel (1998) detected Virgo's intracluster component via a statistical excess of red star counts in a Hubble Space Telescope (HST) Virgo field over that in the Hubble Deep Field. Finally, intracluster stars have been unambiguously identified from the ground via planetary nebula surveys in Fornax (Theuns \& Warren 1997) and Virgo (M\'endez \etal 1997; Ciardullo \etal 1998). Motivated by these results, we have begun a large scale [O~III] $\lambda 5007$ survey of intergalactic fields in Virgo, with the goal of mapping out the distribution and luminosity function of intracluster planetary nebulae (IPN){}. Depending on the efficiency of tidal stripping, Virgo's intracluster component is predicted to contain anywhere from 10\% to 70\% of the cluster's total stellar mass (Richstone \& Malumuth 1983; Miller 1983). A survey of several square degrees of Virgo's intergalactic space with a four meter class telescope should therefore detect several thousand PN, and shed light on both the physics of tidal-stripping and on the initial conditions of cluster formation. Here, we present the first results of our survey. | We report the results of a search in three fields in the Virgo Cluster for intracluster planetary nebulae, and have detected a total of 95 intracluster candidates. From analysis of the numbers of the planetaries, we find that the amount of intracluster light in Virgo is large (at least 22\% of the cluster's total luminosity), distributed non-uniformly, and varies between subclump~A and B{}. By using the planetary nebulae luminosity function, we derive an upper limit of $\sim 12$~Mpc for the distance to the front edge of the Virgo Cluster and use this to show that the cluster must be elongated along our line of sight. We also use the properties of planetary nebulae to suggest that the intracluster stars of Virgo have moderate age and metallicity. The large fraction of intracluster stars found has potentially serious consequences for models of cluster formation and evolution, and for cosmological models. Finally, we note that this survey included less than 0.2\% of the traditional 6 degree core of the Virgo Cluster. Many more intracluster planetary nebulae wait to be discovered. We thank Allen Shafter for some additional off-band observations and Ed Carder at NOAO, for his measurements of our on-band filter so that we could begin our observations on time. We would also like to thank the referee, R. Corradi, for several suggestions that improved the quality of this paper. Figure~1 was extracted from the Digitized Sky Survey, which was produced at the Space Telescope Science Institute under U.S. Government grant NAGW-2166. This work was supported in part by NASA grant NAGW-3159 and NSF grant AST95-29270. \pagebreak | 98 | 3 | astro-ph9803062_arXiv.txt |
9803 | gr-qc9803031_arXiv.txt | Binary systems comprising at least one neutron star contain strong gravitational field regions and thereby provide a testing ground for strong-field gravity. Two types of data can be used to test the law of gravity in compact binaries: binary pulsar observations, or forthcoming gravitational-wave observations of inspiralling binaries. We compare the probing power of these two types of observations within a generic two-parameter family of tensor-scalar gravitational theories. Our analysis generalizes previous work (by us) on binary-pulsar tests by using a sample of realistic equations of state for nuclear matter (instead of a polytrope), and goes beyond a previous study (by C.M.~Will) of gravitational-wave tests by considering more general tensor-scalar theories than the one-parameter Jordan-Fierz-Brans-Dicke one. Finite-size effects in tensor-scalar gravity are also discussed. | The detection of gravitational waves by kilometric-size laser-interferometer systems such as LIGO in the US and VIRGO in Europe will initiate a new era in astronomy. One of the most promising sources of gravitational waves is the inspiralling compact binary, a binary system made of neutron stars or black holes whose orbit decays under gravitational radiation reaction. The observation of these systems will provide important astrophysical information, e.g. masses of neutron stars, and direct distance measurements up to hundreds of Mpc \cite{reviews}. It is also said that detecting gravitational waves from inspiralling binaries should provide rich tests of the law of relativistic gravity in situations comprising strong-field regions (like near a neutron star or a black hole). However, present binary-pulsar experiments already provide us with deep and accurate tests of strong-field gravity \cite{taylor94,twdw92,dt92}. It is therefore interesting to compare and contrast the probing power of present (and foreseeable) pulsar tests with that of future gravity-wave observations. A convenient quantitative way of doing this comparison is to work within a multi-parameter family of physically motivated (and physically consistent) theories of gravity which differ from Einstein's theory in their radiative and strong-field predictions. The most natural framework of this type is the general class of tensor-scalar theories in which gravity is mediated both by a tensor field $g_{\mu \nu}^*$ (``Einstein metric'') and by a scalar field $\varphi$. These theories contain one arbitrary ``coupling function'' $A(\varphi)$ defining the conformal factor relating the pure spin-2 Einstein metric $g_{\mu \nu}^*$ to the ``physical metric'' $\widetilde{g}_{\mu \nu} = A^2 (\varphi) \, g_{\mu \nu}^*$ measured by laboratory clocks and rods. The usual Jordan-Fierz-Brans-Dicke theory \cite{Jordan,Fierz,BransDicke} is the one parameter class of theories defined by a coupling function $A(\varphi) = \exp (\alpha_0 \, \varphi)$. [The coupling strength $\alpha_0$ is related to the often used parameter $\omega$ by $\alpha_0^2 = (2\omega + 3)^{-1}$.] Will \cite{will94} has studied the quantitative constraints on the coupling parameter $\alpha_0$ of Jordan-Fierz-Brans-Dicke theories that could be brought by gravitational-wave observations. His result is that in most cases the bounds coming from gravity-wave observations will be comparable to presently known bounds coming from solar-system experiments (namely, $\alpha_0^2 < 10^{-3}$). This result of Ref.~\cite{will94} seems to suggest that gravitational-wave-based {\it strong-field} tests of gravity do not go really beyond the solar-system {\it weak-field} tests of gravity. We wish, however, to emphasize that this seemingly pessimistic conclusion is mainly due to having restricted one's attention to the special, one-parameter Jordan-Fierz-Brans-Dicke theory. Indeed, in this theory the strength of the coupling of the scalar field $\varphi$ to matter is given by the constant quantity $\alpha_0$ independently of the state of condensation of the gravitational source. As a consequence, the predictions of the theory differ from those of Einstein's theory by a fraction of order $\alpha_0^2$ in all situations: weak-field ones or strong-field ones, alike. By contrast, it has been emphasized in Refs.~\cite{def93,def96} that the more generic tensor-scalar theories in which the strength of the coupling of $\varphi$ to matter, namely \begin{equation} \alpha (\varphi) \equiv \frac{\partial \ln A (\varphi)}{\partial \, \varphi} \, , \label{eq1.1} \end{equation} depended on the value of $\varphi$, allowed for the existence of genuine {\it strong-field effects}, by which the presence of a highly condensed source, such as a neutron star, could generate order-unity deviations from general relativity even if the weak-field limit of $\alpha (\varphi)$ is arbitrarily small. [More precisely, these non-perturbative strong-field effects take place when $\beta (\varphi) \equiv \partial \, \alpha (\varphi) / \partial \, \varphi$ is negative.] This led us to consider, instead of the Jordan-Fierz-Brans-Dicke model $A(\varphi) = \exp(\alpha_0 \varphi)$, the class of theories defined by \begin{equation} A(\varphi) = \exp \left( \frac{1}{2} \, \beta_0 \varphi^2 \right) \, . \label{eq1.2} \end{equation} This class of theories contains two arbitrary (dimensionless) parameters: $\beta_0$ appearing in Eq.~(\ref{eq1.2}), and $\varphi_0$, the asymptotic value of $\varphi$ at spatial infinity. [By contrast, in Jordan-Fierz-Brans-Dicke theory, $\varphi_0$ has no observable effects.] Equivalently, the two parameters in these theories can be defined as being the strength of the linear coupling of $\varphi$ to matter in the weak-field limit $(\varphi \approx \varphi_0)$, \begin{equation} \alpha_0 = \alpha (\varphi_0) = \frac{\partial \ln A(\varphi_0)}{\partial \, \varphi_0} = \beta_0 \varphi_0 \, , \label{eq1.3} \end{equation} and the non-linear coupling parameter \begin{equation} \beta_0 = \frac{\partial \, \alpha (\varphi_0)}{\partial \, \varphi_0} = \frac{\partial^2 \ln A(\varphi_0)}{\partial \, \varphi_0^2} \, . \label{eq1.4} \end{equation} A convenient feature of the two-parameter family of theories (\ref{eq1.2}) is that they have just the amount of generality needed both to parametrize the most general boost-invariant weak-field (``post-Newtonian'') deviations from Einstein's theory, and to encompass nonperturbative strong-field effects. Indeed, on the one hand, the two theory-parameters $(\alpha_0 , \beta_0)$ determine the two well-known Eddington-Nordtvedt-Will parameters, \begin{equation} \overline{\gamma} \equiv \gamma_{\rm Eddington} - 1 = -2 \, \alpha_0^2 / (1 + \alpha_0^2) \, , \label{eq1.5} \end{equation} \begin{equation} \overline{\beta} \equiv \beta_{\rm Eddington} - 1 = \frac{1}{2} \, \beta_0 \, \alpha_0^2 / (1 + \alpha_0^2)^2 \, , \label{eq1.6} \end{equation} which measure the most general, boost-invariant, deviations from general relativity at the first post-Newtonian level. (See \cite{def92,def2pn} for the generalization of the Eddington parameters to the second post-Newtonian level.) On the other hand, when $\beta_0 \lesssim -4$ non-perturbative strong-field effects develop in the theories defined by Eq.~(\ref{eq1.2}). All existing gravitational experiments can be interpreted as constraints on the two-dimensional space of theories defined by Eq.~(\ref{eq1.2}). In other words, we can work within the common $(\alpha_0 , \beta_0)$ plane, and consider each gravitational experiment (be it of weak-field or strong-field nature) as defining a certain exclusion plot within that plane. In some recent work \cite{def96}, we have constructed such exclusion plots, as defined by considering both solar-system experiments and binary-pulsar experiments. The present work will generalize these exclusion plots in several respects: (i)~we shall plot the regions of the $(\alpha_0 , \beta_0)$ plane probed by future gravitational-wave observations of neutron star--neutron star, and neutron star--black hole systems, (ii)~we shall improve our previous study of the probing power of binary-pulsar experiments by considering a sample of realistic nuclear equations of states (instead of the polytropic one used by us before) and by using updated pulsar data, and (iii)~we shall consider the individual constraints on $\alpha_0$ and $\beta_0$ brought by the main solar-system experiments (instead of using published combined limits on $\overline{\gamma}$ and $\overline{\beta}$). This paper is organized as follows: In Sec.~II we summarize our (numerical) approach to computing the various form factors that describe the coupling of the scalar field $\varphi$ to a neutron star. In Sec.~III we generalize Ref.~\cite{will94} in discussing how gravitational wave observations can give quantitative tests of tensor-scalar gravity. In Sec.~IV we combine and compare gravitational-wave tests with binary-pulsar tests and solar-system ones. Our conclusions are presented in Sec.~V, while an Appendix discusses finite-size effects in tensor-scalar gravity. | The main conclusion of the comparison carried out in Figs.~\ref{fig1}--\ref{fig5} is that, in all cases, future LIGO/VIRGO observations of inspiralling compact binaries turn out not to be competitive with present binary-pulsar tests in their {\it discriminating\/} probing of the strong-field, and radiative, aspects of relativistic gravity. This conclusion may seem paradoxical. It should not be interpreted negatively against LIGO/VIRGO observations which, as shown in Figs.~\ref{fig1}--\ref{fig5}, will independently probe strong-field gravity and will exclude regions of parameter space allowed by solar-system experiments. Rather, it is simply a reminder that binary-pulsar experiments are superb tools for probing strong-field and radiative aspects of gravity. It is also somewhat a good news for gravitational wave data analysis (which promises to be already a very challenging task even if one {\it a priori\/} assumes the validity of general relativity; see, e.g., \cite{dis98}). Indeed, our results Figs.~\ref{fig1}--\ref{fig5} indicate that our present experimental knowledge of the law of relativistic gravity is sufficient to justify using general relativity as the standard theory of gravitational radiation. Note that this conclusion explicitly refers to the quantitative probing of plausible\footnote{Within the presently accepted framework for low-energy fundamental physics, namely field theory, the only alternative (long-range) gravity theories which, (i) do not violate the basic tenets of field theory, and (ii) are not already necessarily extremely constrained by existing equivalence-principle tests, are the tensor-scalar gravity theories.} {\it deviations\/} {}from general relativity. At the {\it qualitative\/} level, and also at the {\it non-discriminating\/} quantitative level, LIGO/VIRGO observations will bring invaluable advances in our experimental knowledge of relativistic gravity. First, they will provide the first direct observation of gravitational waves far in the wave zone (while binary pulsar experiments prove the reality of the propagation with finite velocity of the gravitational interaction in the near zone of a binary system). Second, they will (hopefully) lead to superb additional confirmations of general relativity through the observation of the wave forms emitted during the inspiral and coalescence of neutron stars or black holes; see, e.g., \cite{reviews,bs9495}. Independently of the comparison between LIGO/VIRGO tests and binary-pulsar tests, the present paper has provided the first systematic study of the influence of the nuclear equation of state on the theoretical probing power of binary-pulsar tests. In particular, Fig.~\ref{fig2} shows that if the equation of state were as soft as predicted by the simple Pandharipande model, binary-pulsar tests quantitatively supersede solar-system ones in all the region $\beta_0 \lesssim -1$ of parameter space. Even if we consider the less constraining stiff equations of states, the present work confirms the limit \begin{equation} \beta_0 > -4.5 \label{eq5.1} \end{equation} found (modulo $10\%$) in \cite{def96}. We recall that this limit can be interpreted as a limit on the ratio of the two weak-field post-Einstein parameters \begin{equation} \frac{\beta_{\rm Edd} - 1}{\gamma_{\rm Edd} - 1} \approx - \frac{1}{4} \, \beta_0 < 1.1 \, . \label{eq5.2} \end{equation} Finally, we pointed out that the discovery of a binary pulsar with a black-hole companion has the potential of providing a superb new probe of relativistic gravity. The discriminating power of this probe might supersede all its present and foreseeable competitors in measuring $\alpha_0$ down to the level $\alpha_0^2 < 10^{-6}$. | 98 | 3 | gr-qc9803031_arXiv.txt |
9803 | astro-ph9803289_arXiv.txt | We propose a model for the source of the X-ray background (XRB) in which low luminosity active nuclei ($L\sim 10^{43}\ergps$) are obscured ($N\sim 10^{23}\, {\rm cm}^{-2}$) by nuclear starbursts within the inner $\sim 100$~pc. The obscuring material covers most of the sky as seen from the central source, rather than being distributed in a toroidal structure, and hardens the averaged X-ray spectrum by photoelectric absorption. The gas is turbulent with velocity dispersion $\sim {\rm few}\times 100\, \kmps$ and cloud-cloud collisions lead to copious star formation. Although supernovae tend to produce outflows, most of the gas is trapped in the gravity field of the starforming cluster itself and the central black hole. A hot ($T\sim 10^6-10^7$~K) virialised phase of this gas, comprising a few per cent of the total obscuring material, feeds the central engine of $\sim 10^7\, \Msun$ through Bondi accretion, at a sub-Eddington rate appropriate for the luminosity of these objects. If starburst-obscured objects give rise to the residual XRB, then only 10 per cent of the accretion in active galaxies occurs close to the Eddington limit in unabsorbed objects. | The flat spectrum of the X-ray background (XRB) above 1~keV (Marshall et al 1980; Gendreau et al 1995; Chen, Fabian \& Gendreau 1997) is not simply accounted for by the integration of known classes of source, which generally have much steeper spectra. A currently popular model invokes intrinsic absorption in active galaxies with which to flatten the observed X-ray spectrum (Setti \& Woltjer 1989; Madau, Ghisellini \& Fabian 1994; Celotti et al 1995; Comastri et al 1995). The intrinsic absorption must range from column densities of about $10^{22} - 10^{24}\psqcm$, and must cover most (at least $\sim 2/3$)\footnote{Since the XRB has a flat spectrum right down to 1~keV, the fraction of sources with absorption exceeding $10^{21}\, \pcmsq$ must approach 90 per cent; see Section 3} of the sky as seen by the source itself. The combination of different levels of absorption and redshift of the objects can then lead to the observed power-law background spectrum in the 1--10~keV band. The rollover in the spectrum at about 30~keV is due redshift acting on an intrinsic 100~keV break, such as is seen in nearby active galaxies (Zdziarski et al 1995). The geometry of the obscuring material within a typical, X-ray background-contributing, active galaxy is unclear. The nucleus must be mostly surrounded by a typical column density of say $10^{23}\psqcm$. If the material is freely orbiting the nucleus, then it should soon collide with itself, dissipate, and flatten into a disk, unless either a) the orbits are carefully arranged or b) there is sufficient energy to continually throw matter into a wide range of orbital inclinations. We note that what evidence there is for a torus of molecular material around active galaxies does not indicate that its high column density part covers much of the Sky. Indeed, models where the obscuring material is extended on scales $\sim 100\, \pc$ (Granato et al 1997) account for the infrared properties of Seyfert galaxies more successfully than ones which postulate a thick compact ($<1\, \pc$) torus (Pier \& Krolik, 1992, 1993). The first case a) may be accounted for by a warped disk. Pringle (1996) and Maloney, Begelman \& Pringle (1996) have shown that the outer parts of irradiated accretion disks are unstable to warping, the final result of which is that much of the sky is covered by the warp. The details are currently uncertain as to whether most of the sky can be covered and whether such high column densities can be attained in the warped material. We do not pursue that further here. The second case b) may be accounted for by a nuclear starburst, the supernovae of which can provide the energy to push clouds around and so obscure the nucleus. It is that possibility we explore further here. There is indeed much empirical evidence for a connection between nuclear starbursts and active galactic nuclei (Terlevich \& Melnick 1985, Terlevich, D\'{\i}az \& Terlevich 1990; Perry \& Williams 1993, Heckman et al 1995, 1997; Maiolino et al 1997). The extra featureless ultraviolet continuum in Seyfert 2 galaxies (FC2) appears to be due to a nuclear starburst (Terlevich \& Cid-Fernandes 1995; Heckman et al 1997). Turner et al (1997, 1998) in a discussion of the absorption properties of Seyfert 2 galaxies propose that the starburst may be involved there. The striking similarity between the evolution of the luminosity density due to star formation (Madau et al 1996) and that due to QSOs (Boyle \& Terlevich 1998, Dunlop 1998) also suggests a close link between star formation processes and the fueling of QSO accretion. In addition, some sort of starburst activity is inevitable when an AGN forms, since vast amounts of material (e.g., triggered by merging) are funneled into the core of the galaxy giving rise ultimately to the central engine. It is then likely that this extended obscuring starburst is a common feature of all AGN during at least some phase, although for the most massive and luminous objects it might last only for a short time. We make here the proposal that the formation of massive stars in the inner $\sim 100$ pc of a central massive black hole is instrumental in both fuelling the central engine by accretion and obscuring it by distributing cold clouds all around that engine. Since most of the X-ray accretion power in the Universe emerges in the XRB (see e.g. Chen et al 1997) then it is in the obscured starburst mode that most of it takes place. An obscuring starburst also accounts for the optical, narrow-line, appearance of a population of faint hard X-ray sources (NLXGs; Boyle et al 1995; Roche et al 1995; Carballo et al 1995; Griffiths et al 1996; McHardy et al 1998; Hasinger 1996; Almaini et al 1996; Romero-Colmenero et al 1996) which dominate at faint flux levels and so may provide most of the XRB (see, however, Hasinger et al 1998 and Schmidt et al 1998 who cast some doubts on the reality of these objects as a class different to AGN). The broad-line Seyfert spectrum will not be detectable, unless we have a favourable line of sight, or observe in the mid-infrared in an object where the column density is not too high (Granato, Danese \& Franceschini 1997). The combination of a starburst spectrum and a Seyfert narrow-line spectrum will make classification of the objects ambiguous (Iwasawa et al 1997b). There are various examples of galaxies whose spectrum is dominated by a starburst/LINER component at all wavelengths except at hard X-rays where the AGN shows up, obscuration preventing its detection at lower energies. Among them, NGC4945 is obscured by a column $\sim 5\times 10^{24}\, \pcmsq$ (Iwasawa et al 1993) and NGC 6240 by a column $> {\rm few}\times 10^{24}\, \pcmsq$ (Iwasawa \& Comastri 1998). | Nuclear starbursts can play a major role in obscuring low-luminosity AGN. The emerging X-ray spectrum will be moderately to highly absorbed by the cold and warm gas in the star cluster, thus providing the flat spectral shape that is needed to produce the XRB. The column of absorbing gas is likely to vary vith viewing direction in a `random' way due to supernovae produced in the starburst. The average high fraction of the sky that the majority of the faint X-ray sources have to see covered by absorbing gas arises in material related to the star cluster rather than in a hypothetical torus, whose thickness to radius ratio must be large. Arguments favouring a geometry for the obscuring material more isotropic than toroidal have been put forward before. Turner et al (1998) find no obvious dependence of the X-ray properties of NLXGs with the orientation of the host galaxy. Iwasawa et al (1995) also suggested that the obscuring material in the Seyfert 2 galaxy $IRAS$ 18325-5926, which is obscured by a column in excess of $10^{22}\, \pcmsq$ but shows no X-ray reflected component, is likely to be isotropically distributed around the central engine. Indeed the question remains, as in all obscured AGN models for the XRB, as to why the distribution of obscuring material along different lines of sight in the population of these objects is exactly tuned in such a way that it gives rise to the featureless spectrum of the XRB. There is no physical principle which leads to this and some fine-tuning of this distribution is required to produce a detailed model for the XRB. The optical/UV properties of the highly absorbed AGN will be dominated by the starburst itself. Collisions between cold clouds, occuring at several $\times 100 \kmps$ will lead to radiative shocks with some emission at soft X-ray energies. Therefore, the obscuring material will also contribute to the observed soft X-ray emissivity, making some of these objects visible in the ROSAT PSPC band in spite of being absorbed. Near infra-red spectroscopy of NLXGs also lends support to this hybrid (i.e., accretion onto a black hole plus an obscuring starburst) hypothesis. While a number of NLXGs show unambiguous evidence for hidden AGN (eg. highly ionized coronal lines), the non-detection of broad Paschen $\alpha$ requires obscuring columns with $N_H > 10^{23}$~cm$^{-2}$ (Almaini et al 1998). Such thick columns are inconsistent with the large X-ray luminosities observed in ROSAT, unless there is an additional source of soft X-ray flux. A contribution from the activity in the obscuring starburst provides a natural explanation. Gas columns along typical lines of sight in these objects are modest and for normal dust to gas ratios the absorbing material will be optically thin in the mid- and far-infrared (Granato et al 1997) where most of UV and soft X-ray reprocessed radiation will be emitted in a rather isotropic fashion. Surveys at these wavelengths (particularly at $10-15\, \mu$m) should be especially efficient in finding these objects. Barcons et al (1995) measured the average ratio of X-ray (2-10~keV) to mid-infrared emission (the {\it IRAS} 12$\mu$m band) for 12$\mu$m-selected type 1 and type 2 AGN. The starburst obscured AGN are likely to have $f(5\, \keV)/f(12\, \mu{\rm m})\sim 2\times 10^{-7}$, flux ratio that was found for Seyfert 2 galaxies, and therefore an obscured AGN with a 2-10~keV X-ray flux of $10^{-14}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ is likely to emit at a level of $\sim 3$~mJy at 10-15$\mu$m. The {\it European Large Area Infrared Space Observatory Survey} (ELAIS, Oliver et al 1997) is reaching a sensitivity of $\sim 2$~mJy at 15$\mu$m over a survey area of 13${\rm deg}^2$ (some 1000-2000 objects are expected in total). Even if at a 2-10~keV flux of $10^{-14}\, \ergpcmsqps$ only a minor fraction of the sources are starburst obscured AGN, they will certainly show up in large numbers in the ELAIS survey. The recently commissioned Sub-mm Common User Bolometer Array (SCUBA) at the James Clerk Maxwell Telescope in Hawaii offers an ideal opportunity to detect obscured AGN at high redshifts. The re-radiation of nuclear energy in the far infra-red is expected to produce a thermal emission peak at $\lambda \sim 60-100\mu$m, which is redshifted into the sub-millimetre region. A starburst-obscured AGN at redshift $z=2$ with a 2-10~keV X-ray flux of $5\times10^{-15}\, \ergpcmsqps$ is expected to give a 350$\mu$m flux of $\sim 20$mJy. Deep SCUBA observations of hard X-ray selected fields (eg. with ASCA, AXAF or XMM) should reveal these obscured objects. As stated, we expect that the very massive black holes will appear as quasars. Sometimes however a massive black hole ($M_{BH}>10^8\Msun$) may have a more massive surrounding star cluster in which case our scenario may apply and the active nucleus will be hidden. This may explain the low level of X-ray emission from IRAS galaxies in general and the hyperluminous galaxies (e.g. IRAS F15307+3252) in particular (Fabian et al 1996). The X-ray upper limit to that object requires less than $2\times 10^{-4}$ of its power to emerge as X-rays which means that any quasar-like nucleus must be smothered from all directions. \begin{figure} \centerline{\psfig{figure=obscured.fig2.ps,width=0.5\textwidth,angle=270}} \caption{The energy content of the XRB per unit logarithmic energy interval (solid line) as a function of photon energy. The dashed curve shows the contribution of unabsorbed AGN, where it is assumed that they make 50 per cent of the XRB at 1~keV. There is probably a roll-over in their contribution at energies $> 30-50\, \keV$ which is the exponential cutoff observed in local AGN at $\sim 100\, \keV$ redshifted out to $z\sim 2$. The dotted line shows the total energy produced by all objects, before obscuration is taken into account.} \end{figure} A geometrically thin disk structure is expected in obscured AGN in our model on scales of $<1\pc$ where angular momentum begins to be important. Compton scattering, in some cases in an optically-thick medium, is necessary to explain the X-ray emission of some Seyfert 2 galaxies (see Iwasawa et al 1997a for a detailed study of NGC 1068 and Turner et al 1997 for Mrk 3 among other examples). In the case of NGC 1068 the pc-scale disk structure has been mapped at radio wavelengths (Gallimore, Baum \& O'Dea 1997). Our model can therefore account for both Compton-thin and Compton-thick Seyfert 2 galaxies. Much of the extended torus attributed to these objects is just the starburst obscuration, but a planar Compton-thick inner part lies along our line of sight in some of them. If the model proposed here describes correctly the sources of the residual XRB, it has then important implications on how accretion occurs in the Universe. Fig.~2 shows the distribution of the energy in the XRB as a function of photon energy in the range 1 to 100 keV (Fabian \& Barcons 1992). Unobscured AGN (QSOs and Seyfert 1s) make $\sim 50$ per cent of the XRB at 1~keV (Hasinger 1996, McHardy et al 1998, Hasinger et al 1998, Schmidt et al 1998), with an energy spectral index $\sim 1$. If the `shoulder' in the XRB (Fig.~2) is due to obscuration (mostly photoelectric absorption of the photons with energies $< 30\, \keV$), the energy content at 30 keV provides a measure of the total energy produced by accretion in AGN, for the same underlying power law spectrum. What is remarkable then, is that only 10 per cent of the accretion occurs in unobscured objects (assuming that before it is absorbed by the surrounding material, the spectrum generated by accretion in both cases is similar). The remaining 90 per cent occurs at sub-Eddington accretion rates, and a relevant fraction is absorbed and re-radiated at longer wavelengths, particularly in the infrared. This means that most of the sky as seen from an average active nucleus is obscured. Moreover, given the requirement that these low luminosity AGN are at least one order of magnitude fainter than typical broad line objects (see Section 2), we predict that starburst obscured AGN may outnumber brighter, unobscured AGN by a factor of 100 or more if they are to account for the remainder of the XRB. This estimate is in good agreement with the number count predictions for QSOs and NLXGs based on extrapolating deep ROSAT observations (Almaini \& Fabian 1997). If most of accretion in the Universe is highly obscured, then the amount of emitted power per unit galaxy based on optical or UV QSO luminosity functions (So{\l}tan 1982, Phinney 1997), and therefore the mass in black holes in AGN, might have been underestimated. This is due to the fact that 90 per cent of the accretion power would be obscured and re-radiated in the infrared. Also, if most AGN are Advection Dominated Accretion Flows (ADAFs), as it has been proposed by Di Matteo \& Fabian (1997) as the source of the XRB, the implied low mass to energy conversion efficiency also means that the black hole masses would have to be larger, in agreement with the local estimates. | 98 | 3 | astro-ph9803289_arXiv.txt |
9803 | astro-ph9803130_arXiv.txt | \noindent High resolution, deep radio images are presented for two distant radio galaxies, 3C324 ($z=1.206$) and 3C368 ($z=1.132$), which are both prime examples of the radio--optical alignment effect seen in powerful radio galaxies with redshifts $z \gta 0.6$. Radio observations were made using the Very Large Array in A--array configuration at 5 and 8\,GHz, and using the MERLIN array at 1.4 and 1.65\,GHz. Radio spectral index, radio polarisation, and rotation measure maps are presented for both sources. \noindent Radio core candidates are detected in each source, and by aligning these with the centroid of the infrared emission the radio and the optical\,/\,infrared images can be related astrometrically with 0.1 arcsec accuracy. In each source the radio core is located at a minimum of the optical emission, probably associated with a central dust lane. Both sources also exhibit radio jets which lie along the directions of the bright strings of optical knots seen in high resolution Hubble Space Telescope images. The northern arm of 3C368 shows a close correlation between the radio and optical emission, whilst along the jet direction of 3C324 the bright radio and optical knots are co--linear but not co--spatial. These indicate that interactions between the radio jet and its environment play a key role in producing the excess ultraviolet emission of these sources, but that the detailed mechanisms vary from source to source. \noindent 3C368 is strongly depolarised and has an average rest--frame rotation measure of a few hundred rad\,m$^{-2}$, reaching about 1000\,rad\,m$^{-2}$ close to the most depolarised regions. 3C324 has weaker depolarisation, and an average rest--frame rotation measure of between 100 and 200\,rad\,m$^{-2}$. Both sources show large gradients in their rotation measure structures, with variations of up to 1000\,rad\,m$^{-2}$ over distances of about 10\,kpc. | \label{intro} The discovery that the optical, ultraviolet and line emission of powerful high redshift radio galaxies is elongated and aligned along the direction of the radio axis (McCarthy \etal\ 1987, Chambers \etal\ 1987)\nocite{mcc87,cha87}, a phenomenon known as the `alignment effect', indicates a close association between radio source activity and the optical emission of these galaxies. A number of processes have been proposed to account for these radio--optical correlations (see McCarthy 1993 for a review),\nocite{mcc93} the most promising models falling into two broad categories: interaction models in which the correlations arise directly through interactions between the radio jet and its environment, and illumination models in which radiation from a partially obscured active galactic nucleus (AGN) both photo--ionises the emission line gas within an ionisation cone and is scattered by dust grains or electrons producing the optical and ultraviolet alignment. One of the earliest models for the alignment effect, which has remained popular, is star--formation induced by the passage of the radio jets (e.g. Rees 1989).\nocite{ree89} This process has been observed in jet--cloud interactions at low redshifts, for example in Minkowski's object \cite{bro85,bre85} and in the radio lobe of 3C285 \cite{bre93}. There is quite convincing evidence that it also occurs in at least some distant radio galaxies (e.g. Best \etal\ 1997a, Dey \etal\ 1997)\nocite{bes97b,dey97} but, as yet, direct evidence for the presence of young stars in the aligned structures is limited. A large proportion of the powerful distant radio galaxies which have been studied using spectropolarimetry have optical emission which is polarised at up to the 15\% level, and scattered broad lines in their spectra (e.g. Dey \etal\ 1996, Cimatti \etal\ 1996, 1997, and references therein)\nocite{cim96,dey96,cim97a}. A significant proportion of the aligned emission of these galaxies must therefore be scattered light from an obscured AGN. In an optically unbiased sample of distant radio galaxies, however, Tadhunter \etal\ \shortcite{tad97} detected significant polarisation in only 40\% of the galaxies although most showed large ultraviolet excesses. Nebular continuum emission provides another important alignment mechanism. Dickson \etal\ \shortcite{dic95} showed that a combination of free--free, free--bound, and two--photon continuum, as well as the Balmer forest, can contribute a significant proportion (5--40\%) of the ultraviolet emission from the nuclear regions of powerful radio galaxies, and this may be even higher in the extended aligned emission regions (e.g. Stockton \etal\ 1996)\nocite{sto96a}. Since either or both of photo--ionisation by the AGN and shock--ionisation by the radio jet may be responsible for exciting the gas, distinguishing between the jet--cloud interaction model and the illumination model is made more difficult. The structure and properties of the radio emission from these distant radio galaxies can offer important clues to the nature of the radio--optical correlations. Hubble Space Telescope (HST) observations of 28 radio galaxies at redshifts $z \sim 1$, selected from the revised 3CR catalogue of Laing \etal\ \shortcite{lai83}, have shown that for the majority of small radio sources the alignment effect manifests itself in the form of strings of bright knots tightly aligned between the two radio lobes (Best \etal\ 1996, 1997b).\nocite{bes96a,bes97c} Deep, high resolution radio observations of these sources provide a critical test of the interaction\,/\,illumination models: if the locations of the radio jets and the strings of optical activity are co--linear, then the influence of the radio jets must play a key role in the alignment effect. This would not, however, provide direct proof of the jet--induced star formation hypothesis: the emission may be nebular continuum emission from warm gas shocked by the radio jets, or may be associated with an increased scattering efficiency in these regions, for example by jet shocks breaking up cold gas clouds and exposing previously hidden dust grains, thereby increasing the surface area for scattering along the jet axis \cite{bre96b}. One of the major problems in interpreting the radio--optical structures of distant radio galaxies has been accurate relative astrometry of the radio and optical\,/\,infrared frames of reference. Radio observations are automatically in the International Celestial Reference Frame (ICRF), with positional errors of about 0.01 arcsec, but the absolute positions of the HST images are uncertain at the 1 arcsecond level \cite{lat97}. Such astrometric uncertainties mean that bright radio knots may be either correlated or anti--correlated with the position of the luminous optical emission regions. This problem can be alleviated somewhat by multi--frequency radio observations of sufficient sensitivity to detect a flat--spectrum central radio core: this core is expected to be coincident with the nucleus of the galaxy, which itself is located roughly at the centroid of the infrared image of the galaxy. The optical (rest--frame ultraviolet) emission cannot be used for this procedure, since it is often severely affected by both dust extinction (e.g. de Koff \etal\ 1996, Best \etal\ 1997b)\nocite{bes97c,dek96} and the aligned emission. The infrared images also contain a component of aligned emission (e.g. Eisenhardt \& Chokshi 1990; Rigler \etal\ 1992; Dunlop \& Peacock 1993; Best \etal\ 1997b,1998)\nocite{eis90,rig92,dun93,bes98d} but at a significantly lower level, and are seen to show a sharp central peak which can be astrometrically aligned with the radio core. It should be noted that this process can never be 100\% reliable, owing to the ambiguity of identifying radio core components based upon spectral index information alone; the case of 3C356 (Fernini \etal\ 1993, Best \etal\ 1997a) is a case in point. For the radio galaxies in the current paper, however, further radio--optical correlations are observed when this process is employed (see Sections 3 and 4), indicating that the core identifications are most likely correct. The properties of the radio emission also provide a probe of the large--scale physical environment of the radio source. There has been growing evidence that powerful radio galaxies at large redshifts lie in rich (proto--) cluster environments, based upon measures of galaxy cross--correlation functions (e.g. Yates \etal\ 1989)\nocite{yat89}, detections of powerful extended X--ray emission from the vicinity of distant radio sources (e.g. Crawford and Fabian 1996)\nocite{cra96b}, detections of companion galaxies in narrow--band images and with spectroscopy (e.g. McCarthy 1988, Dickinson \etal\ 1997)\nocite{mcc88,dic97a}, and the very bright infrared magnitudes and large characteristic sizes of the radio galaxies themselves, suggesting that they are as large and luminous as brightest cluster galaxies (see Best \etal\ 1998 for a review)\nocite{bes98d}. Radio data provide further evidence in support of this hypothesis. At low redshifts, radio galaxies in rich cluster environments, such as Cygnus A, tend to have very large rotation measures, $RM \gta 1000$\,rad\,m$^{-2}$ \cite{dre87,tay94}. Carilli \etal\ \shortcite{car97} have studied a sample of 37 radio galaxies with redshifts $z > 2$, and detect rotation measures exceeding 1000\,rad\,m$^{-2}$ in 19\% of them. They consider this percentage to be a conservative lower limit, and interpret these large rotation measures as arising from hot, magnetised (proto--)cluster atmospheres surrounding the sources, with field strengths $\sim 1$\,nT. Detailed studies of individual sources with very high rotation measures support this interpretation \cite{car94a,pen97}. In comparison with these very distant sources, the polarisation properties of radio galaxies with redshift $z \sim 1$ have not been well studied at high angular resolution (although see Pedelty \etal\ 1989, Liu and Pooley 1991, Fernini \etal\ 1993, and Johnson \etal\ 1995)\nocite{ped89a,joh95,fer93,liu91a}. We have selected two radio galaxies from the 3CR sample, with redshifts in excess of one: 3C324 ($z=1.206$) and 3C368 ($z=1.132$). These galaxies are both prime examples of the alignment effect; their Hubble Space Telescope (HST) images show highly extended optical morphologies with strings of bright knots close to the radio axis \cite{lon95}. We have carried out sensitive, multi--frequency, polarimetric radio observations of these radio sources at high angular resolution using the VLA and MERLIN. A description of the observations and the data reduction techniques is given in Section~\ref{observs}. In Section~\ref{sec3c324}, we present the radio data for 3C324, and make a comparison with the optical and infrared images. This is repeated for 3C368 in Section~\ref{sec3c368}. These sources are compared in Section~\ref{concs}, in which we also summarise our conclusions. Throughout the paper we assume $H_0 = 50$\,km\,s$^{-1}$\,Mpc$^{-1}$ and $\Omega = 1$, and all positions are given in equinox J2000 coordinates. | \label{concs} We summarise below the main observational results of this work: \begin{enumerate} \item We have detected radio core candidates in each of 3C324 and 3C368. Aligning these relatively flat spectrum knots with the centre of the infrared images enables the radio and the optical\,/\,infrared frames of references to be astrometrically aligned to an accuracy of about 0.1 arcsec. \item The radio cores of each source lie at a minimum of the optical emission, probably associated with obscuration by a central dust lane. Dust may also play an important role in determining the extended morphology of the optical emission. \item Radio jets are detected in both sources, and in each case are co-linear with the strings of bright optical knots. \item These observations sort out the disparity between lower angular resolution observations of the polarisation structures of these radio galaxies \cite{ped89a,fer93}. The higher angular resolution reveals very large gradients in the rotation measure structures, reaching up to 1000\,rad\,m$^{-2}$ over a distance of order 10\,kpc, and regions of strong depolarisation. \item There is a tight correlation between the depolarisation measures and the gradients in the rotation measure, suggesting that the depolarisation is external, and caused by variations within each beamwidth of the Faraday depth of material in the vicinity of the radio source. \end{enumerate} The striking co-linearity of the radio jet with the bright optical knots in the two sources indicate that interactions of the radio jet itself, rather than just the AGN or the radio cocoon, must play a critical role in the origin of the aligned optical emission. Of interest is that, given these strong interactions between the radio jets and their environment, more enhanced radio emission is not seen in these regions, in contrast to that observed in bright radio knots at the sites of jet--cloud interactions at low redshifts (e.g. 4C29.30, van Breugel \etal\ 1986). This suggests that the more powerful radio jets in these high redshift sources are not as strongly disrupted by their interactions. Despite the many similarities between the two sources, there are equally important differences. The northern arm of 3C368 shows a tight correlation between the regions of radio and optical emission, whilst in 3C324 although these are co-linear they are not co-spatial. The nature of the optical (rest--frame ultraviolet) emission of the two galaxies is also very different. 3C324 has a high percentage of spatially extended optical polarisation, $P \approx 11\%$, indicating that illumination by the central AGN must be important in this source \cite{cim96}. In contrast, HST and Keck observations of 3C368 have detected no polarised optical emission \cite{bre96c}. 3C368 is an exceptional source, being probably the most aligned of all of the 3CR radio galaxies and possessing the highest emission line flux. The extent of its radio emission is 73\,kpc, only slightly greater than that of the optical emission. Its structure may be interpreted in terms of it being a relatively young radio source (a few times $10^6$ years), in which on--going interactions between the radio jet and the interstellar medium remain the dominant effects. In the northern region of this source, where the bright optical knots lie coincident with the radio jet, the nebular continuum emission is very luminous and the emission line velocities are high \cite{dic95,sto96a}. Interactions between the radio jet and its environment, possibly including a companion galaxy (see Section~\ref{368radopt}) may brighten the radio jet and accelerate and shock--ionise the warm emission line gas, explaining all of these features. The strong correlation between the radio and optical structures cannot be explained if the gas responsible for the nebular continuum is photoionised by the AGN. Jet--induced star formation in this region cannot be excluded, but there is no direct evidence for this. Any small contribution of scattered light might not be detectable due to the dilution of the polarisation by the very strong nebular continuum contribution. The radio emission of 3C324 is more extended and reaches well beyond the region of bright optical emission. In this source, we may be observing the residual effects after the radio jets have forced a passage through the host galaxy. In the aftermath of the jet shocks the nebular continuum contribution has decreased, and now some scattered light is detected. Cimatti \etal\ \shortcite{cim96} calculate that at most 30 to 50\% of the optical flux density can be associated with scattered radiation, with the rest likely to be nebular continuum emission or light from a young starburst. They suggest that the knot close to the western lobe may currently be undergoing a burst of star formation at a rate of $70 M_{\odot}$\,yr$^{-1}$. Most of the optical emission from 3C324 arises from the bright knotty structures along the radio axis, with no evidence for the scattered emission being evenly distributed over an ionisation cone structure. Therefore, the scattering material in this source must also be preferentially located along the jet direction. Two possible mechanisms for this latter effect would the production of dust in a region of star formation induced by the radio jet, or the disintegration of cooled clumps of optically thick gas by the jet, and the exposure of previously hidden dust grains at their centre \cite{bre96b}. \smallskip At low redshifts, rotation measure values and variations in excess of a few hundred rad\,m$^{-2}$ are seen in two types of radio source, compact steep spectrum radio sources in which the radio lobes lie within the host galaxy (e.g. Taylor \etal\ 1992)\nocite{tay92}, and FR\,II radio sources located in clusters with luminous X--ray cooling flows. For the former, Garrington and Akujor \shortcite{gar96} have suggested that the rotation measure correlates with linear size, with the dense gas located towards the centre of the host galaxy being responsible for the Faraday rotation. For the latter, Taylor \etal\ \shortcite{tay94} found a correlation between the rotation measure and the cooling flow rate, suggesting that the intracluster medium is responsible. We have argued that the large rotation measures observed in the southern lobe of 3C368 are due to emission line gas associated with the host galaxy. The northern lobe of this source, however, and both lobes of 3C324, lie distant from the host galaxy. Dickinson \shortcite{dic97a} has reported the detection of extended X--ray emission at a luminosity of $L_{\rm X} = (8.1 \pm 1.6) \times 10^{44}$\,erg\,s$^{-1}$ from 3C324, a value comparable to that of the Coma cluster. He associates this with cooling intracluster gas. Similarly, 3C368 has been detected using the ROSAT PSPC, with a signal--to--noise ratio of 5, giving $L_{\rm X} \sim 1.7 \times 10^{44}$\,erg\,s$^{-1}$ \cite{cra95}. If these X--ray luminosities are associated with cooling flows, then the cooling flow rates and rotation measures lie roughly upon the correlation of Taylor \etal\ \shortcite{tay94}, suggesting that the high rotation measures of 3C324 and of the regions of 3C368 away from the emission line gas are indeed associated with their intracluster gas. These rotation measure values are not as extreme as, for example, the 20000\,rad\,m$^{-2}$ seen in 3C295, at the centre of a very rich cluster of galaxies at $z = 0.461$ \cite{per91}, but the large values and gradients provide further support for the hypothesis that powerful distant radio galaxies lie in at least moderately rich young cluster environments. | 98 | 3 | astro-ph9803130_arXiv.txt |
9803 | chao-dyn9803022_arXiv.txt | The energy spectrum and the nolinear cascade rates of MHD turbulence is not clearly understood. We have addressed this problem using direct numerical simulation and analytical calculations. Our numerical simulations indicate that Kolmogorov-like phenomenology with $k^{-5/3}$ energy spectrum, rather than Kraichnan's $k^{-3/2}$, appears to be applicable in MHD turbulence. Here, we also construct a self-consistent renomalization group procedure in which the mean magnetic field gets renormalized, which in turns yields $k^{-5/3}$ energy spectrum. The numerical simulations also show that the fluid energy is transferred to magnetic energy. This result could shed light on the generation magnetic field as in dynamo mechanism. | The fluid parcels in a turbulent flow have random motion. However, the random flow velocities obeys certain properties. Since most of the talks in the conference dealt with chaos which also yields random signals, it is important to contrast the difference between chaos and turbulence. Chaos can occurs in a nonlinear system with a few (3-10) degrees of freedom, but a turbulent system has many (order millions or more) degrees of freedom. Also, a turbulent system may be chaotic; may not be chaotic such as in structures (e.g., vortex street); or it may have both chaos and structure coexisting with each other. Turbulence is ubiquitous and is of very practical importance. Some of the major applications are in mixing, aeroplane and high speed vehicle design, atmospheric flows and weather prediction, astrophysical objects like stars, jets, interplanetary medium etc. Even with its wide industrial, practical, and theoretical importance, the understanding of turbulence is very weak. There are many empirical laws from experiments and simulations. There are some phenomenologies, the most famous among these is by Kolmogorov. There only a few mathematically rigorous calculations, primarily Kraichnan's direct interaction approximation, calculations based on renormalization group etc. In this paper we will focus on some of the statistical properties of velocity and magnetic fields in a turbulent magnetohydrodynamic (MHD) plasma. We have attempted to review in a concise manner some of the phenomenological, numerical, analytical work, and observations from the solar wind, with an emphasis on our work (with D. A. Roberts, M. L. Goldstein, J. K. Bhattacharjee, V. Eswaran). Outline of the paper is as follows. Since the existing MHD turbulence phenomenologies are motivated by Kolmogorov's fluid turbulence phenomenology, we review Kolmogorov's arguments in section 2. In section 3 we review the existing MHD turbulence phenomenologies. Section 4 contains the numerical results, which are compared with the predictions of the turbulence phenomenologies. In section 5 we briefly report the cascade rates of fluid and magnetic energies. Section 6 contains a renormalization group scheme which provides a self-consistent procedure to obtain the effective mean magnetic field. Section 7 contains conclusions. | In this paper we have reviewed some of the phenomenological, observational, numerical, and analytic work in the statistical theory of MHD turbulence, with emphasis on our studies. We find that the solar wind observations and the numerical results are inconsistent with the predictions of the existing MHD turbulence phenomenologies. The energy spectrum and the cascade rates appear to be closer the predictions of Kolmogorov-like phenomenology even when the mean magnetic field or the magnetic field of the largest eddies is large compared to the inertial range velocity and magnetic field fluctuations (a region where Kolmogorov-like theory is not expected to hold). We have attempted to obtain Kolmogorov-like energy spectrum in MHD turbulence in presence of arbitrary $B_0$ by postulating that the effective $% B_0$ is scale dependent, unlike what has been taken by Kraichnan and Dobrowolny et al. We have constructed a renormalization group scheme and shown that the self consistent $B_0(k)\propto k^{-1/3}$ and $E(k)\propto k^{-5/3}$. This analysis has been worked out when $E^{+}=E^{-}$ and $r_A=1.$ The generalization to arbitrary parameters, or at least to the limiting cases, is also planned. We will be able to the get the Kolmogorov's constants for MHD\ turbulence analytically using this procedure; these constants will be useful for the large-eddy-simulations (LES) of MHD turbulence. We have also carried out a preliminary study of the cascade rates of fluid and magnetic energies. We find that there is a net transfer of fluid energy to magnetic energy. We are investigating the parameter regimes in which these cascade rates are large enough (more than dissipation rates) so that the total magnetic energy increases with time, what is found in dynamo theory. Lastly, we would like to mention that the MHD turbulence theories are very important for modelling various astrophysical phenomena and plasma processes. We have estimated turbulent heating, nonclassical viscosity and resistivity of the solar wind using the MHD turbulence phenomenologies \cite {Verma:sw}. In this light, search for a satisfactory theory of MHD turbulence appears quite important. We thank all our collaborators, D. A. Roberts, M. L. Goldstein, J. K. Bhattacharjee, and V. Eswaran. MKV\ also thanks V. Subrahmanyam, M. Barma, V. Ravishankar, and D. Sa for numerous useful discussions. | 98 | 3 | chao-dyn9803022_arXiv.txt |
9803 | nucl-th9803012_arXiv.txt | s{We discuss models to calculate one-- and two--neutron capture reactions on light nuclei. These are applied to calculate the reaction rates of $^{15}$N(n,$\gamma$)$^{16}$N, $^{16}$N(n,$\gamma$)$^{17}$N and $^4$He(2n,$\gamma$)$^6$He. The possible astrophysical importance is discussed.} | Neutron capture reactions on light nuclei play a role in various astrophysical scenarios. In the framework of Inhomogeneous Big Bang Models a high neutron flux can bridge the mass 5 and mass 8 gaps. Subsequent neutron capture reactions may trigger a primordial r--process \cite{rau94}. Another site for neutron capture reactions is the high--entropy bubble formed during a type II supernova \cite{woo92,mey92,woo94}. Due to the photodisintegration at very high temperatures an $\alpha$--rich environment is created. When the temperature has dropped heavier elements can be built up mainly by $\alpha$-- and neutron--capture reactions. Again a critical question is how the mass 5 and mass 8 gaps can be bridged. In Section 2 we describe our model to calculate one--neutron capture reactions. This model is applied to the reactions $^{15}$N(n,$\gamma$)$^{16}$N and $^{16}$N(n,$\gamma$)$^{17}$N. In Section 3 we will describe the theory of calculating a two--neutron capture reaction in a three--body model. We calculate the reaction rate of $^4$He(2n,$\gamma$)$^6$He and compare the result with other works. In the final chapter we discuss and summarize our results. | We calculated reaction rates for one--neutron capture on $^{15}$N and $^{16}$N and two--neutron capture on $^4$He. The rate for $^{15}$N(n,$\gamma$)$^{16}$N is in good agreement with both previous calculations and experimental data. In general the neutron capture rates on stable targets are quite well known. The reaction $^{16}$N(n,$\gamma$)$^{17}$N is an example of a capture reaction on an unstable target. We find a considerable enhancement to a previous calculation. This enhancement can be also be observed at various other capture reactions on unstable targets in this mass range \cite{her98}. This fact could influence the reaction path both in the nucleosynthesis of Inhomogeneous Big Bang Models and in the alpha--rich freeze--out of type II supernovae. For the reaction $^4$He(2n,$\gamma$)$^6$He we also find a strong enhancement of the rate compared to the calculation of Fowler \cite{fowler75}. But from the calculation of the inverse photodisintegration rate by Danilin {\it et al}~\cite{danilin93} we deduce an even much higher rate. This is probably due to the fact that this rate also includes the simultaneous capture of two neutrons. But even with this enhanced rate the path via $^4$He(2n,$\gamma$)$^6$He is dominated by the reaction $^4$He($\alpha$n,$\gamma$)$^9$Be at conditions typical for the alpha--rich freeze--out. This is due to the effective destruction of $^6$He through photodisintegration. | 98 | 3 | nucl-th9803012_arXiv.txt |
9803 | astro-ph9803183_arXiv.txt | According to a currently popular paradigm, nuclear activity in quasars is sustained {\it via } accretion of material onto super-massive black holes located at the quasar nuclei. A useful tracer of the gravitational field in the vicinity of such central black holes is available in the form of extremely dense gas clouds within the broad emission-line region (BLR) on the scale of $\sim 1~$parsec. Likewise, the radio sizes of the lobe-dominated radio sources are believed to provide a useful statistical indicator of their ages. Using two homogeneously observed (and processed) sets of lobe-dominated radio-loud quasars, taken from literature, we show that a positive correlation exists between the radio sizes of the quasars and the widths of their broad $H\beta$ emission lines, and this correlation is found to be significantly stronger than the other well known correlations involving radio size. This statistical correlation is shown to be consistent with the largest (and, hence, very possibly the oldest) radio sources harboring typically an order-of-magnitude more massive central engines, as compared to the physically smaller and, hence, probably much younger radio sources. This inference is basically in accord with the "accreting central engine" picture for the radio-loud quasars. | Being too compact to be resolved with even the most advanced optical telescopes, the structure and kinematics of the broad emission line region (BLR), a prime feature of quasars, continues to be a major enigma in the AGN research (see, e.g., Brotherton, 1996; Marziani et al., 1996; Corbin, 1992). Nonetheless, it is a key ingredient to the theoretical models that seek to explain various observable properties of quasars, e.g., their intense $\gamma$-ray emission (e.g., Dermer \& Schlickeiser, 1993; Ghisellini \& Madau, 1996). Deciphering the BLR geometry has therefore been a key objective of many observational programmes. One strategy for this is the so called `reverberation mapping' (e.g., Peterson, 1993; Koratkar \& Gaskell, 1991), from which a positive dependence of the BLR size on the bolometric luminosity has been inferred: $r~\sim 0.06 L_{46}^{0.5}~pc$ (Netzer \& Laor 1993), where $L_{46}$ is the luminosity expressed in the units of $10^{46} erg.s^{-1}$. \begin{table*} \label{ The sample of lobe dominated radio-loud quasars (48 quasars)} \begin{tabular}{lrrcccrrc} \multicolumn {9}{c}{\bf Table 1: Sample of lobe-dominated quasars (42 quasars)}\\ \multicolumn {9}{c}{}\\ \hline \hline \multicolumn {1}{c}{QSO}&\multicolumn {1} {c}{z} & \multicolumn {1}{c} {LAS}& \multicolumn {1}{c}{log($f_c$)}& \multicolumn {1}{c} {$\rm log(R_v)$}&\multicolumn {1}{c} {$\rm M_{v}$}&\multicolumn {2}{c}{W($\rm H\beta$) (km s$^{-1}$)}&\multicolumn {1}{c}{ $\xi$}\\ \cline {7-8} &&\multicolumn {1}{c} {(arcsec)}&&&&B(96)&JB(91)&\\ \hline 0003+158 &0.450&36.0(1)&$-$0.38(a)&1.95&-26.0&4760&...&...\\ 0042+101 &0.583&59.0(3)&$-$0.50(b)&1.79&-24.8&...&17774&...\\ 0044+030 &0.624&18.6(5)&$-$0.42(a)&0.83&-27.2&5100&...&...\\ 0110+297 &0.363&76.2(2)&$-$0.80(b)&1.59&-24.8&...&8702&...\\ 0115+027 &0.670&13.1(2)&$-$0.50(b)& 2.25&-26.5&5000&7591&0.66\\ 0118+034 &0.765&45.0(4)&$-$1.20(b)&2.17&-25.1&...&22403&...\\ 0133+207 &0.425&68.0(2)&$-$1.00(b)&2.51&-24.0&...&17403&...\\ 0134+329 &0.367&1.3(3)&$-$1.14(a)&2.50&-25.7&3800&5863&0.65\\ 0405$-$123&0.575&31.7(6)&$-$0.23(a)&2.36&-28.4&4800&...&...\\ 0414$-$060&0.781&36.4(4)&$-$0.50(a)&1.68&-27.8&8200&16602&0.49\\ 0518+165 &0.759&0.8(9)&$-$0.90(b)&3.94&-24.1&...&4876&...\\ 0538+498 &0.545&0.2(7)&$-$1.40(b)&3.43&-24.2&...&3456&...\\ 0710+118 &0.768&48.0(2)&$-$1.85(a)&1.09&-27.1&20000&17774&1.13\\ 0800+608 &0.689&25.0(10)&$-$0.60(b)&2.65&-24.2&...&6912&...\\ 0837$-$120&0.200&169.0(2)&$-$0.70(a)&1.80&-24.7&6060&...&...\\ 0838+133 &0.680&10.0(7)&$-$0.50(b)&3.16&-25.4&3000&4197&0.71\\ 0903+169 &0.410&50.0(8)&$-$1.40(c)&1.70&-24.5&4400&...&...\\ 0952+097 &0.298&12.5(2)&$<$$-$0.50(c)&$<$1.90&-24.1&3800&...&...\\ 0955+326 &0.530 &1.0(3)&$-$0.07(a)&2.23&-27.1&1380&...&...\\ 1004+130 &0.240&115.0(2)&$-$1.68(a)&0.43&-25.7&6300&9998&0.63\\ 1007+417 &0.613&32.0(2)&$-$0.39(a)&2.17&-27.0&3560&6912&0.52\\ 1048$-$090&0.334&83.0(2)&$-$1.31(a)&1.66&-24.9&5620&...&...\\ 1100+772 &0.311&30.0(1)&$-$0.91(a)&1.63&-25.8&6160&7961&0.77\\ 1103$-$006&0.423&21.0(2)&$-$0.14(a)&2.20&-25.8&6560&...&...\\ 1111+408&0.730&13.2(2)&$-$1.80(b)&1.89&-25.3&...&9134&...\\ 1137+660&0.652&44.2(2)&$-$1.07(a)&1.83&-27.1&6060&8702&0.69\\ 1223+252&0.268&67.0(2)&$-$1.59(b)&0.99&-23.8&...&9319 &...\\ 1250+568&0.321&1.5(2)&$-$1.62(a)&2.01&-23.6&4560&6295&0.72\\ 1305+069&0.599&46.5(4)&$<$$-$1.20(a)&$<$1.52&-26.1& 6440&...&...\\ 1351+267&0.310&190.0(2)&$-$0.27(a)&1.72&-24.3&8600&...&...\\ 1425+267&0.366&230.0(1)&$-$0.40(a)&1.03&-26.2&9410&...&...\\ 1458+718&0.905&2.1(2)&$-$1.00(d)&2.74&-27.3&3000 &...&...\\ 1512+370&0.371&54.0(1)&$-$0.71(a)&1.88&-25.6&6810&...&...\\ 1545+210&0.264&70.0(1)&$-$1.32(a)&1.74&-24.4&7030&...&...\\ 1618+177&0.555&48.0(3)&$-$0.67(a)&1.98&-26.5&7000&...&...\\ 1622+238&0.927&21.7(2)&$-$1.72(d)&1.63&-26.7&7100&...&...\\ 1704+608&0.371&55.0(1)&$-$1.91(a)&0.73&-26.6&6560&...&...\\ 1742+617&0.523&40.0(11)& $-$2.00(b)&1.61&-24.0&...&9751&...\\ 1828+487&0.691&14.0(11)&$-$0.50(b)&4.05&-26.2&...&9998&...\\ 2135$-$147&0.201&150.0(2)&$-$1.13(a)&1.97&-24.9&7300&11479&0.63\\ 2251+113&0.323&9.8(2)& $-$1.52(a)&1.00&-25.8&4160&8702&0.48\\ 2308+098&0.432&108.0(1)&$-$0.78(a)&1.41&-26.3&7970&...&...\\ \hline \multicolumn {9}{l}{{\bf References for LAS:} 1 : Kellermann et al. (1994), 2 : Nilsson et al. (1993), 3 : Singal (1988), }\\ \multicolumn {9}{l}{4 : Kapahi (1995), 5 : Price et al. (1993), 6 : Morganti et al.(1993), 7 : Bogers et al. (1994),}\\ \multicolumn {9}{l}{ 8 : Bridle et al (1994),9 : Akujor et al. (1993), 10 : Jackson et al. (1990), 11 : Reid et al. (1995)}\\ \multicolumn {9}{l}{{\bf References for $f_c$:} a : Wills \& Browne (1986), b : Jackson \& Browne (1991a), }\\ \multicolumn {9}{l}{c : Brotherton (1996), d : Wills et al. (1992), e : Kellermann et al. (1994)}\\ \end{tabular} \end{table*} An early indication about the BLR geometry came from an empirical study of a heterogeneously selected sample of radio-loud quasars with the characteristic core-lobe type radio structure. The study revealed a statistically significant anti-correlation between the {\it prominence} of the radio core relative to the lobe, and the FWHM of the $H\beta$ broad emission line ($W$). It was thus inferred that the BLR clouds are predominantly confined to a rotating disk-shaped region surrounding the quasar nucleus and oriented roughly perpendicular to the jet axis (Wills \& Browne, 1986; hereafter WB86). The significance of this correlation was found to be considerably lower in a subsequent study, however (Jackson \& Browne, 1991b). In a recent work, it has been proposed that the absolute {\it visual} magnitude, $M_v$, of the quasar (plus its host galaxy) provides a more reliable measure of the intrinsic power of the central engine, and therefore the beamed radio core flux normalized by $M_v$ is a better indicator of the orientation of the jet relative to the line-of-sight (Wills \& Brotherton, 1995; Brotherton, 1996). Adopting this new parameter for the core-prominence and employing a larger sample of quasars , these authors have found a conspicuous anti-correlation between the core-prominence and the $W(H\beta)$, thereby supporting the disk-like BLR geometry inferred earlier by WB86 (at least for the $H\beta$ emitting clouds). It is now widely believed that the large widths of the BLR emission lines are a manifestation of the deep gravitational potential at the centers of quasars, possibly due to super-massive black-holes. The accretion of the material postulated for sustaining the quasar luminosity is expected to steadily increase the mass of the central engine during the lifespan of the nuclear activity. If this conjecture is basically right, the question arises: {\sl do we see in the data any evidence for the postulated growth of the central mass} (e.g, in the dynamics of the BLR clouds) ? In the present study we shall examine this issue by employing published measurements of radio-loud, steep-spectrum quasars. While the needed reliable estimator of the age is generally not available for individual sources, the overall radio size can serve as a useful statistical measure of age (since most radio sources are believed to steadily grow in size with time; cf. Sect. 3). Thus, we wish to examine here the nature of any relationship between the observed linear sizes of steep-spectrum radio quasars and the widths of their broad H$\beta$ emission lines. | 98 | 3 | astro-ph9803183_arXiv.txt |
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9803 | astro-ph9803057_arXiv.txt | The method of obtaining confidence intervals on a subset of the total number of parameters ($p$) of a model used for fitting X-ray spectra is to perturb the best-fitting model until, for each parameter, a range is found for which the change in the fit statistic is equal to some critical value. This critical value corresponds to the desired confidence level and is obtained from the $\chi^{2}$ distribution for $q$ degrees of freedom, where $q$ is the number of interesting parameters. With the advent of better energy-resolution detectors, such as those onboard {\it ASCA} ({\it Advanced Satellite for Cosmology and Astrophysics}) it has become more common to fit complex models with narrow features, comparable to the instrumental energy resolution. To investigate whether this leads to significant non-Gaussian deviations between data and model, we use simulations based on \asca data and we show that the method is still valid in such cases. We also investigate the weak-source limit as well as the case of obtaining upper limits on equivalents widths of weak emission lines and find that for all practical purposes the method gives the correct confidence ranges. However, upper limits on emission-line equivalent widths may be over-estimated in the extreme Poisson limit. | \label{intro} The procedure for generating confidence intervals for the parameters of models used to fit X-ray spectra with the \chisq minimization technique is well established (e.g. Lampton, Margon and Bower 1976; Avni 1976). One uses the fact that if $\chi ^{2} _{\rm true}$ is the value of the statistic from a given experiment, calculated for the `true model' (known only to nature) and $\chi_{\rm min}^{2}$ is the value of the statistic obtained from the best-fitting model, then from many repetitions of the same experiment $\Delta \chi^{2} = \chi^{2}_{\rm true} - \chi^{2} _{\rm min}$ has a probably distribution like $\chi ^{2}$ with $p$ degrees of freedom ($\equiv \chi^{2}_{p}$) where $p$ is the total number of free parameters. Confidence intervals generated for the $p$ parameters from the $\chi^{2}_{p}$ distribution are then {\it joint} confidence intervals for all $p$ parameters. If we are interested only in a subset $q$ of the $p$ parameters then the confidence intervals are generated from the $\chi^{2}_{q}$ distribution. The crucial difference between $p$-parameter and, say, one-parameter errors is as follows. In the former case the $P\%$ confidence intervals from $\chi^{2}_{p}$ will {\it simultaneously} enclose the true values of all parameters in $P\%$ of all experiments. In the latter case, the $P\%$ intervals from $\chi^{2}_{1}$ will enclose the true values of any of the parameters in $P\%$ of all experiments, {\it but it will be a different $P\%$ subset of experiments for each parameter}. In a particular case, the number of `interesting' parameters (i.e. $q$) is determined by the scientific problem being posed. In principle one generates the $\Delta \chi ^{2}$ space by stepping through a $q$-dimensional grid of parameter values. The use of $\chi ^{2}$ in model-fitting requires that there are a sufficient number of photons per energy bin for the statistical variations to be Gaussian. However, with the advent of better energy resolution X-ray detectors it is increasingly becoming the case that the Gaussian approximation cannot be made unless the spectrum is binned, sacrificing valuable information. In such cases the statistical variation in counts per bin is Poisson and one must use the maximum likelihood ratio (hereafter, `$C$ statistic') to optimize the model parameters (see Cash 1979). Cash (1979) demonstrated that the $C$ statistic can be used to generate confidence intervals in an analogous manner to $\chi ^{2}$ since $\Delta C$ has a probability distribution like $\chi ^{2}$ except for terms of order $\alpha/ \sqrt{n}$ where $\alpha$ depends on the model and $n$ is the number of photons carrying information about the parameter in question. Hereafter, we shall use $C$ for the sake of generality. Lampton \etal (1976) warned that (at that time), there was no general proof that the technique for projection of the subset of $q$ parameters did not depend explicitly on model linearity. Then, Avni (1976), using some non-linear models in simulations of {\it Uhuru} data, showed that the $\chi^{2}_{q}$ region worked in these particular cases, but there was still no general proof. Such a proof was presented by Cash (1976), showing that the $\chi^{2}_{q}$ region worked for any data set, {\it provided that the deviations are Gaussian}. We must remember that X-ray detectors now have much better energy resolution and sensitivity than they did then and accordingly the models have become much more complex. Both Lampton \etal (1976) and Avni (1976) used {\it Uhuru} spectral responses with seven energy bins between 1 and 7 keV, and simple power-law or plasma models. It is not clear whether, for example, a model including a narrow emission line whose intrinsic width is comparable to the energy resolution of the detector, would introduce non-Gaussian deviations between model and data. The purpose of this paper is to check whether the method for parameter estimation currently in use gives the correct confidence intervals even with the new generation of improved energy-resolution instrumentation. We are particularly interested in models with features such as emission lines and absorption edges (which occur frequently in a wide range of X-ray sources), whose widths in energy space are comparable to the instrumental energy resolution. We also wish to check the Poisson regime, when the source count rate is low, and when upper limits must be obtained on the equivalent width of weak or non-detected emission-line features. The structure of the paper is as follows: \S \ref{simulations} describes basic simulations and models used in the investigation; \S \ref{results} describes the basic results; \S \ref{lineew} demonstrates the equivalence of emission-line equivalent width and intensity confidence regions; \S \ref{weaksource} describes results for the extreme Poisson limit and \S \ref{weakline} describes results pertaining to measuring upper limits on weak emission-line features. Our conclusions are stated in \S \ref{conclusions}. | \label{conclusions} We have verified that for all practical purposes, the method of generating confidence intervals for a subset, $q$, of $p$ model parameters, using the $\chi^{2}_{q}$ distribution can still be used even if the model has components which are narrow compared to the instrumental energy resolution (such as emission lines and absorption edges). We have also investigated the weak-source and weak emission-line limits and find the method to work, except for the extreme Poisson limit when there may be one or less total counts per energy bin. In this case, equivalent widths of emission lines may be somewhat over-estimated. It must be remembered, however, that the $\chi^{2}_{q}$ confidence ranges can say nothing of the {\it simultaneous} confidence ranges of the other $p-q$ parameters. For example, suppose one observes an active galaxy or X-ray binary and measures the magnitude of an X-ray reflection continuum component (due to Compton-thick scattering) and the equivalent width of an iron-K line and quotes, as is common practice, 90\% confidence errors for one interesting parameter ($\Delta \chi^{2} = 2.7$). Now, the relation between the iron-K line equivalent width and the strength of the reflection continuum can be predicted from a theoretical model, so one can determine whether the measured values are consistent with such a model, within the errors. However, one-parameter errors (as used , for example, in Zdziarski \etal 1996) are inappropriate since these confidence ranges are not simultaneous. One must use two-parameter errors in such a case. \vspace{2cm} Much of this work was done during an extended stay at the Institute of Space and Astronautical Science (ISAS), Japan, during the summer of 1993 and two weeks in November 1996. The author thanks everyone in the X-ray astronomy group at ISAS for their great hospitality. The author would also like to thank Peter Serlemitsos, Richard Mushotzky, Andy Fabian, and Andy Ptak for useful discussions, and Keith Arnaud for generally maintaining XSPEC and fixing bugs promptly. The author is very grateful to Dr. W. Cash for correcting some serious errors in an earlier version of the paper and is also indebted to an anonymous referee for making some extremely important points which led to a complete revision of the paper. \newpage | 98 | 3 | astro-ph9803057_arXiv.txt |
9803 | astro-ph9803327_arXiv.txt | We present {\it Faint Object Camera (FOC)} ultraviolet images of the central $14 \times 14 \arcsec$ of Messier 31 and Messier 32. The hot stellar population detected in the composite UV spectra of these nearby galaxies is partially resolved into individual stars, and their individual colors and apparent magnitudes are measured. We detect 433 stars in M~31 and 138 stars in M~32, down to detection limits of $m_{F275W}$~=~25.5~mag and $m_{F175W}$~=~24.5~mag. We investigate the luminosity functions of the sources, their spatial distribution, their color-magnitude diagrams, and their total integrated far-UV flux. Comparison to {\it IUE} and {\it HUT} spectro-photometry and {\it WFPC2} stellar photometry indicates consistency at the 0.3~mag level, with possible systematic offsets in the {\it FOC} photometry at a level less than this. Further calibrations or observations with the {\it Space Telescope Imaging Spectrograph (STIS)} will be necessary to resolve the discrepancies. Our interpretation rests on the assumption that the published {\it FOC} on-orbit calibration is correct. Although M~32 has a weaker UV upturn than M~31, the luminosity functions and color-magnitude diagrams of M~31 and M~32 are surprisingly similar, and are {\it inconsistent} with a majority contribution from any of the following: PAGB stars more massive than 0.56~$M_\odot$ (with or without associated planetary nebulae), main sequence stars, or blue stragglers. Both the luminosity functions and color-magnitude diagrams are {\it consistent} with a dominant population of stars that have evolved from the extreme horizontal branch (EHB) along tracks with masses between 0.47 and 0.53 $M_\odot$. These stars are well below the detection limits of our images while on the zero-age EHB, but become detectable while in the more luminous (but shorter) AGB-Manqu$\acute{\rm e}$ and post-early asymptotic giant branch (PEAGB) phases. The {\it FOC} observations require that only a very small fraction of the main sequence population (2\% in M~31 and 0.5\% in M~32) in these two galaxies evolve though the EHB and post-EHB phases, with the remainder evolving through bright PAGB evolution that is so rapid that few if any stars are expected in the small field of view covered by the {\it FOC}. A model with a flat EHB star mass distribution reproduces the {\it HUT} and {\it IUE} spectra of these two galaxies reasonably well, although there is some indication that an additional population of very hot (T$_{\rm eff} > 25000$ K) EHB stars may be needed to reproduce the {\it HUT} spectrum of M~31 near the Lyman limit, and to bring integrated far-UV fluxes of M~31 and M~32 into agreement with {\it IUE}. In addition to the post-EHB population detected in the {\it FOC}, we find a minority population ($\sim$ 10\%) of brighter stars that populate a region of the CMD that cannot be explained by canonical post-HB evolutionary tracks. The nature of these stars remains open to interpretation. The spatial distributions of the resolved UV-bright stars in both galaxies are more centrally concentrated than the underlying diffuse emission, implying that stellar populations of different age and/or metallicity might be responsible for each component. | The spectra of elliptical galaxies and spiral galaxy bulges exhibit a strong upturn shortward of 2700~\.{A}, dubbed the ``UV upturn.'' At the time of its discovery, the existence of a hot stellar component went against the traditional picture of early-type galaxies. The canonical view of ellipticals held that these galaxies contained a cool, passively evolving population of old stars. The pioneering UV observations of ellipticals -- with the {\it Orbiting Astronomical Observatory (OAO)} (Code \& Welch 1979\markcite{CW79}) and the {\it International Ultraviolet Explorer (IUE)} (Bertola, Capaccioli, \& Oke 1982\markcite{BCO82}) -- could only sample the Rayleigh-Jeans tail of the hot UV flux, with poor signal-to-noise and resolution. Early explanations for the source of the UV upturn covered a wide range of candidates, including massive young stars, hot horizontal branch stars, planetary nebula nuclei, and several binary scenarios (see Greggio \& Renzini 1990\markcite{GR90} for a complete review). The presence of young stars would imply ongoing star formation in early-type galaxies, while the evolved candidates suggested that old stellar populations could be efficient UV emitters. Characterized by the $m_{1550}-V$ color, the UV upturn shows surprisingly strong variation (ranging from 2.05--4.50 mag) in nearby quiescent early-type galaxies (Bertola et al.\ 1982\markcite{BCO82}; Burstein et al.\ 1988\markcite{B88}), even though the spectra of ellipticals at longer wavelengths are qualitatively very similar. A large sample of UV measurements demonstrated that the UV upturn is positively correlated with the strength of Mg$_2$ line absorption in the $V$ band, in the sense that the $m_{1550}-V$ color is bluer at higher line strengths, opposite to the behavior of optical color indices (Burstein et al.\ 1988\markcite{B88}). Opposing theories have been devised to explain this correlation. For example, Lee (1994\markcite{L94}) and Park \& Lee (1997\markcite{PL97}) suggest that the UV flux originates in the low metallicity tail of an evolved stellar population with a wide metallicity distribution. Their reasoning is that the more massive ellipticals formed earlier, and that it is actually the {\it mean} metallicity that is higher (and driving the optical indices) in the older and bluer galaxies. In contrast, several groups (Brocato et al.\ 1990\markcite{BMMT90}; Bressan, Chiosi, \& Fagotto 1994\markcite{BCF94}; Greggio \& Renzini 1990\markcite{GR90}; Horch, Demarque, \& Pinsonneault 1992\markcite{HDP92}; Brown et al.\ 1997\markcite{B97}; Yi, Demarque, \& Oemler 1997\markcite{YDO97}) argue that metal-rich horizontal branch stars and their progeny are responsible for the UV flux. Under the metal-rich hypothesis, the canonical trend in horizontal branch morphology (i.e., redder HBs with higher metallicity) is reversed at high metallicity, due to increased helium abundance and possibly a higher mass loss rate on the red giant branch, resulting in the production of hot UV-efficient stars on the extreme horizontal branch (EHB). These hypotheses lead to different ages for the stellar populations in these galaxies. Ages exceeding those of Galactic globular clusters are required under the low-metallicity Park \& Lee (1997\markcite{PL97}) hypothesis, while ages as low as 8 Gyr are allowed in the Bressan et al.\ (1994\markcite{BCF94}) model. In these two scenarios, the EHB stars are drawn from either tail of the metallicity distribution. However, it is also possible, indeed perhaps more likely, that the EHB stars arise from progenitors near the peak of the metallicity distribution (cf.\ Dorman, O'Connell, \& Rood 1995\markcite{DOR95}), but represent a relatively rare occurrence. The correlation of $m_{1550}-V$ with the global metallicity of the galaxy might indicate that this rare path of stellar evolution becomes less so at high metallicity and helium abundance. The {\it Hopkins Ultraviolet Telescope (HUT)}, designed for medium resolution ($\approx$~3~\.{A}) spectroscopic observations of faint extended UV sources down to the Lyman limit at 912~\.{A}, offered a new perspective on these populations. With observations of two galaxies on the Astro-1 mission (M~31 and NGC~1399), Ferguson et al.\ (1991\markcite{F91}) demonstrated that young, massive stars cannot be a significant contributor to the UV upturn. There is a lack of strong \ion{C}{4} absorption expected from such stars, and the continuum flux decreases from 1050~\.{A} down to the Lyman limit. Such a decrease implies that the UV flux is dominated by stars with temperatures $\leq 25000$~K and is incompatible with emission by a population of young stars having a normal initial mass function. Ferguson et al.\ (1991\markcite{F91}) also suggested that a bimodal distribution on the horizontal branch was needed to reproduce the shape of spectra from the near-UV to the far-UV, otherwise the spectra would be flatter than observed. Six more galaxies were observed on the Astro-2 mission (M~49, M~60, M~87, M~89, NGC~3115, and NGC~3379), and with these data Brown et al.\ (1997\markcite{B97}) demonstrated that a two-component population of high-metallicity post-HB stars could reproduce the UV light seen in nearby ellipticals. In this model, most ($> 80$\%) of the UV-producing stars were undergoing post-asymptotic-giant-branch (PAGB) evolution, and the remainder were evolving along AGB-Manqu$\acute{\rm e}$ paths from the extreme horizontal branch. Although in the minority, these AGB-Manqu$\acute{\rm e}$ stars can produce the majority of the flux, because their lifetimes are orders of magnitude longer than those of the PAGB stars. Because even the brightest of these galaxies are faint and extended in the UV, studies of the UV upturn have focused mostly on the composite spectral energy distributions of ellipticals. However, the UV imaging capabilities of {\it HST} have now opened the possibility of studying the resolved UV population, at least in the nearest galaxies. Attempts to do this prior to the {\it HST} refurbishment were undertaken by King et al.\ (1992\markcite{K92}), for M~31, and by Bertola et al.\ (1995\markcite{BBB95}), for M~31 and M~32. King et al.\ (1992\markcite{K92}) obtained a pre-COSTAR {\it Faint Object Camera (FOC)} observation of a $44 \times 44 \arcsec$ field in the center of M~31, using the F175W filter and the F/48 relay. They found more than 100 sources that they identified as PAGB stars. Intermediate-mass ($M > 0.6~M_{\odot}$) PAGB stars are short-lived, and given a population size constrained by the fuel consumption theorem, the large number of detected stars implied that these were low-mass PAGB stars. King et al.\ (1992\markcite{K92}) estimated that these stars accounted for approximately 20\% of the UV light, with the rest unresolved, presumably coming from EHB stars and their AGB-Manqu$\acute{\rm e}$ descendants, which could account for this unresolved light without violating fuel consumption constraints. Bertola et al.\ (1995\markcite{BBB95}) used the {\it FOC} to image M~31, M~32, and NGC~205 with the combination of the F150W and F130LP filters on the F/48 relay. Although the optical luminosity enclosed by the {\it FOC} field was higher in M~32 than in M~31, they found far fewer UV sources in the M~32 field. Because M~32 has a weaker UV upturn and lower metallicity, the UV light was expected to originate in PAGB stars of higher mass (and shorter lifetimes) than those in M~31; so, the relative numbers of detected sources were in line with these expectations. However, the luminosity functions in M~31 and M~32 appeared similar, in contrast to expectations when comparing a population of less massive PAGB stars to a population of more massive ones. Such a puzzle may be partly explained if the PAGB stars are enshrouded in dust during the early part of their evolution away from the AGB. Both previous {\it FOC} studies faced daunting challenges in untangling the uncertainties in the {\it FOC} sensitivity calibration and the red leak of the filters. King et al.\ (1992\markcite{K92}) adopted the best in-flight calibration at the time, and used ground calibrations of the filter and photocathode response to assess the effects of red leak. They also determined that the pre-COSTAR PSF required a huge aperture correction of 2.6~mag. Bertola et al.\ (1995\markcite{BBB95}) derived an independent calibration based upon {\it IUE} observations of NGC~205, M~31, and M~32. In the Bertola et al.\ calibration, the nominal {\it FOC} F150W+F130LP efficiency curve required multiplication by factors of 0.21--0.92 (varying with wavelength) in order to produce agreement between the {\it FOC} and {\it IUE}. Checking their revised calibration against common stars in the King et al.\ (1992\markcite{K92}) F175W images, Bertola et al.\ determined that the F175W efficiency curve also required revision, such that the peak in the efficiency curve was at 28\% of its nominal value, but with increased red leak from longer wavelengths. We demonstrate in \S\ref{secf48} that the pre-COSTAR {\it FOC} data calibration was seriously in error. We have used the refurbished, recalibrated {\it FOC} to follow up these earlier studies with deeper UV images of the M~31 and M~32 cores, in order to further characterize the evolved stellar populations in these galaxies. Our observations use the F175W and F275W filters to determine color-luminosity relationships in these galaxies, and to compare them to the predictions of stellar evolutionary theories. | \label{secdis} Since the discovery of the UV upturn phenomenon, a great deal of research has tried to characterize the stellar population producing the UV light in spiral bulges and ellipticals. The range of candidates has included young massive stars, hot horizontal branch stars, planetary nebula nuclei (PAGB stars), and several binary scenarios (see Greggio \& Renzini 1990\markcite{GR90} for a complete review). Earlier work relied primarily upon {\it IUE} and {\it HUT} spectra of composite stellar populations, and pre-COSTAR {\it FOC} observations (King et al.\ 1992\markcite{K92}; Bertola et al.\ 1995\markcite{BBB95}) of the brightest stars. Our current understanding of the populations in ellipticals and spiral bulges predicts that a significant population of hot HB and post-HB stars should be evolving from the extreme horizontal branch and following either AGB-Manqu$\acute{\rm e}$ or post-early-AGB evolution (see Brown et al.\ 1997\markcite{B97} and references therein). The prediction is mainly based upon fuel consumption constraints that rule out PAGB stars as the sole UV producers; the PAGB stars are just too short-lived to efficiently produce the required UV light. Our {\it FOC} observations confirm the existence of hot post-EHB stars in the centers of M~31 and M~32, and also confirm that PAGB stars are not the predominant component of the UV-bright population. The existence of these EHB stars also implies that the horizontal branches in both M~31 and M~32 are at least somewhat extended, because ``red clump'' horizontal branches are unlikely to produce the stars seen in our color-magnitude diagrams. Future observations, given the improved UV capabilities of HST with the installation of {\it STIS}, should be able to reach the HB and thus characterize the entire evolved population more directly. We note that significant systematic uncertainties may exist in our {\it FOC} photometry, at the 0.3 mag level. Our results should be considered with these uncertainties in mind; however, the uncertainties are not large enough to allow a completely different interpretation of the detected stellar population. As discussed in \S\ref{secehbcmd}, we find a minority population ($\sim$ 10\%) of brighter stars that cannot be explained by canonical post-HB evolutionary tracks. Although PEAGB or very low-mass PAGB stars evolve through this region of the CMD, their evolution during this phase is too rapid to produce the 35 stars seen in M~31 while maintaining consistency with fuel consumption constraints and observed integrated spectra of M~31. The nature of these stars remains unexplained as of this writing. The fraction of light in the resolved UV population in the center of M~31 and in the center of M~32 is consistent with expectations from the {\it HUT} and {\it IUE} spectra of these galaxies. The far-UV light in M~31 can be explained by a main sequence population where 98\% of the stars channel through intermediate-mass PAGB evolution and 2\% through EHB evolution. The far-UV light in M~32 can be explained by a population where 99.5\% of the stars channel through intermediate-mass PAGB evolution and 0.5\% through EHB evolution. Our simple model populations can account for most of the far-UV light, although there is room in both galaxies for a contribution from very low-mass hot EHB stars, which would be below our detection limits even in the AGB-Manqu$\acute{\rm e}$ phase. We found that the stellar populations in M~31 and M~32 do not appear remarkably different in our {\it FOC} UV images. Fewer stars appear to be entering the UV-efficient EHB paths in M~32, and this explains both the smaller absolute number of detected stars and the smaller fraction of resolved UV light. The similarities of the luminosity functions of M~31 and M~32 were unexpected, given the dramatically different $m_{1550}-V$ colors of the two galaxies. This finding suggests that, while the fractional mass in EHB stars is probably sensitive to the properties of the overall stellar population, (e.g., metallicity, age, or helium abundance), the mass distribution on the EHB may not be as sensitive to these parameters. Certainly, the differences can not be investigated without deeper images that can resolve more of the evolved population. {\it STIS} observations planned for late 1998 will reach the HB in M~32 and provide direct evidence for the horizontal branch distribution. | 98 | 3 | astro-ph9803327_arXiv.txt |
9803 | astro-ph9803111_arXiv.txt | We have obtained high-resolution echelle spectra of 18 solar-type stars that an earlier survey showed to have very high levels of \ion{Ca}{2} H and K emission. Most of these stars belong to close binary systems, but 5 remain as probable single stars or well-separated binaries that are younger than the Pleiades on the basis of their lithium abundances and \ha\ emission. Three of these probable single stars also lie more than 1 magnitude above the main sequence in a color-magnitude diagram, and appear to have ages of 10 to 15 Myr. Two of them, HD 202917 and HD 222259, also appear to have a kinematical association with the pre-main sequence multiple system HD 98800. | Two years ago we published a survey of \ion{Ca}{2} H and K emission strengths in more than 800 southern solar-type stars (\cite{h96} 1996; hereafter Paper I), determined from low-resolution ($R\approx2,000$) spectra obtained at CTIO. The purpose of the survey was to provide at least rough estimates of the ages of the individual stars, and to examine the distribution of emission strengths in a large and unbiased sample. The very existence of a simple or single relationship between chromospheric emission (CE) and age is debatable, but there is ample evidence that CE declines steadily with age (\cite{sk72} 1972; \cite{sdj} 1991). In other words, we are confident of a general decline of CE with age because of observations of stars in clusters, and also because CE is so intimately tied to stellar rotation and we know that rotation declines with age in solar-type stars. But we also know that the CE levels of individual stars vary due to rotational modulation of active regions, long-term cycles, and other phenomena, and we also know that not all stars reach the Zero-Age Main Sequence (ZAMS) with the same angular momentum, and thus we are not so sure that there is a unique CE-age relation that applies to all stars. The survey of Paper I included stars from a G-dwarf sample defined using the combination of two-dimensional spectral types (\cite{hk1} 1975; \cite{hk2} 1978; \cite{hk3} 1982; \cite{hk4} 1988) and Str\"omgren photometry (\cite{o88} 1988, \cite{o93} 1993), and it turned up two groups of stars that we have studied further. The first group we called ``Very Active'' because they exhibited CE levels well above any seen in the field stars of the earlier survey of \cite{vp80} (1980). The second group is called ``Very Inactive'' and consists of stars that appear to have activity levels well below the Sun's. The Very Inactive stars will be the subject of a future paper. Here we concentrate on the Very Active stars, and we are led to examine them in detail for several reasons. First, stars with very high levels of activity are that way because they rotate rapidly, and they rotate rapidly either because they are very young (and have not yet lost much of their initial angular momentum), or because they are in a close binary system (where the companion's tidal forces lead to synchronous rotation). Both types of systems are interesting, and both types offer laboratories for the study of the rotation-activity relation. A second reason for undertaking this detailed study is to find very young stars in the immediate solar neighborhood (i.e., within about 50 pc). Even if they were evenly mixed in the Galaxy, very young stars would be rare just because their ages ($\la100$ Myr) represent such a small fraction of the age of the Galactic disk. But such stars are {\it not} evenly mixed because they form in discrete regions and take time to be dispersed into the field. There may be a few stars near the Sun that are as young as, say, the Pleiades, but they are few indeed. However, even small numbers matter since they imply that many more such stars lie in the much vaster volume of the greater solar neighborhood (i.e., within $\sim100$ pc). Some of these very young field stars, such as the HD 98800 system, may be examples of the elusive post-T Tauri star class; see, e.g., \cite{s98} (1998). The moderate-resolution spectra obtained for Paper I were centered on the \ion{Ca}{2} H and K lines to determine the level of CE. We have obtained higher-resolution echelle spectra to confirm their high activity levels (at \ha), to test for youth (with Li), and to observe each star several times to search for radial velocity changes indicative of close companions. | We present high resolution echelle spectroscopy of 18 Very Active southern solar-type stars identified in the \ion{Ca}{2} H and K survey of \cite{h96} (1996). We find evidence from line doubling or radial velocity variations that 13 of these are members of short-period, close binary systems. Activity in these stars may thus be due to the presence of a close companion, rather than youth; however, it is entirely possible that some of the objects (HD 106506, 119022, 155555AB, and HD 222259B) are young close binaries. Just a few exposures of high resolution but very modest S/N seems to be a highly effective and efficient means of identifying active close binary systems. Based on our \ha\ observations, we confirm that the remaining five of the eighteen stars are also Very Active, and find no evidence that such activity is due to membership in a close binary system. Four of these stars also have significant Li abundances, comparable to or larger than Pleiades stars of similar color. We consider these five stars to be young candidates, i.e., stars whose chromospheric activity seems associated solely with youth. We note that two of the young candidates (HD 202917 and HD 222259) appear to have the same $U,V$ velocities as HD 98800, a rare example of a post-T Tauri star in the field, according to \cite{s98} (1998). While the formation site of these stars is unclear, the three stellar systems may be part of a small group of very young stars sharing a common history and origin. The $UVW$ velocities of HD 175897 are in excellent agreement with those of the Pleiades. Two other interesting objects are HD 106506 and HD 119022, which were noted above to be possible young objects excluded from our final best young candidate list due to their duplicity. Both of these objects demonstrate sharp features in the broader Na D lines and stronger features of other metals. We do not observe any radial velocity shifts of these sharp features. This would suggest an interstellar or circumstellar origin for the sharp D-line features. These two objects are among the most distant in our sample: the Hipparcos-based distances are ${\sim}125$ pc, but seems too near for an interstellar origin. There is thus the intriguing possibility that these features arise from circumstellar material, perhaps not unlike that surrounding ${\beta}$ Pictoris or early-type shell stars similarly inferred from the presence of sharp features in a broader stellar absorption line. Comparison of the spectral type and photometry for HD 119022 suggests a moderate-sized reddening of perhaps 0.1 to 0.15 mag in $E(B-V)$. Higher quality spectra and high resolution imaging of these objects would be of great interest. Finally, Li abundances or upper limits were derived for the sharp-lined inactive \ha\ standards employed in our program. We report detectable Li for the solar \teff\ star HD 76151 and for the late F-stars HD 38393 and 45067. There may be a Li abundance difference in the two similar wide components of HD 158614; better knowledge of their fundamental parameters from spatially resolved photometry and spectroscopy is needed. | 98 | 3 | astro-ph9803111_arXiv.txt |
9803 | astro-ph9803039_arXiv.txt | Multi-epoch VLBA observations of the maser in NGC 4258 have yielded a 4\% geometrical distance to the galaxy. The potential scientific payoffs of finding similar objects at large distances, in the Hubble flow, are considerable. In this contribution, I discuss the plausibility of detecting high-redshift water masers, and describe a search strategy that we have implemented to realize this objective. | VLBA observations of the water masers in NGC 4258 reveal a nearly edge-on, slightly warped, extremely thin disk in nearly perfect Keplerian rotation around a central binding mass of $3.5\times10^{7}$\,M$_{\odot}$ (Watson \& Wallin 1994; Greenhill {\it et al.} 1995; Miyoshi {\it et al.} 1995; Moran {\it et al.} 1995; Herrnstein, Greenhill, \& Moran 1996). The VLBA observations of the NGC 4258 maser provide insight into the structure and kinematics of the accretion disk, and NGC 4258 is an exceptional laboratory for the study of AGN accretion phenomenon and the connection between accretion disks and jet emission (Herrnstein {\it et al.} 1997; Herrnstein {\it et al.} 1998a). They can also be used to derive a precise geometric distance to NGC 4258. The upper panel of Figure~1 is a schematic representation of the best-fitting NGC~4258 disk model, as derived from the positions and line-of-sight (LOS) velocities of the masers. As the disk rotates, the 'systemic' masers along the near edge of the disk drift in position and LOS velocity by about 30\,$\mu$as\,yr$^{-1}$ and 9\,km\,s$^{-1}$, respectively, with respect to the apparently stationary 'high-velocity' masers in the plane of the sky (Herrnstein 1997a\&b). The expressions at the bottom of Figure~1 illustrate that {\it both} the LOS accelerations ($\dot{v}_{LOS}$) {\it and} the proper motions ($\dot{\theta}_{x}$) can be used to derive a purely geometric distance ($D$) to NGC 4258. Because the actual space velocities ($v_{rot}$) and angular radii ($\theta_{R}$) of the systemic masers cannot be measured directly, these acceleration and proper motion distances are model dependent. Fortunately, however, most of the model dependence resides in $\theta_{R}$, and the two distance estimates together provide a geometric distance estimate that is largely model independent. With this in mind, we have observed NGC 4258 with the VLBA at 3--4 month intervals for the last three years. The LOS accelerations and proper motions provided by the first five epochs of these data yield a geometric distance of $7.3\pm0.3$\,Mpc (Herrnstein 1997a; Herrnstein {\it et al.} 1998b). \begin{figure}[p] % \plotone{herrnf1.eps} \caption {The top panel is a schematic representation of the sub-parsec molecular disk in NGC 4258. The bottom panels show the expected acceleration and proper motion vectors superposed on actual VLBA data of the systemic masers taken in May of 1995.} \end{figure} | 98 | 3 | astro-ph9803039_arXiv.txt |
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9803 | astro-ph9803333_arXiv.txt | We present the results of adaptive-optics imaging of the $z=0.2923$ QSO PG\,1700+518 in the $J$ and $H$ bands. The extension to the north of the QSO is clearly seen to be a discrete companion with a well-defined tidal tail, rather than a feature associated with the host galaxy of PG\,1700+518 itself. On the other hand, an extension to the southwest of the QSO (seen best in deeper, but lower-resolution, optical images) does likely comprise tidal material from the host galaxy. The SED derived from images in $J$, $H$, and two non-standard optical bands indicates the presence of dust intermixed with the stellar component. We use our previously reported Keck spectrum of the companion, the SED found from the imaging data, and updated spectral-synthesis models to constrain the stellar populations in the companion and to redetermine the age of the starburst. While our best-fit age of 0.085 Gyr is nearly the same as our earlier determination, the fit of the new models is considerably better. This age is found to be remarkably robust with respect to different assumptions about the nature of the older stellar component and the effects of dust. | PG\,1700+518 ($z=0.2923$), one of the more luminous low-redshift QSOs, shows a bright, arc-like structure extending about 2\arcsec\ to the north of the QSO (Hutchings, Neff, \& Gower \markcite{hut92b}1992; Stickel \etal\ \markcite{sti95}1995). We recently presented a spectrum of this extension, showing that it is dominated in the optical by a stellar population $\sim10^8$ years old (Canalizo \& Stockton \markcite{can97}1997; hereinafter Paper 1). We argued that it is plausible that this starburst and the QSO activity were both triggered by a recent interaction, and that spectroscopic dating of starbursts in QSO hosts and strongly interacting companions could lead to the development of an empirical evolutionary sequence for QSOs. From the imaging data we had on PG\,1700+518 at that time, we were unsure whether the extension was a companion galaxy or a feature associated with the host galaxy of the QSO. Here we describe the results of adaptive-optics (AO) imaging and show that the object is indeed a discrete companion galaxy undergoing strong interaction with the QSO host. We combine these data with previous imaging in two optical bandpasses to give a spectral-energy distribution (SED) from 0.42 \micron\ to 1.28 \micron\ in the rest frame in order to place additional constraints on the stellar populations. We then use these constraints and fits of new spectral synthesis models to the Keck LRIS spectrum to obtain a more reliable age for the post-starburst component. | \subsection{The Morphology of the Companion Galaxy} The AO images of PG\,1700+518 are shown in Fig.\ 1, in original, PSF-subtracted, and {\it plucy}-restored versions. The extension to the north of the QSO, which has the appearance of an arc-like or ``boomerang'' shape in the best previous ground-based images (Hutchings \& Neff \markcite{hut92}1992; Stickel \etal\ \markcite{sti95}1995; Paper 1), retains that general appearance in our higher-resolution images, but it now also shows discrete condensations within this overall structure. It is also clearly a separate companion galaxy: there is a rather abrupt dropoff in surface brightness just south of the brightest condensation, whereas we would expect to see more continuity with the inner regions if this were a tail associated with the host galaxy. Such a connection should stand out in spite of our using the QSO to model the PSF, since we expect to be sensitive to any non-elliptically-symmetric features. The companion, though distinct, is apparently in the process of merging with the QSO host galaxy. The main features of the companion are consistent between the $J$ and $H$ images and must be real. The bright condensation $a$ is probably the nucleus of the companion. The apparent tidal tail, curving to the north and east, contains another condensation, $b$, which may be a bright star-forming region or even a dwarf galaxy forming from the tidal debris (\eg\ Duc \& Mirabel \markcite{duc94}1994; Hunsberger, Charlton, \& Zaritsky \markcite{hun96}1996). The nature of features at lower surface brightnesses is less certain because of the effect of noise on the deconvolution. There is clearly material to the east of $a$, looking like another condensation in the $J$ image, but more like an arc in the $H$ image. The peak to the south of $a$ on the $J$ image is apparently an artifact of inadequacies in the PSF model, but there is some evidence in these images for bridge-like material between the companion and the QSO. If most of the luminous material north of the QSO is associated with the companion, is there any evidence for tidal debris from the QSO host? Stickel \etal\ \markcite{sti95}1995 noted a possible faint extension to the SW of the QSO, which can also be seen in Fig.\ 2 of Paper 1. Here we show this feature in two optical bandpasses in Fig.\ 2, where we have slightly oversubtracted the wings of the PSF profile to show the SW extension more clearly. We suggest that this is likely a counter-tidal feature from the QSO host. The inner, higher-surface-brightness contours of the host galaxy appear to be aligned nearly E--W (Paper 1). As we were completing this {\it Letter,} we were informed that Hines \etal\ \markcite{hin98}1998 had obtained a Hubble Space Telescope NICMOS image of PG\,1700+518. They also conclude that the extension to the N is a companion, but they see what we have described as a tidal tail rather as part of a ring. \subsection{The SED of the Companion and the Age of the Starburst} In Paper 1, we presented a spectrum of the companion, corrected for contamination from the QSO. We modeled the SED as a superposition of two simple stellar populations (instantaneous bursts). Our age estimates, based on a $\chi^2$ fit of the Bruzual \& Charlot \markcite{bru93}(1993) models to the spectrum of the companion, gave 0.09 (+0.04, $-0.03$) Gyr and 12.25 Gyr, respectively, for the two components. Two recent developments encourage us to try to refine these estimates: better models are now available (Bruzual \& Charlot \markcite{bru98}1998), and we now have images in bandpasses covering a wide spectral range, which can additionally constrain the stellar populations. We use the same Keck Low-Resolution Imaging Spectrograph (LRIS) data from Paper 1. Because we need fairly high resolution for detailed fitting to features in the spectrum, we restrict ourselves to the solar-metallicity spectral-synthesis models based on the Gunn \& Stryker \markcite{gun83}(1983) and Jacoby, Hunter, \& Christian \markcite{jac84}(1984) spectra. We first discuss fits to the Keck LRIS spectrophotometry, considered in isolation; then we show how the models must be modified to take into account the SED over a wider wavelength region. Although the fit to the spectrum shown in Fig.\ 3 is from our final model, it is typical of the quality of fit we find for other models we discuss here. If we try a similar approach to that of Paper 1, using only simple stellar populations from the newer models (Bruzual \& Charlot \markcite{bru98}1998), we obtain nearly the same age for the younger population (0.10 Gyr), but a much younger age for the older population (1.8 Gyr). While this new model fits the Keck spectrum much better than did the model given in Paper 1, and we could no doubt improve the fit even more by adding a third, older population, the morphological evidence suggests another approach. The presence of a strong tidal tail indicates a dominant, pre-existing disk component, and the evidence for recent star formation indicates a gas-rich system. Such a galaxy would be expected to have been forming stars over the entire lifetime of the disk. We therefore consider another range of models: those in which there has been an exponentially-decaying rate of star formation on which is superposed a single recent starburst. Specifically, we assume that the younger population can be modelled as a burst (there is no appreciable difference in this case whether the burst is taken to be instantaneous or has a finite duration of several Myr) and that the older population has an exponentially-decaying star formation beginning 10 Gyr ago. We find that the age of the younger population and the quality of the fit are insensitive to the decay rate of star formation in the older component over a reasonable range (time constants $\sim3$--10 Gyr), but that we cannot get as good fits to the observed spectrum for either an old instantaneous burst or a constant star-formation rate. For definiteness, we assume an exponential time constant of 5 Gyr. Considering only the Keck spectrophotometry, we find the best $\chi^2$ fit by combining this underlying population with a starburst with an age of 114 Myr. We now deal with additional constraints from the SED at longer wavelengths. We determine the SED of the companion over a rest-frame range from 0.42 \micron\ to 1.28 \micron\ from the AO $J$ and $H$-band images and images in two non-standard optical bands centered at 5442 and 7248 \AA, with FWHM of 1002 and 1260 \AA, respectively (see Fig.\ 2; these bandpasses have been designed to avoid strong emission lines at the redshift of PG\,1700+518). Figure 4 shows the flux densities in these four filters for a 1\arcsec\ diameter aperture centered 0\farcs4 E and 2\farcs3 N of the QSO. This region is far enough from the QSO that the photometry should not be very sensitive to the chosen scaling or other errors in the PSF subtraction process. While it does not exactly coincide with the region covered by the Keck LRIS spectrum, and there could be small differences in the two SEDs, we attempt to find a model that will fit both. We have explored a wide variety of SEDs involving superpositions of both simple and composite stellar populations based on the Bruzual \& Charlot \markcite{bru98}(1998) models; none of these gives a satisfactory fit both to the Keck LRIS spectrum and to the wide-band SED. A pure old population comes moderately close to matching the overall SED, but it fails completely to reproduce either the strong Balmer absorption spectrum or the continuum shape of the restframe UV---blue spectrum. On the other hand, the ``reddest'' reasonable model we can find that adequately fits the Keck spectrum and the optical photometry falls well below the $J$ and $H$ points (see Fig.\ 4). Even this latter model is rather unphysical, comprising a pure young population (to fit the UV---blue spectrum) and a pure old stellar population (to attempt to fit the IR photometry), with nothing in between; however, the addition of any intermediate-age component only makes the fit worse. Clearly one plausible way to raise the relative flux at longer wavelengths while retaining an early-type spectrum at shorter wavelengths is to include some sort of reddening due to dust. However, simply applying a Galactic reddening law (\ie\ assuming a dust screen between the object and the observer) is neither appropriate nor very effective. Screen-like reddening sufficient to force a fit to the overall SED results in very strong variation in extinction across the range of the Keck spectral fit, making it difficult to fit simultaneously the lines and continuum. One needs a means of effecting a substantial reddening of the IR flux with respect to the optical flux without changing the slope in the UV---blue region too much. Witt, Thronson, \& Capuano \markcite{wit92}(1992) have shown that dust distributions that are more-or-less coextensive with the stellar distribution, in addition to being more realistic, can provide exactly this sort of reddening. Such models differ from the standard reddening law (due to intervening dust) in two main respects: (1) some of the blue light removed along the sightline is compensated by scattering of blue light emitted in other directions into the line of sight, and (2) optical depth effects ensure that most of the light appearing at short wavelengths has suffered little extinction, while, at longer wavelengths, a larger portion of the total stellar distribution contributes to the emergent flux. We have used the family of ``dusty galaxy'' models calculated by Witt \etal\ \markcite{wit92}(1992) to produce sets of modified Bruzual \& Charlot \markcite{bru98}(1998) models. The dust and the stars are both assumed to have constant density within a sphere, and the only variable is the optical depth to the center at a specific wavelength. While these models are highly artificial, they are sufficient to show the general nature of the reddening due to embedded dust, and they give us a good fit to both the LRIS spectrophotometry and the wide-band SED. We have included in Fig.\ 4 the model we have found to give the best fit to both the Keck spectral data and the overall SED. This same model is shown in Fig.\ 3, both as fit to the Keck spectrum and as individual components: the starburst, represented by an 85-Myr-old instantaneous burst, and the underlying population having star formation beginning 10 Gyr ago and exponentially decaying with a time constant of 5 Gyr. Both components are reddened by the curve given by the Witt \etal\ \markcite{wit92}(1992) ``dusty galaxy'' model with a $V$-band central optical depth of 6. The fit of the model to the Keck LRIS spectrophotometry (Fig.\ 3) is remarkably good, reproducing most individual features as well as the continuum slopes quite accurately. We can account for most of the deviations. Poor fits to H$\beta$ (and, to a lesser extent, H$\gamma$) absorption are due to distortion of these profiles by general emission around the QSO not specifically associated with the companion: there is both a positive component and a negative component (at a higher velocity, from the region on the opposite side of the QSO that was used to subtract off scattered QSO light; see Paper 1 for details). Similarly, the excess peak near 3868 \AA\ is due to [\ion{Ne}{3}] emission. The broad dip between 4500 and 4600 \AA\ is due to excess subtraction of \ion{Fe}{2} emission, which is apparently spatially variable (Paper 1). We have found that we can improve the fit of the \ion{Ca}{2} $K$ line by including a secondary burst $\sim2$ Gyr ago, which may be evidence for a previous close passage of the companion, but this evidence alone is slender enough that we have chosen not to include this added complexity to the model. The only significant discrepancy that we cannot explain is the poor fit to H10 $\lambda3798$; this may simply be due to a glitch in our observed spectrum. \subsection{Towards an Evolutionary Sequence for QSOs} Although the SED model we have presented is the best fit we have found to the data, subject to the constraints of simplicity and astrophysical reasonableness that we have chosen, we can find other models that fit nearly as well. However, in exploring various models, we have found that, for reasonably plausible combinations of young and old stellar populations and reddening curves (\ie\ those that come close to fitting both the LRIS spectrophotometry and the overall SED), there is very little spread in the age of the young population. We consistently obtain ages in the range from 75 to 100 Myr, and we believe that this is a reasonable estimate of the uncertainty in our determination, subject only to any remaining uncertainty in the Bruzual \& Charlot \markcite{bru98}(1998) models themselves. We regard this robustness in the age determination as a hopeful sign for our ongoing efforts to develop an age sequence for triggering events for QSOs lying in the transition region between ultraluminous IR galaxies and the classical QSO population in the far-IR two-color diagram (Paper 1; Canalizo \& Stockton \markcite{can98}1998; Stockton \markcite{sto98b}1998). This enthusiasm is only slightly dampened by the fact that the AO imaging shows clearly that our age determination is for the interacting companion to PG\,1700+518, rather than for the host galaxy itself. Obtaining similar quality spectrophotometry for the extension of the host galaxy to the SW, which is both considerably fainter and closer to the QSO, would be extremely difficult. Nevertheless, its colors appear to be quite similar to those of the companion, and models of starbursts in interacting pairs (\eg\ Mihos \& Herquist \markcite{mih96}1996) indicate that star formation peaks at times of closest passage in both participants. The projected $\sim2$\arcsec\ ($\approx7$ kpc) separation between the QSO and companion and the projected $\sim5$ kpc tail length of the companion are entirely consistent with a close passage 85 Myr ago, if projection factors are $\sim0.5$ and mutual velocities are $\sim150$ km s$^{-1}$. We therefore believe that the age of starburst in a close, clearly interacting companion is a good surrogate for an age determined from the QSO host galaxy itself for purposes of attempting to define an evolutionary sequence. | 98 | 3 | astro-ph9803333_arXiv.txt |
9803 | astro-ph9803043_arXiv.txt | A preliminary 1.4 GHz RLF at redshift of about 0.14 is derived from the {\it Las Campanas Redshift Survey} (LCRS) and the NVSS radio data. No significant evolution has been found at this redshift in comparison to the 'local' RLF. | 98 | 3 | astro-ph9803043_arXiv.txt |
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9803 | astro-ph9803275_arXiv.txt | We consider the possibility that cosmic magnetic field, instead of being uniformly distributed, is strongly correlated with the large scale structure of the universe. Then, the observed rotational measure of extra-galactic radio sources would be caused mostly by the clumpy magnetic field in cosmological filaments/sheets rather than by a uniform magnetic field, which was often assumed in previous studies. As a model for the inhomogeneity of the cosmological magnetic field, we adopt a cosmological hydrodynamic simulation, where the field is passively included, and can approximately represent the real field distribution with an arbitrary normalization for the field strength. Then, we derive an upper limit of the magnetic field strength by comparing the observed limit of rotational measure with the rotational measure expected from the magnetic field geometry in the simulated model universe. The resulting upper limit to the magnetic field in filaments and sheets is ${\bar B}_{fs} \la 1 \mu{\rm G}$ which is $\sim10^3$ times higher than the previously quoted values. This value is close to, but larger than, the equipartition magnetic field strength in filaments and sheets. The amplification mechanism of the magnetic field to the above strength is uncertain. The implications of such a strength of the cosmic magnetic field are discussed. | The strength and morphology of the intergalactic magnetic fields remain largely unknown, because it is intrinsically difficult to observe them (for recent reviews, see Kronberg 1994; Beck et al. 1996; Zweibel \& Heiles 1997). While the magnetic fields inside typical galaxies are observed to be of order of $3-10\mu{\rm G}$, recent observations indicate that the magnetic fields of the similar strength are also common in core regions of rich clusters (for a recent discussion see En{\ss}lin \etal 1997). On the other hand, the upper limit for a large-scale field strength placed by the observed rotation measure (RM) of quasars is about $B_{IGM} \la 10^{-9} L_{\rm rev,Mpc}^{-1/2}{\rm G}$ (Kronberg 1994). Here, $L_{\rm rev,Mpc}$ is the reversal scale of the magnetic fields in units of Mpc in comoving coordinates. It was assumed the fields are uniform in direction within the reversal scale and varied as a comoving passive field in the expansion of the universe. The above value is close to the limit placed by the observed anisotropy in the cosmic microwave background radiation is $B_{IGM} < 6.8 \times 10^{-9} (\Omega_o h^2)^{1/2} {\rm G}$ (Barrow, Ferreira \& Silk 1997). Here, a uniform magnetic field was assumed. $h$ is $H_o$ in unit of $100~{\rm km/s/Mpc}$. According to a recent popular view based on both observational and theoretical cosmology, the dominant nonlinear structure in the universe is a web-like network of filaments (\eg Bond, Kofman \& Pogosyan 1996; Shectman \etal 1996). At a lower density contrast, however, bubbly walls whose intersections are in fact filaments are the dominant structures. Such large scale structures form via gravitational instability, and then turbulent motions inside the structures as well as streaming motions along the structures necessarily exist (\eg Kang \etal 1994). Hence, magnetic fields in the early universe, if existed, should have been modified and amplified by those flow motions (\eg Kulsrud \& Anderson 1992; Kulsrud \etal 1997). As a result, we expect that {\it the ``geometry'' of cosmic magnetic field, which means the spatial distribution of the field strength and its orientation, should be correlated with the large-scale nonlinear structures of the universe}. In other words, the field strength increases with the matter density and its orientation tends to align with the sheets and filaments, while its random component is associated with the local turbulent motions. The magnetic fields should be strong along the walls of cosmic bubbles, even stronger along the filamentary superclusters, and the strongest inside the clusters of the galaxies, while they are very weak inside the voids. Based on the proposition that the large scale magnetic field is correlated with the large scale structure of the universe, we consider a way to estimate the field strength which relies on the observational data. The cosmic magnetic field together with free electrons in the intergalactic medium (IGM) induces the Faraday rotation in polarized radio waves from extra-galactic sources. The observational RM data of quasars show a systematic growth of RM with redshift, $z$, and limit RM to $\sim 5~{\rm rad~m^{-2}}$ or less at $z=2.5$ (Kronberg \& Simard-Normandin 1976; Kronberg 1994 and references therein). Adopting a model for the distribution of the large scale magnetic field, these data can be used to constrain the strength of the magnetic field. For example, the upper limit estimated by Kronberg (1994) was based on a model in which the field is uniform within the reversal scale, the orientation is random, and there are $N=l_s/L_{rev}$ reversals along the line of sight to a source $l_s$ Mpc away from us. In the present paper, we re-derive the upper limit by taking a new model in which the intergalactic magnetic field is mostly confined within filaments and sheets but very weak inside voids, rather than the simple model of uniform field. Of course the field would be strongest inside clusters, but their contribution is usually excluded in the observed RM. If we take only the geometric consideration that the field distribution along the lines of sight to radio sources has a small filling factor, high only inside filaments and sheets but low inside voids, then it is obvious that the expected magnetic field strength in filaments and sheets should be higher than the value derived from the uniform field model. But this geometric model is too simple, since the electron distribution as well as the field direction are uncertain. Here, we take a more practical approach by adopting numerical data in a simulated universe which includes a magnetic field, which is evolved during the large scale structure formation. With the magnetic field distribution, specific to this simulation model, we argue that the magnetic field strength in filaments and sheets is limited to ${\bar B}_{fs} \la 1~\mu{\rm G}$. We should emphasize that {\it this limit should be approximately valid regardless of details of any model for the origin of the large scale magnetic field, provided that it is correlated with the large scale structure of the universe}. In the next section, we describe the model simulation and the resulting magnetic field geometry. In \S3, we explain the procedure to calculate the upper limit of the magnetic field strength in filaments and sheets constrained by the observed RM. In the final section, we discuss the implications of the possible existence of $\sim 1~\mu{\rm G}$ or less magnetic field in filaments and sheets. | Our result indicates that, with the present value of the observed limit in RM, the existence of magnetic field of up to $\sim 1~\mu{\rm G}$ in filaments and sheets can not be ruled out, if the cosmic magnetic field is preferentially distributed in these nonlinear structures, rather than uniformly distributed in the intergalactic space. {\it We emphasize that this result is independent of the details of how the magnetic field originated.} However, it is not certain if the dynamo processes such as the one considered in Kulsrud \etal (1997) can amplify a week seed field to the above field strength. Interestingly, there is confirmation of such a magnetic field in one case. Kim \etal (1989) reported the possible existence of a intercluster magnetic field of $0.3-0.6~\mu{\rm G}$ in the plane of the Coma/Abell 1367 supercluster. They reached the result from the observation using high dynamic range synchrotron images at $327~{\rm MHz}$. Clearly, this result should be confirmed in other regions of the sky with the newly available possibilities of low radio frequency interferometry such as with the Giant Meterwave Radio Telescope (GMRT) in India, or the new receiver systems on the Very Large Array (VLA) in the USA. We may compare the above upper limit with the strength of the magnetic field whose energy is in equipartition with the thermal energy of the gas in filaments and sheets. The equipartition magnetic field strength can be written as \begin{equation} B = 0.33 h \sqrt{T \over 3 \times 10^6 {\rm K}} \sqrt{\rho_b \over 0.3 \rho_c}~\mu{\rm G}. \label{b_equip} \end{equation} The fiducial values $T = 3 \times 10^6 {\rm K}$ and $\rho_b = 0.3 \rho_c$ may considered to be appropriate for filaments (see, Kang \etal 1994 and Figure 3). The values appropriate for sheets should be somewhat smaller. So, our upper limit set by RM is close to, but several times larger than, the equipartition magnetic field strength. Filaments and sheets with the above fiducial temperature and gas density can contribute radiation at X-ray wavelengths to the cosmological background. Using the above values and a typical size of $10h^{-1}$Mpc, the luminosity of filaments and sheets in the soft X-ray band ($0.5{-}2\,$keV) is expected to be $\la 3\times 10^{41}{\rm erg~s^{-1}}$. Actually, Soltan \etal (1996) and Miyaji \etal (1996) reported the detection of extended X-ray emission from structures of a comparable size and a luminosity of $\approx 2.5 \times 10^{43}{\rm erg~s^{-1}}$, which is correlated to Abell clusters. These observations suggest that the temperature and gas density outside clusters could even be considerably larger than the fiducial values assumed above. In any case, if these observations can be further confirmed, and the equipartition of magnetic field and gas energies is assumed, it would imply the existence of a $\sim 1\mu{\rm G}$ magnetic field on large scales outside galaxy clusters. The possible existence of strong magnetic field of $1~\mu{\rm G}$ or less in filaments and sheets has many astrophysical implications, some of which we briefly outline in the following: The large-scale accretion shocks where seed magnetic fields could be generated are probably the biggest shocks in the universe with a typical size $\ga$ a few (1-10) Mpc and very strong with a typical accretion velocity $\ga$ a few $1000~{\rm km}~{\rm s}^{-1}$. The accretion velocity onto the shocks around clusters of a given temperature, or a give mass to radius ratio $M_{cl}/R_{cl}$, is smaller in a universe with smaller $\Omega_o$, and is given as $v_{acc} \approx 0.9-1.1 \times 10^3~{\rm km~s^{-1}} [(M_{cl}/R_{cl})/(4 \times 10^{14}{\rm M}_{\odot}/{\rm Mpc})]^{1/2}$ in model universes with $0.1 \le \Omega_o \le 1$ (Ryu \& Kang 1997a). With up to $\sim 1~\mu{\rm G}$ or less magnetic field around them, the large-scale accretion shocks could serve as possible sites for the acceleration of high energy cosmic rays by the first-order Fermi process (Kang, Ryu \& Jones 1996; Kang, Rachen \& Biermann 1997). Although the shocks around clusters would be the most efficient sites for acceleration, those around filaments and sheets could make a significant contribution as well (Norman, Melrose \& Achterberg 1995). With the particle diffusion model in quasi-perpendicular shocks (Jokipii 1987), the observed cosmic ray spectrum near $10^{19}{\rm eV}$ could be explained with reasonable parameters if about $10^{-4}$ of the infalling kinetic energy can be injected into the intergalactic space as the high energy particles (if an $E^{-2}$ spectrum of cosmic rays is assumed; for a slightly steeper spectrum as suggested by radio relic sources, the efficiency is closer to 0.1, En{\ss}lin \etal 1998). The discoveries of several reliable events of high energy cosmic rays at an energy above $10^{20} {\rm eV}$ raise questions about their origin and path in the universe (a recent review is P. Biermann 1997), since their interaction with the cosmic microwave background radiation limits the distances to their sources to less than 100 Mpc, perhaps within our Local Supercluster. The Haverah Park and Akeno data indicate that their arrival directions are in some degree correlated with the direction of the Supergalactic plane (Stanev \etal 1995; Hayashida \etal 1996; Uchihori \etal 1996). In Biermann, Kang \& Ryu (1996), we noted that if the magnetic field of $\sim 1\mu{\rm G}$ or less exists inside our Local Supercluster and there exist accretion flows infalling toward the supergalactic plane, it is possible that the high energy cosmic rays above the so-called GZK cutoff ($E> 5\times 10^{19}$ eV) can be confined to the supergalactic plane sheet, an effect analogous to solar wind modulation. In each case a shock front pushes energetic particles upstream as seen from its flow. Obviously, this effect would occur only for a small part of phase space, namely those particles with a sufficiently small initial momentum transverse to the sheet. This would explain naturally the correlation between the arrival direction of the high energy cosmic rays and the supergalactic plane. Also, confinement means that for all the particles captured into the sheets, the dilution with distance $d$ is $1/d$ instead of $1/d^2$, increasing the cosmic ray flux from any source appreciably with respect to the three-dimensional dilution. So we may see sources to much larger distances than expected so far. On the other hand, particles with a larger initial momentum transverse to the sheet would be strongly scattered, obliterating all source information from their arrival direction at Earth. With the magnetic field in the intergalactic medium, charged particles passing it would not only experience deflection. It would also smear out their arrival direction as well as delay their arrival time. Plaga (1995) suggested that an exhibition of this latter effect would be the delay of the arrival times of $\gamma$-rays from a cascade caused by photons from highly time-variable extragalactic sources. However, with strong fields of $\sim 1~\mu{\rm G}$ or less in filaments and sheets intervening between the sources and us, the expected consequence would be a strong smearing of the sources rather than the delay of arrival times (Kronberg 1995). But, for details, calculations following the propagation of photons and charged particles should be done. Another interesting exploration is to study the radiation emission arising from filaments and sheets strong magnetic field of $1~\mu{\rm G}$ or less. We noted that the original estimate of the magnetic field in the plane of the Coma/Abell 1367 supercluster was based on a synchrotron radio continuum measurement (Kim \etal 1989). We also noted that the thermal Bremsstrahlung emission in the soft X-ray band may provide an interesting check on the work presented here. Whether these structures can contribute at other wavelengths, such as $\gamma$-ray energies, to the cosmological background remains an an unanswered but challenging question at this time. These issues will be discussed elsewhere. | 98 | 3 | astro-ph9803275_arXiv.txt |
9803 | astro-ph9803153_arXiv.txt | We report $J$, $H$, and $K$ photometry of 86 stars in 40 fields in the northern hemisphere. The fields are smaller than or comparable to a 4$\times$4~arcmin field-of-view, and are roughly uniformly distributed over the sky, making them suitable for a homogeneous broadband calibration network for near-infrared panoramic detectors. $K$ magnitudes range from 8.5 to 14, and $J-K$ colors from -0.1 to 1.2. The photometry is derived from a total of 3899 reduced images; each star has been measured, on average, 26.0 times per filter on 5.5 nights. Typical errors on the photometry are $\sim$~0\fm012. | The widespread availability of near-infrared (NIR) panoramic detectors has rendered possible many scientific programs which were unfeasible with single-element photometers. New and more sophisticated data acquisition and reduction techniques have successfully exploited the capabilities of the new technology while, at the same time, photometric calibration has generally relied on older pre-existing standard networks based on ``one-pixel'' photometry. Such networks include the SAAO system of Glass (\cite{glas}), expanded and rationalized by Carter (\cite{carter90}, \cite{carter95}), comprising probably the most comprehensive and best-observed list available; the MSO system defined by Jones \& Hyland (\cite{jone80}, \cite{jone82}), now more or less supplanted or absorbed into the AAO system (Allen \& Cragg \cite{alle}); the ESO system (Engels \cite{enge}; Bouchet, Schmider, \& Manfroid \cite{bouchet}); and the system which appears to the greatest extent to have inherited the original mantle of the photometry of H. L. Johnson in the 1960s, the CIT (Caltech/Cerro-Tololo) system of Frogel et al. (\cite{frog}) updated by Elias et al. (\cite{elia82}). This last forms the basis of the unpublished, but widely used, hybrid standard star set maintained at the 3.8~m UK Infrared Telescope (UKIRT). All these comprise relatively bright stars suitable for photometry at small- and medium-sized telescopes with instruments which do not have the limited well capacities of array elements. However, when observed with array detectors even on moderate-sized telescopes, stars with K $\lesssim$ 7.5 produce saturated pixels unless defocussed or observed in non-standard modes with extremely short on-chip integration times, neither of which stratagem is conducive to precise and homogeneous calibration. The ``UKIRT Faint Standards'' (Casali \& Ha-warden \cite{casa}) upon which the calibration of this work is based comprise a new set of much fainter stars chosen and observed at UKIRT to facilitate observations with panoramic detectors with limited dynamic range. However, they are relatively few in number and isolated, occur mostly around the celestial equator, and many of the stars are too faint to be useful for small- or medium-sized telescopes. Only preliminary results are currently available for the UKIRT Faint Standards, although this situation is actively being remedied by the expansion and reobservation of the list at UKIRT. We present here a set of $J$, $H$, and $K$ photometric measurements, obtained with the Arcetri NICMOS3 camera, ARNICA. The photometry comprises 86 stars in 40 fields observable from the northern hemisphere. The selection of the standard fields is described in Section~2, followed by a discussion of observing and data reduction techniques in Section~3. Section 4 presents the photometry and a comparison with other photometric systems. | We have presented new NIR photometry for 86 stars in 40 fields. The sky coverage is relatively uniform, and ideal for observatories with $\delta\,\gtrsim\,30^\circ$. On average, stars have been observed on more than five different nights, and typical errors on the photometry are 0\fm012. We find some indication of a color transformation between the ARNICA (NICMOS) photometry reported here and the original UKIRT (InSb) system, but a definitive statement awaits a larger data set designed specifically to determine such a transformation. | 98 | 3 | astro-ph9803153_arXiv.txt |
9803 | astro-ph9803015_arXiv.txt | The BeppoSAX satellite has recently opened a new way towards the solution of the long standing gamma-ray bursts' (GRBs) enigma, providing accurate coordinates few hours after the event thus allowing for multiwavelength follow-up observational campaigns. The BeppoSAX Narrow Field Instruments observed the region of sky containing GRB970111 16 hours after the burst. In contrast to other GRBs observed by BeppoSAX no bright afterglow was unambiguously observed. A faint source (1SAXJ1528.1+1937) is detected in a position consistent with the BeppoSAX Wide Field Camera position, but unconsistent with the IPN annulus. Whether 1SAXJ1528.1+1937 is associated with GRB970111 or not, the X-ray intensity of the afterglow is significantly lower than expected, based on the properties of the other BeppoSAX GRB afterglows. Given that GRB970111 is one of the brightest GRBs observed, this implies that there is no obvious relation between the GRB gamma-ray peak flux and the intensity of the X-ray afterglow. | The comprehension of the nature of the Gamma-Ray Bursts (GRBs) is a long-standing problem of a world-wide scientific community since the announcement of their discovery (Klebesadel \etal 1973). Many observational (Fishman \& Meegan 1995) and theoretical (Lamb 1995; Paczynski 1995) efforts did not succeed in understanding the origin of GRBs. The launch of the BeppoSAX satellite (Boella \etal 1997a) revolutionized the field, opening a new observational window soon after the GRB event. Due to its Gamma Ray Burst Monitor (GRBM, 40--700~keV, Frontera \etal 1997a; Feroci \etal 1997a) and its Wide Field Cameras (WFCs, 2--26~keV, Jager \etal 1997) this satellite is capable of detecting GRBs in the gamma-ray band and accurately localizing them in X-rays through a coded mask proportional counter. Five GRBs, amongst those simultaneously detected by the GRBM and the WFCs, were promptly analyzed, allowing multiwavelength follow-up observational campaigns. The first result is the BeppoSAX discovery of the X-ray afterglow of GRB970228 (Costa \etal 1997, Costa \etal 1997a) and the discovery of a related optical transient by ground-based telescopes (van Paradijs \etal 1997). Further results have been achieved with the detection of the X-ray afterglows of GRB970402 (Feroci \etal 1997b, Piro \etal 1997a), GRB970508 (Costa \etal 1997c, Piro \etal 1997b) and GRB971214 (Heise \etal 1997a, Antonelli \etal 1997). From GRB970508 an indication of an extragalactic origin has been derived through the detection of an optical transient (Bond, 1997; Djorgovski \etal 1997) and the measurement of its optical spectrum (Metzger \etal 1997), providing a lower limit of 0.835 for the redshift of the possible GRB optical afterglow. One of the most intriguing mysteries of GRB emitters is possibly solved, but the overall picture is far from clear. In fact, out of the five events for which BeppoSAX performed rapid follow up searches of a X-ray counterpart, one (GRB970111) has given a result that is significantly different from the other four. The celestial location of GRB970111 was observed by BeppoSAX just 16 hours after the GRB event, and no unambiguous evidence for an X-ray afterglow was found. A new faint source (1SAX J1528.1+1937) was detected at a flux level that is much lower than that expected on the basis of the properties of the other GRBs later observed by BeppoSAX. Here we present this detection, discuss its association with GRB970111 and the diversity from the other four BeppoSAX GRBs. | The BeppoSAX follow-up observation of the error box of GRB970111 was the first prompt follow-up observation of a GRB ever performed by an X-ray satellite. Before BeppoSAX the time-scale of a possible X-ray emission from GRB remnants was completely unknown. This first basically non-detection, therefore, could only be interpreted as an upper limit to the time-scale of the decline of an X-ray afterglow or to its flux. Now, with the detection of the X-ray afterglows of GRB970228, GRB970402, GRB970508 and GRB971214, BeppoSAX has set up a new scenario in which GRB970111 seems misplaced. Also the detection of the X-ray afterglow of a GRB (GRB970828, Remillard \etal 1997; Murakami \etal 1997) by the RossiXTE and the ASCA satellites supports the general framework for the GRBs' afterglow built by BeppoSAX. GRB970228, GRB970402 and GRB970828 showed a similar behavior, with a fading X-ray counterpart continuously decaying from the GRB main emission into the afterglow following an approximate $t^{-1.3}$ law. In the case of GRB970228, the spectral analysis (Frontera \etal 1997b) confirms the continuity between the latest GRB emission and the X-ray counterpart detected after few hours. This temporal behaviour could be explained in the framework of the fireball model (Cavallo \& Rees 1978; Rees \& Meszaros 1992) as a highly radiative expansion of a relativistic shell. GRB970508 has shown a X-ray counterpart whose decay is more complicated than the above three. The above model could still account for this different behavior, but it needs to invoke a non-uniform surrounding medium, with a density scaling as $r^{-2}$ (Vietri 1997). Whether 1SAX J1528.1+1937 is related to GRB970111 or not, this gamma-ray burst had a much faster decay than observed for any of the others. In order to make this clear, we compare a hypothetic power-law decay of GRB970111 with the ``typical'' power-law decay of GRB970228 reported in Costa \etal (1997a). Therefore, in Fig. 3 we assume that the new X-ray source is associated with the GRB and impose a power-law decay of the afterglow starting from the WFC mean flux at a time centred on the GRB X-rays duration. The needed power-law index is -1.5. \begin{figure} \epsfxsize=\hsize \centerline{\epsffile{fig_3.ps}} \vspace{-0.5cm} \caption[]{X-ray (2--10~keV) decay law of the candidate counterpart of GRB970111, compared to GRB970228. The dot-dashed and the solid horizontal lines are the mean X-ray flux for GRB970228 and GRB970111, respectively. The inclined dashed line is the decay law suggested by Costa et al. (1997a) for GRB970228. The inclined solid line shows the power-law index, 1.5, needed for connecting the WFC GRB970111 mean flux and the 1SAXJ 1528.1+1937 flux} \label{fig:Decay} \end{figure} Alternatively, assuming that 1SAX J1528.1+1937 is not related to GRB970111 we can derive a lower limit to the power-law index by using the upper limit of the MECS flux in the region of sky defined by the error box, to obtain a value very similar to the 1.5 value given above. Trying to extract GRB970111 from the group as an intrinsically different event, we note that its gamma-ray fluence is about more than three times larger than the largest of the other three. On the other hand, even if this GRB is of the ``No High Eenergy'' type (that is, it shows only weak emission above 300 keV, Pendleton et al. 1997), the ratio between the X-ray (2--10~keV) and gamma-ray (40--700~keV) fluences is about 4\%, to be compared to 20\% (2--10~keV) for GRB970228 (Frontera \etal 1997b), 5\% (2--10~keV) for GRB970402 (Nicastro \etal 1997) and 40\% (2--26~keV) for GRB970508 (Piro \etal 1997b). GRB970111 appears therefore as the one (together with the April event) with the less efficient low X-rays mechanism for energy release. Furthermore, no optical source was found in the WFC error box changing its intensity more than 0.5 magnitudes at a level of B=23 and R=22.6 from about 19 hours to about one month later (Castro-Tirado \etal 1997; Gorosabel \etal 1998). A radio search at 1.4 GHz (Frail \etal 1997) and at millimetric wavelength (Smith \etal 1997) did not find a counterpart to 1SAXJ1528.1+1937. These results support the idea that the optical, radio and millimetric channels are unefficient as well. Since GRB970111 was one of the brightest events ever detected in gamma-rays, one may conclude that its gamma-ray channel was efficient enough to dissipate most of the energy generated in the burst. Alternative interpretations of the lack of X\-/optical\-/radio afterglow of the GRB970111 may be either a very rapidly evolving afterglow, with a decay law faster than observed in the other BeppoSAX GRB afterglows, or the absence of an afterglow source. The former hypothesis would be in agreement with the model by Tavani (1997) of a decay behavior represented by a power law with index $-21/8$ due to the observation in a fixed energy band (2--10~keV) of a synchrotron emission spectrum with a rapidly evolving critical frequency. Alternatively, the latter situation could be due, as an example, to the scenario in which the event that caused the GRB occurred in a region in which the interstellar medium density is low enough (perhaps the external regions of a host galaxy) to justify the absence of an external shock, possibly responsible for the afterglow emission in the other cases (Katz \& Piran 1997). | 98 | 3 | astro-ph9803015_arXiv.txt |
9803 | gr-qc9803028_arXiv.txt | The energy range between the grand unification scale $M_{\rm GUT}\sim 10^{16}$ GeV and the Planck scale $\mpl\simeq 1.22\times 10^{19}$ GeV is crucial for fundamental physical questions and for testing current ideas about grand unification, quantum gravity, string theory. Experimental results in this energy range are of course very difficult to obtain. From the particle physics point of view, there are basically two important experimental results that can be translated into statements about this energy range (see e.g.~\cite{rev} for recent reviews): (i) the accurate measurement of gauge coupling constants at LEP that, combined with their running with energy, shows the unification of the couplings at the scale $\mgut$, provided that the running is computed including the supersymmetric particles in the low energy spectrum. And (ii) the negative results on proton decay; the lower limit on the inverse of the partial decay width for the processes $p\ra e^+\pi^0$ and $p\ra K^+\bar{\nu}$ are $5.5\times 10^{32}$ yr and $1.0\times 10^{32}$ yr, respectively, and imply a lower bound $\mgut \gsim 10^{15}$ GeV, which excludes non-supersymmetric SU(5) unification. Further improvement is expected from the SuperKamiokande experiment, which should probe lifetimes $\sim 10^{34}$ yr. From the cosmological point of view, informations on this energy range can only come from particles which decoupled from the primordial plasma at very early time. Particles which stay in thermal equilibrium down to a decoupling temperature $T_{\rm dec}$ can only carry informations on the state of the Universe at $E\sim T_{\rm dec}$. All informations on physics at higher energies has in fact been obliterated by the successive interactions. The condition for thermal equilibrium is that the rate $\Gamma$ of the processes that mantain equilibrium be larger than the rate of expansion of the Universe, as measured by the Hubble parameter $H$~\cite{KT}. The rate is given by $\Gamma =n\sigma |v|$ where $n$ is the number density of the particle in question, and for massless or light particles in equilibrium at a temperature $T$, $n\sim T^3$; $|v|\sim 1$ is the typical velocity and $\sigma$ is the cross-section of the process. Consider for instance the weakly interacting neutrinos. In this case the equilibrium is mantained, e.g., by electron-neutrino scattering, and at energies below the $W$ mass $\sigma\sim G_F^2\langle E^2\rangle \sim G_F^2T^2$ where $G_F$ is the Fermi constant and $\langle E^2\rangle$ is the average energy squared. The Hubble parameter during the radiation dominated era is related to the temperature by $H\sim T^2/\mpl$. Therefore~\cite{KT} \be \left(\frac{\Gamma}{H}\right)_{\rm neutrino} \sim \frac{G_F^2T^5}{T^2/\mpl}\simeq\left( \frac{T}{1\rm MeV}\right)^3\, . \ee Even the weakly interacting neutrinos, therefore, cannot carry informations on the state of the Universe at temperatures larger than approximately 1 MeV. If we repeat the above computation for gravitons, the Fermi constant $G_F$ is replaced by Newton constant $G=1/\mpl^2$ (we always use units $\hbar =c=k_B=1$) and at energies below the Planck mass \be \left(\frac{\Gamma}{H}\right)_{\rm graviton} \sim \left(\frac{T}{\mpl}\right)^3\, . \ee The gravitons are therefore decoupled below the Planck scale. (At the Planck scale the above estimate of the cross section is not valid and nothing can be said without a quantum theory of gravity). It follows that relic gravitational waves are a potential source of informations on very high-energy physics. Gravitational waves produced in the very early Universe have not lost memory of the conditions in which they have been produced, as it happened to all other particles, but still retain in their spectrum, typical frequency and intensity, important informations on the state of the very early Universe, and therefore on physics at correspondingly high energies, which cannot be accessed experimentally in any other way. It is also clear that the property of gravitational waves that makes them so interesting, i.e. their extremely small cross section, is also responsible for the difficulties of the experimental detection.\footnote{Thinking in terms of cross-sections, one is lead to ask how comes that gravitons could be detectable altogheter, since the graviton-matter cross section is smaller than the neutrino-matter cross section, at energies below the $W$-mass, by a factor $G^2/G_F^2\sim 10^{-67}$ and neutrinos are already so difficult to detect. The answer is that gravitons are bosons, and therefore their occupation number per cell of phase space can be $n_k\gg 1$; we will see below that in interesting cases, in the relic stochastic background we can have $n_k\sim 10^{40}$ or larger, and the squared amplitude for exciting a given mode ot the detector grows as $n_k^2$. So, we will never really detect gravitons, but rather classical gravitational waves. Neutrinos, in contrast, are fermions and for them $n_k\leq 1$.} With the very limited experimental informations that we have on the very high energy region, $\mgut \lsim E\lsim \mpl$, it is unlikely that theorists will be able to foresee all the interesting sources of relic stochastic background, let alone to compute their spectra. This is particularly clear in the Planckian or string theory domain where, even if we succeed in predicting some interesting physical effects, in general we cannot compute them reliably. So, despite the large efforts that have been devoted to understanding possible sources, it is still quite possible that, if a relic background of gravitational waves will be detected, its physical origin will be a surprise. In this case a model-independent analysis of what we can expect might be useful. In this paper we discuss whether, from the experience gained with various specific computations of relic backgrounds, it is possible to extract statements or order of magnitude estimates which are as much as possible model-independent. These estimates would constitute a sort of minimal set of naive expectations, that could give some orientation, independently of the uncertainties and intricacies of the specific cosmological models. We discuss typical values of the frequencies involved and of the expected intensity of the background gravitational radiation, and we try to distinguish between statements that are relatively model-independent and results specific to given models. The paper is written having in mind a reader interested in gravitational-wave detection but not necessarily competent in early Universe cosmology nor in physics at the string or Planck scale, and a number of more technical remarks are relegated in footnotes and in an appendix. We have also tried to be self-contained and we have attempted to summarize and occasionally clean up many formulas and numerical estimates appearing in the literature. The organization of the paper is as follows. In sect.~2 we introduce the variables most commonly used to describe a stochastic background of gravitational waves. We give a detailed derivation of the relation between exact formulas for the signal-to-noise ratio, and approximate but simpler characterizations of the characteristic amplitude and of the noise. The former variables are convenient in theoretical computations while the latter are commonly used by experimentalists, so it is worthwhile to understand in some details their relations. In sect.~3 we apply these formulas to compute the sensitivity to a stochastic background that could be obtained with a second Virgo interferometer correlated with the first, and we compare with various others detectors. We find that in the Virgo-Virgo case the noise which would give the dominant limitation to the measurement of a stochastic background is the mirror thermal noise, and we give the sensitivity for different forms of the relic GW spectrum. In sect.~4 and 5 we discuss estimates of the typical frequency scales. We examine the statements leading to the conclusion that Virgo/LIGO will explore the Universe at temperatures $T\sim 10^7$ GeV (sect.~4), and in the appendix we discuss some qualifications to this statement. In sect.~5 we discuss the possibility to reach much higher energy scales, including the typical scales of grand unification and quantum gravity. In sect.~6 we discuss different scenarios, depending on how the inflationary paradigm is implemented. Characteristic values of the intensity of the spectrum and existing limits and predictions are discussed in sect.~7. Sect.~8 contains the conclusions. | Present GW experiments have not been designed especially for the detection of GW backgrounds of cosmological origin. Nevertheless, there are chances that in their frequency window there might be a cosmological signal. The most naive estimate of the frequency range for signals from the very early Universe singles out the GHz region, very far from the region accessible to ground based interferometers, $f<$ a few kHz, or to resonant masses. To have a signal in the accessible region, one of these two conditons should be met: either we find a spectrum with a long low-frequency tail, that extends from the GHz down to the kHz region, or we have some explosive production mechanism much below the Planck scale. As we have discussed, both situations seem to be not at all unusual, at least in the examples that have been worked out to date. The crucial point is the value of the intensity of the background. With a very optimistic attitude, one could hope for a signal, present just in the 10Hz-1~kHz band, with the maximum intensity compatible with the nucleosynthesis bound, $\hogw\sim$ a few $\times 10^{-6}$ (or even $10^{-5}$, stretching all parameters to the maximum limit). Such an option is not excluded, and the fact that such a background is not predicted by the mechanisms that have been investigated to date is probably not a very strong objection, given our theoretical ignorance of physics at the Planck scale and the rate at which new production mechanisms have been proposed in recent years, see the reference list. However, with more realistic estimates, on general grounds it appears difficult to predict a background that in the kHz region exceeds the level $\hogw\sim$ a few $\times 10^{-7}$, independently of the production mechanism. This should be considered the minimum detection level at which a significant search can start. Such a level is beyond the sensitivity of first generation experiments, unless a ground based interferometer is correlated with a second interferometer, located at a distance small enough so that a significant correlation is possible, and large enough to decorrelate local noises. A few tens of kilometers would probably be the right order of magnitude. In this case we could detect a signal at the level $2\times 10^{-7}$ in one year of integration time, with SNR=1.65, and the level of a few $\times 10^{-8}$ could be reached with longer integration time (and possibly allowing for a slightly worse confidence level; this could make some sense in a stochastic search because the SNR increases with time, and the hint of a signal at low SNR would provide a strong motivation for pursuing the search). The difficulties of such a detection are clear, but it should be stressed that the payoff of a positive result would be enormous, opening up a window in the Universe and in fundamental high-energy physics that will never be reached with particle physics experiments. \vspace*{5mm} {\bf Acknowledgments.} I am very grateful to Adalberto Giazotto for many interesting discussions and stimulating questions, which prompted me to write down this paper. I thank Valeria Ferrari and Raffaella Schneider for discussing with me their unpublished results on astrophysical backgrounds. I also thank for useful discussions or comments on the manuscript Pia Astone, Danilo Babusci, Carlo Baccigalupi, Alessandra Buonanno, Massimo Cerdonio, Eugenio Coccia, Stefano Foffa, Maurizio Gasperini, Marco Lombardi, Emilio Picasso, Riccardo Sturani and Andrea Vicer\`e. \appendix | 98 | 3 | gr-qc9803028_arXiv.txt |
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9803 | astro-ph9803290_arXiv.txt | \noindent We describe a method of determining the system parameters in non-eclipsing interacting binaries. We find that the extent to which an observer sees the shape of the Roche-lobe of the secondary star governs the amount of distortion of the absorption line profiles. The width and degree of asymmetry of the phase-resolved absorption line profiles show a characteristic shape, which depends primarily on the binary inclination and gravity darkening exponent. We show that, in principle, by obtaining high spectral and time resolution spectra of quiescent cataclysmic variables or low mass X-ray binaries in which the mass-losing star is visible, fitting the shape of absorption line profiles will allow one to determine not only the mass function of the binary, but also the binary inclination and hence the mass of the binary components. | The determination of the binary inclination in non-eclipsing interacting binaries has been problematic for many years. One method widely used to determine the binary inclination in dwarf novae and the soft X-ray transients is to measure the ellipsoidal variations of the late-type star (Warner 1995; van Paradijs \& McClintock 1995). These variations are caused by the companion star presenting differing aspects of its distortion to the observer, giving rise to a double-humped modulation whose amplitude is strongly dependent on the binary inclination. However, one problem in measuring the ellipsoidal variations of the secondary star is that the accretion disc can contribute a significant amount of flux at optical wavelengths (e.g. typically about 10--50 per cent in the soft X-ray transients). This contribution must be accounted for if the binary inclination is to be determined using optical light curves (see Charles 1996 and references within). Here we describe a technique which uses the information available about the shape of the Roche-lobe of the secondary star, and its effect on the shape of the absorption line profiles. The parameters we measure are insensitive to the disc contribution, but dependent on the binary inclination and gravity darkening exponent. We will first give a brief description of the model and then describe the effect of the mass ratio, inclination, limb and gravity darkening on the shape of the line profiles. Finally, we will simulate data and proceed to fit it using the model. | By obtaining high spectral and time resolution spectra of cataclysmic variables or low mass X-ray binaries, systems in which the absorption features of the secondary star can be seen, one can perform a radial velocity study of the secondary star and also determine the binary mass ratio (see Marsh, Robinson \& Wood 1994 for a full description of this kind of study). However, as pointed out in this paper, it is also possible to extract the binary inclination by fitting the absorption lines profiles of the secondary star. Some of the assumptions inherent in the model which predicts the shape of the line profiles should be noted. The main assumption in the model is that the observed surface of constant optical depth coincides with the Roche potential surface. However, we find that this only changes the shape of the line profiles by less than 1 per cent. Zonal flow patterns and/or dark- hot-spots can have an appreciable effect on the shape of the line profiles. However, the extent of these effects depends very much on the magnitude of the zonal flows and the size of the spots. Simulations show that only very low velocities are predicted for zonal flows (Martin \& Davey 1995). The characteristic modulations in the width of the line profiles are, in principle, similar to the ellipsoidal modulations, which are due to the observer seeing differing aspects of the gravitationally distorted secondary star as it orbits a compact object. However, this method uses the velocity information across the star as well as the projected area. It should also be noted that unlike modelling the ellipsoidal variations in the optical, where the accretion disc contribution must be taken into account, this method is independent of the disc contribution (as long as the disc contribution is not such that it totally swamps the secondary star features). This implies all emission and continuum sources, i.e. from a bright spot or a gas stream, will only affect the determination of the veiling factor and not the shape of the absorption lines. However, care must be taken in choosing the lines to use for this kind of analysis; the lines must be clear of any weak emission features arising from the disc or accretion flow. The determination of the binary inclination is essential if one wants to obtain the mass of the binary components in non-eclipsing binaries. The method described here can be applied to bright cataclysmic variable stars or low mass X-ray binaries, in which one can resolve the absorption lines of the secondary star, and where the amount of X-ray heating is small. The effect of heating implies that one cannot use temperature sensitive absorption lines such as the Na $\sc i$ 8183-8184 \AA\ doublet in the modelling (see Friend et al. 1990 and references within for the effects of irradiation). Such systems are the dwarf novae and the soft X-ray transients. This method allows one to determine all the system parameters from a high spectral resolution (few km~s$^{-1}$) spectroscopic study. We have described a method of determining the binary inclination in non-eclipsing interacting binaries by fitting the shape of the absorption lines arising from the secondary star. We find that the amount of distortion of the absorption line profiles is primarily due to the extent to which an observer sees the shape of the Roche-lobe of the secondary star. We show that, in principle, by obtaining high spectral and time resolution spectra of quiescent dwarf novae or the soft X-ray transients, where the disc is low, fitting the shape of absorption line profiles will allow one to determine the binary inclination. Our simulations show that previous efforts to determine the inclination are flawed, and give systematically lower values for the inclination. | 98 | 3 | astro-ph9803290_arXiv.txt |
9803 | astro-ph9803259_arXiv.txt | \cite{peeb1} showed that in the gravitational instability picture galaxy orbits can be traced back in time from a knowledge of their current positions, via a variational principle. We modify this variational principle so that galaxy redshifts can be input instead of distances, thereby recovering the distances. As a test problem, we apply the new method to a Local Group model. We infer $M = 4$ to $8 \times 10^{12} M_{\sun} $ depending on cosmology, implying that the dynamics of the outlying Local Group dwarves are consistent with the timing argument. Some algorithmic issues need to be addressed before the method can be applied to recover nonlinear evolution from large redshift surveys, but there are no more difficulties in principle. | Phase space is six--dimensional and therefore six numbers for each particle will specify the dynamics completely. The standard approach is to define initial conditions as six numbers (positions and velocities in three dimensions) and integrate those forward in time. This is the usual $N$-body approach to the problem. Alternatively, it is not necessary to define those six numbers at only one time. It is equally well possible to split them, such that three numbers will be known at an initial time, and three at a final time, where the former are derived from physical arguments about the inital state, and the latter are provided by data on the current state of the system. This is the boundary value approach. The usual boundary value for structure formation by hierarchical clustering at initial times is the gravitational instability requirement that initial peculiar velocities must vanish: therefore, it is possible to express the orbits as a sum of growing modes. This is what perturbation theory is designed to do, and linear perturbation theory and its extension, the Zel'dovich approximation are in wide use. However, as structure formation is non-linear on small scales, a non-linear formalism would be more useful. Non-linear perturbation theory exists (e.g. \cite{nus}, \cite{gram93a}, \cite{gram93b}, \cite{buch}) but is very complicated and has convergence problems, i.e. the dynamics at early times can be fitted to a high accuracy only at the expense of a good fit at later times. A different method (cf \cite{peeb1}) of addressing the boundary value problem is to use the variational principle. The basic method is to start with a parameterisation of the orbits which satisfies the boundary conditions and then adjust the parameters until a stationary point of the action is found. (A variant, suggested by \cite{gia93} and implemented by \cite{susp}, parameterises the density and velocity fields rather than the orbits.) This may seem like perturbation theory because it is a method which attempts to get better and better approximations of galaxy orbits but there are important differences. The main one is that perturbation theory attempts to fit early times even at the expense of accuracy at later times, whereas the variational principle spreads out the errors more uniformly over all times. The variational principle is also algorithmically more straightforward to do to higher orders. The main disadvantage of Peebles' original method is that it requires the input of distances to recover redshifts. Redshifts, however, are easy to measure, whereas distances are not. We therefore change to recovering distances from redshifts. The recent nearly all-sky redshift surveys QDOT and PSC$z$ provide a strong motivation for variational methods in redshift space. \cite{shay93} and \cite{shay95} have pointed out that this could be achieved by treating the distance boundary condition as a parameter and fitting that until the recovered redshifts agree with the measured ones. Another approach is to modify the variational principle until the boundary conditions are of the desired form. We take such an approach and so does \cite{whit98} but the details in his treatment differ from ours. The attractive feature of this approach is that it only requires small modifications of Peebles' elegant original method. In this paper we develop a redshift space variational method, apply it to a small system --- the Local Group ---, and point out the algorithmic problems that need to be solved to extend to larger systems. Various aspects of the Local Group dynamics have been studied by \cite{peeb1}, \cite{peeb2}, \cite{peeb4}, and \cite{dunn}. The new feature of our analysis is that we attempt to constrain the masses of the Milky Way, M31, and an underlying distribution of unclustered matter by using a likelihood approach. Two points need to be made about these variational methods in general: Firstly, the true solutions of the variational action need not be a minimum even though the colloquial use of `least action' has stuck (cf \cite{peeb2}). Secondly, the equations solved are the same as for $N$-body integration, only the approximations used to solve them are different. In neither of the two approaches do particles have to be galaxies -- they may well be samplers of some underlying distribution function and as \cite{branch} and \cite{dunn2} have pointed out, applying the variational method using galaxies only, neglecting biasing and unclustered matter, leads to wrong results. | The $N$-body check proves that the code works and is ready to be extended to larger systems. The main problem with a system like the Local Group is the occurrence of multiple solutions. We have no guarantee that the solutions we find are the real ones, even if they fit all redshifts perfectly and predict distances which are not too far out from the measured ones. In fact, some tests suggests that at least for certain masses, several possible solutions are very close to each other so it is easy to pick the wrong one. This problem will not occur in larger systems, so they should in fact be easier to deal with. In future work, two issues will need to be addressed: First, we need an approximate and more efficient way of doing the force calculations in equations \ref{reals} and \ref{reds}. The direct sum that we have used needs to be replaced by a standard $N$-body method. Second, we need more efficient convergence. As mentioned above, the main difficulty lies in preventing the coefficients from finding two-cycles instead of fixed points of the iteration. A faster yet still robust method for dealing with this problem is needed. Even when these problems are solved, the variational calculations will never have as many particles as $N$-body simulations. The reason for pursuing this method is that unlike the $N$-body calculation, which can only reproduce the current state of the system in a statistical sense, the variational method can fit the current state exactly. {\vskip 10pt} | 98 | 3 | astro-ph9803259_arXiv.txt |
9803 | astro-ph9803284_arXiv.txt | Ongoing searches for supernovae (SNe) at cosmological distances have recently started to provide a link between SN Ia statistics and galaxy evolution. We use recent estimates of the global history of star formation to compute the theoretical Type Ia and Type II SN rates as a function of cosmic time from the present epoch to high redshifts. We show that accurate measurements of the frequency of SN events in the range $0<z<1$ will be valuable probes of the nature of Type Ia progenitors and the evolution of the stellar birthrate in the universe. The {\it Next Generation Space Telescope} should detect of order 20 Type II SNe per $4'\times 4'$ field per year in the interval $1<z<4$. | The remarkable progress in our understanding of faint galaxy data made possible by the combination of HST deep imaging (Williams \etal 1996) and ground-based spectroscopy (Lilly \etal 1995; Ellis \etal 1996; Cowie \etal 1996; Steidel \etal 1996), has recently permitted to shed some light on the evolution of the stellar birthrate in the universe, to identify the epoch $1\lta z\lta 2$ where most of the optical extragalactic background light was produced, and to set important contraints on galaxy formation scenarios (Madau \etal 1998; Steidel \etal 1998). While one of the biggest uncertainties in our knowledge of the emission history of the universe is probably represented by the poorly constrained amount of starlight that was absorbed by dust and reradiated in the IR at early and late epochs, one could also imagine the existence of a large population of relatively old or faint galaxies still undetected at high-$z$, as the color-selected ground-based and {\it Hubble Deep Field} samples include only the most actively star-forming young objects. It is then important at this stage to devise different observational strategies, free of some of the biases that plague current galaxy surveys, and to make testable predictions for future astronomical capabilities, such as SIRTF, FIRST and NGST. Here, we shall focus our attention on the rate of supernova (SN) explosions in the universe. An obvious reason to consider SNe is purely observational, i.e. the fact that they are very bright objects, with luminosities as high as $10^{10}~L_\odot$, and are point-like sources, making their detection possible even at very large distances/redshifts. More in general, the evolution of the SN rate with redshift contains unique information on the star formation history of the universe, the initial mass function (IMF) of stars, and the nature of the binary companion in Type Ia events. All are essential ingredients for understanding galaxy formation, cosmic chemical evolution, and the mechanisms which determined the efficiency of the conversion of gas into stars in galaxies at various epochs (e.g. Madau \etal 1996; Madau, Pozzetti, \& Dickinson 1997; Renzini 1997). While the frequency of ``core-collapse supernovae'', SN~II and possibly SN~Ib/c, which have short-lived progenitors (e.g. Wheeler \& Swartz 1993) is essentially related, for a given IMF, to the instantaneous stellar birthrate of massive stars, Type Ia SNe -- which are believed to result from the thermonuclear disruption of C-O white dwarfs in binary systems -- follow a slower evolutionary clock, and can then be used as a probe of the past history of star formation in galaxies (e.g. Branch \etal 1995; Ruiz-Lapuente, Canal, \& Burkert 1997; Yungelson \& Livio 1998). The recent detection of Type Ia SNe at cosmological distances (Kim \etal 1997; Garnavich \etal 1998; Perlmutter \etal 1998) allow for the first time a detailed comparison between the SN rates self-consistently predicted by stellar evolution models that reproduce the optical spectrophotometric properties of field galaxies, and the observed values. In this {\it Letter} we show how accurate measurements at low and intermediate redshifts of the frequencies of Type II(+Ib/c) and Ia SNe could be used as an independent test for the star formation and heavy element enrichment history of the universe, and significantly improve our understanding of the intrinsic nature and age of the populations involved in the SN explosions. A determination of the amount of star formation at early epochs is of crucial importance, as the two competing scenarios for galaxy formation, monolithic collapse -- where spheroidal systems formed early and rapidly, experiencing a bright starburst phase at high-$z$ (Eggen, Lynden-Bell, \& Sandage 1962; Tinsley \& Gunn 1976) -- and hierarchical clustering -- where ellipticals form continuosly by the merger of disk/bulge systems (White \& Frenk 1991; Kauffmann \etal 1993) and most galaxies never experience star formation rates in excess of a few solar masses per year (Baugh \etal 1998) -- appear to make rather different predictions in this regard. We show how, by detecting Type II SNe at high-$z$, the {\it Next Generation Space Telescope} should provide an important test for distinguishing between different scenarios of galaxy formation. | We have investigated the link between SN statistics and galaxy evolution. Using recent determinations of the star formation history of field galaxies from the present epoch to high-$z$, and a simple model for the evolutionary history of the binary system leading to a Type Ia event, we have computed the theoretical Type Ia and Type II SN rates as a function of cosmic time. While significant uncertainties still remain in these estimates, we believe the calculations presented in this {\it Letter} offer a first, realistic glimpse to the evolution of the cosmic supernova rates with cosmic time. Our main results can be summarized as follows. \begin{itemize} \item At the present epoch, the predicted Type II(+Ib/c) frequency appears to match remarkably well the observed local value. The obvious caveat is the well known fact that the SN II rate is a sensitive function of the lower mass cutoff of the progenitors, $m_l$. Values as low as $m_l=6\,\msun$ (Chiosi, Bertelli, \& Bressan 1992) or as high as $m_l=11\,\msun$ (Nomoto 1984) have been proposed in the literature: adopting a lower mass limit of 6 or 11 $\msun$ would increase or reduce our Type II rates by a factor 1.5, respectively. Note also that rates obtained from traditional distant (beyond 4 Mpc) sample might need to be increased by a factor of 1.5--2 because of severe selection effects against Type II's fainter than $M_V=-16$ (Woltjer 1997). \item In the interval $0\lta z\lta 1$, the predicted rate of SN Ia is a sensitive function of the characteristic delay timescale between the collapse of the primary star to a WD and the SN event. Accurate measurements of SN rates in this redshift range will improve our understanding of the nature of SN~Ia progenitors and the physics of the explosions. Ongoing searches and studies of distant SNe should soon provide these rates, allowing a universal calibration of the Type Ia phenomenon. \item While Type Ia rates at $1\lta z\lta 2$ will offer valuable information on the star formation history of the universe at earlier epoch, the full picture will only be obtained with statistics on Type Ia and II SNe at redshifts $2<z<4$ or higher. At these epochs, the detection of Type II events must await the {\it Next Generation Space Telescope} (NGST). A SN II has a typical peak magnitude $M_B\approx -17$ (e.g. Patat \etal 1994): placed at $z=3$, such an explosion would give rise to an observed flux of 15 nJy (assuming a flat cosmology with $q_0=0.5$ and $H_0=50\,h_{50}\kmsmpc$) at 1.8 \micron. At this wavelength, the imaging sensitivity of an 8m NGST is 1 nJy ($10^4$ s exposure and $10\sigma$ detection threshold), while the moderate resolution ($\lambda/\Delta \lambda=1000$) spectroscopic limit is about 50 times higher ($10^5$ s exposure per resolution element and $10\sigma$ detection threshold) (Stockman \etal 1998). The several weeks period of peak rest-frame blue luminosity would be stretched by a factor of $(1+z)$ to few months. Figure 3 shows the cumulative number of Type II events expected per year per $4'\times 4'$ field. Depending on the history of star formation at high redshifts, the NGST should detect between 7 (in the merging model) and 15 (in the monolithic collapse scenario) Type II SNe per field per year in the interval $2<z<4$. The possibility of detecting Type II SNe at $z\gta5$ from an early population of galaxies has been investigated by Miralda-Escud\'e \& Rees (1997). By assuming these are responsible for the generation of all the metals observed in the Lyman-$\alpha$ forest at high redshifts, a high baryon density ($\Omega_bh_{50}^2=0.1$), and an average metallicity of $0.01Z_\odot$, Miralda-Escud\'e \& Rees estimate the NGST should observe about 16 SN II per field per year with $z\gta 5$. Note, however, that a metallicity smaller by a factor $\sim 10$ compared to the value adopted by these authors has been recently derived by Songaila (1997). For comparison, the models discussed in this {\it Letter} predict between 1 and 10 Type II SNe per field per year with $z\gta 4$. \end{itemize} | 98 | 3 | astro-ph9803284_arXiv.txt |
9803 | astro-ph9803237_arXiv.txt | We present quantitative statistical evidence for a $\gamma$-ray emission halo surrounding the Galaxy. Maps of the emission are derived. EGRET data were analyzed in a wavelet-based non-parametric hypothesis testing framework, using a model of expected diffuse (Galactic $+$ isotropic) emission as a null hypothesis. The results show a statistically significant large scale halo surrounding the center of the Milky Way as seen from Earth. The halo flux at high latitudes is somewhat smaller than the isotropic $\gamma$-ray flux at the same energy, though of the same order ($O(10^{-7}$--$10^{-6})$ ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$ above 1 GeV). | The study of diffuse high-energy $\gamma$-ray emission has proceeded along two complementary paths. The first is via the use of parametric methods (see e.g. \citeasnoun{hunter}), where a physically motivated model described by a small number of parameters is fit to the data. The alternate approach (see e.g. \citeasnoun{chen} and \citeasnoun{willis}) is non-parametric, where one attempts to find the flux distribution as a function of position without reference to a particular model. The two classes complement each other in the sense that the strengths of one are often the weaknesses of the other. For instance, while a parametric fit usually gives some global measure of how good the fit is, but does not tell you where the model fails (e.g., ``the model did not account for this blob of flux over there''). The non-parametric approach can give this information, but unlike the parametric analysis, the quantitative assessment of the results is complicated by the effects of statistical bias. This latter point should be born in mind when examining the results presented in this paper. Non-parametric analysis of photon-limited data is beset by certain difficulties, key of which is that Poisson noise is neither stationary nor additive. Thus, the results of what would be a straightforward analysis (e.g. smoothing by a Gaussian kernel) for data contaminated by white noise are more difficult to interpret. We therefore apply a new wavelet-based technique which has the following characteristics: \begin{itemize} \item{Rigorous treatment of Poisson statistics.} \item{Assessment (in some sense) of the statistical significance of the results.} \item{Spatial adaptivity, in that structures at different size scales are recovered automatically.} \end{itemize} In this paper we apply non-parametric analysis to EGRET data taken during Phases 1-4 of the Compton Gamma Ray Observatory (CGRO) mission. Below we describe the non-parametric analysis method, and present results from EGRET data showing an extended halo of $\gamma$-ray emission apparently surrounding the center of the Milky Way. We show that this halo is statistically very strong in the data, and not obviously attributable to any systematic effect of the analysis or instrument, from which we conclude that it is most likely of astrophysical origin. We conclude with some brief discussion on the possible origins and implications of the $\gamma$-ray halo. This work expands upon that first presented in \citeasnoun{dixon1} and \citeasnoun{dixon2}. | We have presented strong statistical evidence for a large scale anisotropic excess in high energy $\gamma$-rays. Examination of our maps indicates that this excess may originate in the Galactic halo, though our results in no way rule out a local origin. The spectrum appears to be broadband; detailed investigation of the spectral properties will be the subject of future work. The origin of the halo is unclear. Emission from inverse Compton in a large halo \cite{moskalenko} appears to be a good candidate but it remains to be seen whether it can account for the entire observations in detail. A definitive answer on this topic will perhaps require a ``smoking gun'' in this or another energy band, or may have to wait for the GLAST mission. Further, as noted by previous authors (\citeasnoun{smialkowski}; \citeasnoun{chen}; \citeasnoun{strong1}), the existence of such an extended excess may impact estimates of the extragalactic $\gamma$-ray background. | 98 | 3 | astro-ph9803237_arXiv.txt |
9803 | astro-ph9803001_arXiv.txt | We carry out CCD photometry of galaxies in the $5.25$ square region centered on Coma cluster down to $M_R=-16.0$, beyond the limit of conventional morphological classification. We use the angular two-point correlation function as well as radial profiles in order to characterize the luminosity segregation. We find strong luminosity segregation for our total sample over the magnitude range of $-20 \leq M_R \leq -16$, which is not entirely accounted for in terms of the morphology-density relation that is known to exist only for bright galaxies. We use a single consistent parameter, the degree of luminosity concentration, to parameterize the morphology of galaxies over the wide magnitude range, where both giant and dwarf galaxies are included. Galaxies with high central concentration (HCC) show strong luminosity segregation, \ie their clustering strength depends strongly on luminosity while those with low central concentration (LCC) show almost no luminosity segregation. Radial density profile shows that brighter HCC-type galaxies tend to more strongly concentrate near the cluster center while LCC-type galaxies do not show such a dependence on luminosity. We show that these results are tenable against the contamination by field galaxies and uncertainties in our method of classification. | Some studies suggest that luminous galaxies are clustered more strongly than faint galaxies. This phenomenon is referred to as the luminosity segregation (hereafter LS). The LS can be interpreted as a result of either primordial effects or environmental effects. Explanations based on the primordial effect include the biased Cold Dark Matter (CDM) model in which galaxies form at high peaks in the density field (e.g., \cite{kai84}; \cite{dav85}). It is known that the LS is naturally predicted by this model. Valls-Gabaud, Alimi, \& Blanchard (1989) pointed out that the correlation strength has a positive dependence on galaxy luminosity on the basis of the biased-CDM model. White et al. (1987) also predicted that clustering strength is a strong function of the circular velocity of galaxies, which in turn correlates with the luminosity as indicated by Tully-Fisher or Faber-Jackson relations. In terms of environmental effects, the LS can be regarded as a result of frequent merging or other dynamical mechanisms in the vicinity of the cluster core. Observational evidence for the LS has been uncertain and controversial. Some found positive results (e.g., Capelato et al. 1980; Dom\'inguez-Tenreiro \& Pozo-Sanz 1988; Davis et al. 1988), and others found negative results (e.g., Phillips \& Shanks 1987; Einasto 1991). The clustering property of galaxies is also correlated with their morphology, which is so-called the morphology segregation, or morphology-density relation (\cite{dre80b}). If the morphology segregation is the fundamental correlation, the LS would be naturally expected. Early-type galaxies show a stronger degree of clustering than late-type galaxies, and at the same time, early-type galaxies are on the average brighter than late-type galaxies (e.g., \cite{eep88}). Consequently, galaxies which are more strongly clustered are brighter than those which are less clustered, which is the LS. On the contrary, if the LS is an essential correlation, the morphology segregation would also be expected. Accordingly, it is critical to see if the LS is observed {\it within a given morphological type} in order to disentangle the coupling of the LS and morphology segregation. Only a few such studies have been made so far (\cite{ein91}; \cite{lov95}). There are also few studies on the LS among dwarf galaxies. Binggeli, Tammann, \& Sandage ($1987$) found in the Virgo cluster that nucleated dwarf ellipticals (dEs) are more strongly concentrated towards the cluster center than nonnucleated dEs. Ferguson \& Sandage ($1989$) found in the Virgo and Fornax clusters that the faint ($M_B>-13.3$) nonnucleated dEs have the distribution identical to that of the E/S0 galaxies and bright nucleated dEs. They also found that the faint nonnucleated dEs are more strongly concentrated on the cluster center than the bright nonnucleated dEs. Thompson \& Gregory ($1993$) showed in the Coma cluster that both dEs and dwarf spheroidals have the same distributions as that of giant early-type galaxies. Morphological classifications of these studies are based on the eye inspection and Thompson \& Gregory's criteria used for Coma dwarfs are slightly different from those of Sandage and collaborators used for Virgo dwarfs. One of the reasons why the LS has not been examined systematically is the difficulty in sampling a large number of faint galaxies with known absolute magnitude. Clusters of galaxies are good targets to address the problem of the LS in high density environments. In particular, in nearby clusters we can sample intrinsically faint galaxies, which are critically important to the study of the LS. However, nearby clusters have such a large apparent size that we cannot survey whole the cluster with a CCD which has a small physical size. In this study, we present the angular two-point correlation function as well as the radial profiles of galaxies in the Coma cluster on the basis of a large homogeneous sample covering a wide magnitude range $-20 \leq M_R \leq -16$, where both giant and dwarf galaxies are included. We examine if there is a difference in these properties between the galaxies which have high central concentration of the surface brightness distribution and those with low central concentration. Our sample is made available by three new techniques; CCD mosaic, semi-automated data reduction/analysis software, and quantitative and objective classification of morphological type of galaxies based on surface photometry parameters. In section $2$, we briefly explain our imaging observation of the Coma cluster. In section $3$, we describe the data reduction procedures, calibration methods which are special to our camera, and the method of constructing a homogeneous galaxy sample. In section $4$, we describe star-galaxy discrimination, evaluation of the number of contaminated field galaxies, and the method of classifying galaxies according to the surface brightness concentration. We present our results in section $5$. Finally, reliability of our results is discussed in section $6$ in terms of the effects of uncertainties in our classification and contamination by background galaxies. A comparison with previous studies are given in section $7$. | We have carried out a wide-field galaxy survey in the Coma cluster region with a mosaic CCD camera to study the clustering properties of cluster members. We have investigated the luminosity segregation (LS) quantitatively by measuring the angular two-point correlation function and radial distribution over the magnitude range of $-20 \leq M_R \leq -16$, where both giant and dwarf galaxies are included. Our analysis of the galaxy distribution based on the morphology-classified galaxy samples with the unprecedentedly deep limiting magnitude has yielded the following main results: \begin{enumerate} \item We have found that the galaxies with a high central concentration in surface brightness profile (the HCC type) have strong luminosity segregation while the galaxies with a low central concentration (the LCC type) show almost no luminosity segregation, \ie the strength of clustering of the LCC-type galaxies does not depend on luminosity. \item We have found strong segregation in luminosity for the total sample of Coma cluster galaxies. This is because the majority of the total sample is comprised of the HCC-type galaxies which show strong luminosity segregation. \item Brighter HCC-type galaxies tend to more strongly concentrate near the cluster center than fainter HCC-type galaxies, while the LCC-type galaxies do not show such dependence on luminosity in the density profile. \end{enumerate} We have shown that these results are tenable against the contamination by field galaxies and uncertainty in our method of morphological classification. Our results suggest that it would be necessary to consider this type-dependence of the LS in the study of clustering evolution with $\acf$. In order to see the universality of our results for clusters of galaxies in general, further studies for other clusters are necessary to investigate the luminosity and morphology dependence of galaxy distribution. It is also desired to assemble a large sample of spectroscopic data for clusters of galaxies to construct samples of confirmed cluster members. Probably both of a primordial and an environmental effects would have influenced the clustering properties of galaxies in clusters. It is desirable to theoretically evaluate the LS due to each effect quantitatively. In addition, the observational investigation of the LS in high-z clusters would directly distinguish the effects, although morphological classification and field correction would become more difficult in such clusters. | 98 | 3 | astro-ph9803001_arXiv.txt |
9803 | astro-ph9803147_arXiv.txt | We report on observations of $^{12}$\coa emission from the chemically young starburst galaxy Mrk~109. These observations were part of a study to constrain the O$_2$/CO ratio in metal--deficient galaxies, which were motivated by theoretical work that suggests the possible enhancement of the O$_2$/CO ratio in chemically young systems. Five low metallicity ($Z \leq 0.5 Z_{\odot}$) IRAS galaxies at redshifts $z > 0.02$ (required to shift the 118.75 GHz $^{16}$O$_2$ line away from the atmospheric line) were searched for CO emission. We detected the CO line in only Mrk~109. From O$_2$ observations of Mrk~109, we achieved an upper limit for the O$_2$ column density abundance ratio of $N(\oo)/N(\co) < 0.31$. These results provide useful constraints for the theoretical models of chemically young galaxies. We argue that either most of the molecular gas in Mrk~109 does not reside in dark clouds ($A_{V} \ga 5$), or the standard equilibrium chemistry models are inadequate for metal--poor systems. The molecular gas mass implied by the CO observations of Mrk~109 is $M({\rm H}_2) \simeq 4 \times 10^{9} \msun$, and the CO data are consistent with a central starburst induced by the interaction with a nearby companion. | Although studying the molecular gas content of metal--deficient galaxies is challenging, such efforts are essential for our understanding of the formation and evolution of galaxies. By studying chemically young galaxies in the local universe, we gain insight into the processes which occurred at early times ($z\sim 1-5$) for the metal--rich spiral and elliptical galaxies found at the current epoch. Unfortunately, chemically young galaxies tend to be dwarf galaxies and are difficult to detect in CO (Combes 1985; Arnault et al. 1988; Sage et al. 1992; Israel, Tacconi, \& Baas 1995). Due to the observational difficulties, very little is currently understood about the molecular gas content of chemically young galaxies. One of the main uncertainties is the CO to H$_2$ conversion factor. Wilson (1995) and Arimoto, Sofue, \& Tsujimoto (1996) have found an empirical relationship indicating an increase in the CO to H$_2$ conversion factor with decreasing metallicity. However, detailed studies of the molecular clouds in the LMC and SMC suggest other factors, besides metallicity, have an important role on the CO to H$_2$ conversion factor (Rubio 1997; Israel 1997). Although significant questions remain in our understanding of the CO to H$_2$ factor, the molecular chemistry that occurs in metal--deficient galaxies is even more uncertain. Molecular chemistry calculations are quite complicated and typically involve networks of hundreds to several thousand reactions (Graedel, Langer, \& Frerking 1982; Herbst \& Leung 1989 [HL89]; Langer \& Graedel 1989 [LG89]; Bergin, Langer, \& Goldsmith 1995 [BLG95]). Since the chemistry models still have difficulties in explaining several of the observed abundance ratios in Galactic molecular clouds, very little modeling has been devoted to metal--poor galaxies. One notable exception is the theoretical study of the chemistry in LMC and SMC molecular clouds by Millar \& Herbst (1990) [MH90]. In general, the molecular chemistry models are sensitive to a variety of parameters, such as the density, temperature, and ionization field, but the dominant parameter for determining the relative molecular abundances of the carbon and oxygen species is the carbon to oxygen ratio (LG89). The models predict that the O$_2$/CO ratio decreases exponentially with increases in the C/O ratio (LG89). This could have interesting ramifications on the molecular abundances in chemically young galaxies. Observations with the Hubble Space Telescope have shown that the C/O abundance ratio increases with increasing metallicity in {\sc Hii} galaxies (Garnett et al. 1995). Therefore, we could expect to find lower C/O ratios and correspondingly larger O$_2$/CO ratios within dark molecular clouds in chemically young galaxies. Based on these simple ideas, the evolution of the O$_2$/CO abundance ratio as a function of metallicity in galaxies has been recently quantified for a variety of conditions and parameters governing the IMF and star formation histories (Frayer \& Brown 1997 [FB97]). At low metallicities, FB97 calculate lower C/O ratios and enhanced O$_2$/CO ratios (O$_2$/CO $\sim 1$) within dark ($A_{V} > 5$) molecular clouds. At solar metallicities and above, the O$_2$/CO ratio is expected to decrease by several orders of magnitude. Molecular oxygen has yet to be detected conclusively outside our solar system. Due to atmospheric attenuation, the ground--based Galactic searches have been limited to observing the rarer $^{16}$O$^{18}$O isotope (Liszt \& Vanden Bout 1985 [LV85]; Goldsmith et al. 1985; Combes et al. 1991 [C91]; Fuente et al. 1993; Mar\'{e}chal et al. 1997 [M97]). The most sensitive Galactic studies have provided upper limits of approximately O$_2$/CO $<0.1$. Extragalactic searches for the redshifted $^{16}$O$_2$ lines have provided more sensitive limits (O$_2$/CO $<0.01$) but have also been unsuccessful (Liszt 1985, Goldsmith \& Young 1989 [GY89], C91, Liszt 1992 [L92]; Combes \& Wiklind 1995). The most sensitive limit to date is O$_2$/CO $<0.002$ ($1\sigma$) from an absorption line study toward the radio source B0218+357 (Combes, Wiklind, \& Nakai 1997). All of these previous searches for O$_2$ have concentrated on chemically rich systems, or molecular clouds of unknown metallicity. Since the theoretical models suggest the possible enhancement of O$_2$ in chemically young galaxies, we have carried out a search for O$_2$ in metal--deficient galaxies. From the literature, we have compiled a list of metal--poor ($Z\leq0.5Z_{\sun}$) IRAS galaxies with redshifts of $z>0.02$. There are only about 20 galaxies satisfying these constraints, which were primarily drawn from the samples of Salzer \& MacAlpine (1988) and Dultzin--Hacyan, Masegosa, \& Moles (1990). Unfortunately, most of these galaxies are relatively weak IRAS sources $I(100\micron) \la 1$~Jy, and none have reported CO detections. This is not terribly surprising since metal--deficient galaxies have lower amounts of dust. The metallicities for the galaxies in our sample were derived from their oxygen abundances calculated using standard {\sc Hii} region analysis techniques whenever the 4363\AA\,[{\sc Oiii}] line was observed (Osterbrock 1989, case B). For galaxies with no 4363\AA\,line, the techniques of McGaugh (1991) were used to estimate the oxygen abundance. In this initial study, we observed the strongest IRAS sources satisfying the metallicity and redshift constraints. | In this paper, we report the nondetection of \coa emission in four distant metal--deficient galaxies and the detection of CO emission in Mrk~109. The \coa observations of Mrk~109 imply the presence of $4 \times 10^{9} \msun$ of molecular gas within the central few kpc of Mrk~109. The observational data favor a recent interaction of Mrk~109 with a nearby companion which has induced a starburst in the central regions of Mrk~109. The low metallicity, high gas fraction, and the kinematics are consistent with a young starburst. Based on the competing photodissociation and chemistry effects, we expect the largest global O$_2$/CO ratios for systems with metallicities in the range of $0.1 < Z/Z_{\sun} < 0.5$ (FB97). Although the metallicity of Mrk~109 falls within this optimal range, we fail to detect O$_2$ emission in Mrk~109. The observed intensity upper limit is $I(\oo)/I(\co) < 0.15$. These observations provide the first constraint on the abundance of O$_2$ in a metal--deficient galaxy. We derive a column density abundance limit of $N(\oo)/N(\co)<0.31$ for Mrk~109. These results suggest that most of the molecular gas in Mrk~109 does not reside in dark clouds, and/or that the gas--phase steady--state chemistry models do not apply for Mrk~109. It would be challenging to significantly improve the O$_2$/CO limit for chemically young environments derived from this work with current ground--based instrumentation. Future observations with satellites, such as ODIN (Hjalmarson 1997), of the LMC and SMC should provide more useful limits. | 98 | 3 | astro-ph9803147_arXiv.txt |
9803 | astro-ph9803093_arXiv.txt | Geomagnetic effects distort the zenith angle distribution of sub--GeV and few--GeV atmospheric neutrinos, breaking the up--down symmetry that would be present in the absence of neutrino oscillations and without a geomagnetic field. The geomagnetic effects also produce a characteristic azimuthal dependence of the $\nu$--fluxes, related to the well known east--west effect, that should be detectable in neutrino experiments of sufficiently large mass. We discuss these effects quantitatively. Because the azimuthal dependence is in first order independent of any oscillation effect, it is a useful diagnostic tool for studying possible systematic effects in the search for neutrino oscillations. | 98 | 3 | astro-ph9803093_arXiv.txt |
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9803 | astro-ph9803020_arXiv.txt | We report the first detection of an inverse Compton X-ray emission, spatially correlated with a very steep spectrum radio source (VSSRS), 0038-096, without any detected optical counterpart, in cluster Abell~85. The ROSAT PSPC data and its multiscale wavelet analysis reveal a large scale (linear diameter of the order of 500 $h^{-1}_{50}$ kpc), diffuse X-ray component, in excess to the thermal bremsstrahlung, overlapping an equally large scale VSSRS. The primeval 3 K background photons, scattering off the relativistic electrons can produce the X-rays at the detected level. The inverse Compton flux is estimated to be $(6.5\pm 0.5)\times 10^{-13}$ erg~s$^{-1}$~cm$^{-2}$ in the 0.5--2.4~keV X-ray band. A new 327 MHz radio map is presented for the cluster field. The synchrotron emission flux is estimated to be $(6.6\pm 0.90) \times 10^{-14}$ erg s$^{-1}$ cm$^{-2}$ in the 10--100 MHz radio band. The positive detection of both radio and X-ray emission from a common ensemble of relativistic electrons leads to an estimate of $(0.95\pm 0.10) \times 10^{-6}$ G for the cluster-scale magnetic field strength. The estimated field is free of the `equipartition' conjecture, the distance, and the emission volume. Further, the radiative fluxes and the estimated magnetic field imply the presence of `relic' (radiative lifetime $\ga 10^{9}$ yr) relativistic electrons with Lorentz factors $\gamma \approx$ 700--1700, that would be a significant source of radio emission in the hitherto unexplored frequency range $\nu \approx$ 2--10 MHz. | Apart from dark-matter, the diffuse intracluster medium has the two main constituents: the bremsstrahlung emitting hot ($T \sim 10^{7-8}$ K), tenuous ($n_0\sim 10^{-(3-4)}$ cm$^{-3}$), and diffuse thermal gas, and the higher energy, relativistic particles (cosmic rays), with electrons emitting the magneto-bremsstrahlung (synchrotron) radiation. The presence of diffuse, large scale (0.2--1.0 Mpc) synchrotron sources with very steep spectra, is known in several clusters (e.g. Coma, Abell 2256, 2319, 85, etc.). The origin of these relativistic particles is currently not well understood but in view of the absence of any optical counterparts, these radio sources are believed to be the remnants (`relics' and `halos') of once active radio galaxies. These remnants are prevented from rapid fading from expansion by the thermal pressure of the surrounding intra cluster gas (Baldwin and Scott 1973). The main energy losses of their relativistic electrons come from synchrotron emission and the inverse Compton scattering of the 3K background radiation. Their presence imply that magnetic fields on similar large scales may exist in the intracluster space (see Feretti \& Giovannini 1996, and Kronberg 1994 for reviews). That this magnetic field on large scales, may be a general property of clusters of galaxies, is suggested by the Faraday rotation data on radio sources observed in the direction of several clusters (Kim et al. 1991, Kronberg 1994). Estimates of the field strengths, based on Faraday rotation (Kim et al. 1991) or the `equipartition' hypothesis (e.g. Feretti \& Giovannini 1996), both give a value $B \sim$ 0.5 to 1.0~$\mu$G. Currently, the origin and evolution of these fields are not well understood due to observational difficulties in estimating the magnetic fields with sufficient accuracy and to the small number of actual estimates that has been attempted so far for a few clusters only (Kronberg 1994). Another method of considerable merit to estimate the cluster scale magnetic field is the detection of co-spatial inverse Compton (IC) X-ray emission with the synchrotron emission plasma. The ubiquitous 3K microwave background photons, scattering off the relativistic electrons (the IC/3K process), should produce a diffuse X-ray `glow' associated with the radio plasma (Feenberg \& Primakoff 1948; Hoyle 1965; Baylis et al. 1967; Felten \& Morrison 1966; Harris \& Grindlay 1979; Rephaeli \& Gruber 1988). The non-relativistic (thermal) analogue of this process is the well known Sunyaev-Zel'dovich effect (Rephaeli 1995). A possible detection of inverse Compton X-ray photons from the scattering of 3K microwave background would not only provide a magnetic field estimate based on a different physical process, but would also provide information on as yet unknown non-thermal X-ray component in galaxy clusters. The inverse Compton method has several advantages: both the magnetic field strength, and the electron energy spectrum are obtainable from the observed IC/3K and the synchrotron emission fluxes, and the field so obtained is independent of the `equipartition' conjecture (Pacholkzyck 1970), the distance, and the emission volume. Despite a number of attempts over the last 20 years, the detection of IC/3K radiation has proved elusive, very often with only upper limits to the X-ray fluxes and lower limits to magnetic fields (e.g. Harris \& Romanishin 1974, Rephaeli 1977, Rephaeli et al. 1987, Rephaeli \& Gruber 1988, Bazzano et al. 1990, Rephaeli et al. 1994, Harris et al. 1995). The principal constraints arise from rarity of diffuse, steep spectrum cluster radio sources, limited sensitivity, and spatial/spectral resolution of low frequency (10--300 MHz) radio or X-ray telescopes, and confusion from cluster thermal emission. However, correlating the radio and the EINSTEIN satellite data, Bagchi (1992) suggested the possibility of IC/3K emission in Abell~85. Recently, using the ROSAT PSPC data, Laurent-Muehleisen et al. (1994) and Feigelson et al. (1995) claimed the detection of IC/3K X-ray emission associated with the radio lobes of the galaxy Fornax-A, further supported by Kaneda et al. (1995), employing the ASCA X-ray spectral data. This work employs the good spatial resolution, high sensitivity and spectral capabilities of the ROSAT X-ray satellite, to search for co-spatial IC/3K emission from the diffuse radio source 0038-096 located in the rich cD cluster Abell~85. We present evidence for what might be the first detection of the IC/3K emission from the `relic' relativistic electrons residing in the intra-cluster medium, presumably the remnants of an once active radio galaxy that is presently unidentified. We also present a new 327 MHz radio map of the cluster field. Combining the radio and X-ray data, we estimate the magnetic field strength in the co-spatial emission volume. We then briefly discuss its implications. For a cluster redshift of 0.0555 (Pislar et al. 1997), 1 arcmin corresponds to $97 h^{-1}_{50}$ kpc, with the Hubble constant expressed in units of 50~km~s$^{-1}$~Mpc$^{-1}$. | We have presented an evidence for correlated X-ray emission with the diffuse VSSRS in Abell~85. What are the possible emission processes (other than the IC/3K) that can give rise to the excess X-ray emission? The possibility that the emission is either the extension of the synchrotron radiation to the $\sim$1 keV X-ray band, or that arises from the synchrotron self Compton process (Rees 1967, Harris et al. 1994), can be ruled out from the observed very steep spectral shape for $\nu \ge $ 1 GHz, signifying a dearth of necessary high energy photons. Further possibilities that the emission is generated by an active galaxy or that it arises from radio-jet cluster-medium interaction also appear remote due to the relaxed, `halo' type morphology of the VSSRS, and the absence of any optical counterpart. Recent X-ray imaging and optical spectroscopy of clusters have shown the presence of thermal X-ray emission associated with secondary substructures possibly in the process of gravitational merger (e.g. Briel et al. 1991, Mohr et al. 1993, Henry and Briel 1993, Burns et al. 1994). Significantly, the presence of radio-halo sources appears to be correlated with the evidence for recent mergers, very high gas temperatures (7--14 keV), large velocity dispersions ($\approx 1300$~km~s$^{-1}$), and the absence of both cooling flows and a single dominant central galaxy (BM type II or III; Edge et al. 1992, Tribble 1993, Feretti and Giovannini 1996). It has been proposed that these data possibly indicate the heating of intracluster gas and the reaccelaration of cosmic ray particles with stochastic amplification of magnetic fields, powered by the energy released in the ongoing mergers (Tribble 1993). Can the excess X-ray emission seen with the VSSRS in A85 come from such a process? Although imaging data alone can not resolve this question, it is interesting to note that in terms of its physical properties, unlike the other radio-halo host clusters such as the Coma, Abell~2163, 2218, 2255,2256, and 2319, the cluster Abell~85 contains a central dominant cD galaxy, a central cooling flow of about 100 M$_{\odot}$~yr$^{-1}$, and relatively `cool' gas at a temperature of T $\sim$ 4 keV (Pislar et al. 1997; Lima Neto et al. 1997). Durret et al. (1998) have detected a filamentary structure, visible both in optical and in X-rays, linking Abell 85 to the neighbouring cluster Abell 87. They present evidence for a possible merger of matter in this filament with the southern region of Abell 85, in the `south-blob' region. Markevitch et al. (1998) detect an enhancement in the gas temperature in the same region, possibly related to the merger. The position of the VSSRS, however, is not located there as it is shifted about 3 arcmin (300$h_{50}^{-1}$ kpc) to the northwest from the supposed shocked region and has a cooler temperature for the thermal matter in its vicinity. The resolution of the question of exact physical process behind the excess X-ray emission would require high spectral (and spatial) resolution imaging spectroscopy data. In the hard X-ray energy regime of $\approx$ 10--20 keV, the contribution from thermal bremsstrahlung would drop sharply, whereas the non-thermal inverse Compton flux would be visible as an extra component with a steep power law spectrum of spectral index $\alpha \approx$ 1. Based on the data presented in this work, we predict the possible IC/3K flux of $\approx 2.8 \times 10^{-13}$ erg s$^{-1}$ cm$^{-2}$, and the thermal bremsstrahlung flux of $\approx 3.0 \times 10^{-13}$ erg s$^{-1}$ cm$^{-2}$ in the hard X-ray band of 5--10 keV. These fluxes are well within the reach of the currently operative BeppoSAX telescope and the upcoming AXAF and the XMM missions. The nature of the other excess X-ray and the steep spectrum radio emission detected over the `south-blob' (cf. below) could possibly also be understood with the spectroscopic X-ray data. With the evidence currently available to us, energetically the most feasible mechanism for the X-ray emission from the VSSRS appears to be the IC/3K process. From theory (e.g. Ginzburg 1989), the (total) energy loss ratio can be shown to be ${f_{IC} \over f_{S}} \approx U_{rad}/ {B^{2} \over {8 \pi}}$, where $U_{rad} \approx 5 \times 10^{-13}$ erg~cm$^{-3}$ is the energy density of the cosmological black-body radiation field, and ${B^{2} \over {8 \pi }} \approx 4 \times 10^{-14}$ erg~cm$^{-3}$ is the observed magnetic energy density. Therefore, if the observed radiation is from IC/3K process, we expect $f_{IC} / f_{S} \approx 13$, which is comparable to the observed ratio $f_{IC} / f_{S} \approx 10$ (the difference is mainly attributed to the finite bandwidths of our data). If the excess X-rays are produced by the thermal bremsstrahlung process, it is difficult to understand why the observed ratio is comparable to the ratio expected if IC/3K were the emission mechanism. This again suggests that the detected X-rays are indeed a product of inverse Compton scattering. This data on ratio of fluxes has enabled us to obtain a model independent magnetic field value, $B = 0.95 \pm 0.10 \mu$G, for the diffuse emission volume, located at $\sim 700 h_{50}^{-1}$ kpc from the cluster centre. It is apparent (Fig. 2) that another diffuse X-ray excess (the `south-blob') is located at $\alpha = 00^{\rm h}~41.8^{\rm m}$, $\delta = -09^{\rm d}~27.5^{\rm m}$. In the same region, two radio components, without any definite optical counterparts, are observed with the OSRT with $> 20 \sigma$ signal (162 mJy: northern source, 148 mJy: southern source). A spectral index limit, $\alpha \ > 2.3$, is obtained for each of them, based on their non-detection in the new VLA 1.4 GHz sky-survey (Condon et al. in preparation) to the $\approx$ 5 mJy flux density limit. The nature of this excess X-ray emission is currently not well understood (Lima Neto et al. 1997). Is it possible that a second IC/3K X-ray source is located here, co-spatial with yet another relic radio plasma? If true, this can provide another magnetic field estimate at $\approx 1 h_{50}^{-1}$ Mpc from the cluster centre. Detailed, low-frequency ($<327$ MHz) radio and X-ray spectral data are necessary to explore this possibility. Finally, the estimated magnetic field and the observed radiative fluxes from the VSSRS imply the presence of relativistic electrons with Lorentz factors $\gamma \approx$ 700--1700 (for 0.5--2.4 keV band, $\nu_{x} \propto \nu_{bg} \gamma^{2} $). The population of electrons in this energy range would in turn produce significant radio emission ($\sim $ 1400 Jy at 2 MHz) in the presently unexplored frequency range $\nu_{syn} \approx $ 2--10 MHz ($\nu_{syn} \propto B\;\gamma^{2}$). The radiative lifetime of such `relic' electrons is $\ga 10^{9}$ yr (Harris \& Grindlay 1979), comparable to the time-scale for evolution of clusters. Although this frequency range is below the ionospheric cutoff, the future generation radio telescopes observing from above the earth's atmosphere or from the far side of the moon (Burns 1990), could make such measurements. Such low frequency radio data in association with sensitive X-ray data would prove invaluable in probing the physics of the intracluster media in large scale structures. It is encouraging to note that the inverse Compton X-ray technique has the great potential, not only in measuring the elusive intra-cluster magnetic fields, but also in probing the hitherto unexplored population of very old relativistic electrons residing in the diffuse intra-cluster space. | 98 | 3 | astro-ph9803020_arXiv.txt |
9803 | astro-ph9803216_arXiv.txt | \d\/ in NGC 1536 is possibly the least luminous and energetic Type II supernova discovered to date. The entire light curve is subluminous, never reaching $M_V = -14.65$. The radioactive tail follows the \co\ decay slope. In the case of nearly complete trapping of the $\gamma$-rays, the \ni\ mass derived from the tail brightness is extremely small, $\sim 0.002$ \M. At discovery the spectra showed a red continuum and line velocities of the order of 1000 \kms. The luminosity and the photospheric expansion velocity suggest that the explosion occurred about 50 days before discovery, and that a plateau probably followed. Model light curves and spectra of the explosion of a 26 \M\ star successfully fit the observations. Low mass models are inconsistent with the observations. The radius of the progenitor, constrained by the prediscovery upper limits, is \r0\ \ltsim 300 \R. A low explosion energy of $\sim 4 \times 10^{50}$ ergs is then required in the modeling. The strong \ion{Ba}{2} lines in the photospheric spectra are reproduced with a solar abundance and low $T_{\rm eff}$. A scenario in which the low \ni\/ mass observed in \d\/ is due to fall--back of material onto the collapsed remnant of the explosion of a 25--40 \M\/ star appears to be favored over the case of the explosion of an 8--10 \M\ star with low \ni\/ production. | \d\/ was serendipitously discovered on Jan. 14.15 U.T. (\cite{demello}) in NGC~1536, a morphologically disturbed spiral galaxy belonging to a high density group (\cite{maia}). Although NGC~1536 ($v_{\rm helio}=1296$ \kms, RC3) is one of the galaxies patrolled visually by Rev. Evans for his SN search (private communication), he missed the SN because its brightness apparently never exceeded his detection limit. The first spectrum showed the main features of type II SNe, but also revealed peculiarities which make this object unique. Particularly noteworthy were the extremely slow expansion velocities, the red color and the strong \ion{Ba}{2} lines (\cite{tur98}). A campaign of photometric and spectroscopic observations was therefore promptly started at ESO and CTIO. The complete data set will be presented elsewhere. | \d\/ is among the least luminous Type II supernovae known to date. The low luminosity persists also in the tail, where the decline rate is close to that of a \co-powered tail. \d\ is also characterized by a low temperature and a very low expansion velocity. These features are well reproduced by an explosion model with a relatively small progenitor radius (300 \R), massive ejecta ($\sim$ 24 \M), low explosion energy ($4 \times 10^{50}$ ergs), and small ejected \ni\/ mass ($\sim$ 0.002 \M). A similarly small amount of \ni\ has recently been proposed for another SN~II, SN~1994W (\cite{soll}), but this object was rather luminous and its tail declined faster than the \co\ slope. The neutrino-heating mechanism of massive star explosions is not well understood, so the explosion energy and the amount of \ni\ ejected as a function of progenitor mass are difficult to predict from the hydrodynamical models. According to pre-supernova model calculations, stars more massive than $\sim$ 25 \M\ form an Fe core in a greater gravitational potential because of the significantly smaller C/O ratio and the consequently weak carbon shell burning (\cite{nom93}; \cite{woowae}). Therefore, if the efficiency of neutrino heating does not change with exploding mass, more massive stars tend to produce lower explosion energies, and they suffer from more fall-back of material onto the collapsed remnant. The mass cut is then placed further out. However, the progenitor must be less massive than the Wolf-Rayet progenitors (i.e., \ltsim 30 - 40 \M) because of the presence of an envelope with a mass of at least 15 \M\ above the photosphere at day 50. Stars in the range $\sim$ 10 - 25 \M\ are likely to explode with typical explosion energies of $\sim 1 \times 10^{51}$ ergs, and to produce 0.07 - 0.15 \M\ of \ni, as indicated by the brightness and light curve shapes of SN 1993J, SN 1987A and of Type Ibc supernovae (\cite{nom95}; \cite{young}). Stars of 8 - 10 \M\/ are predicted to produce little \ni\ and other heavy elements (\cite{nom84}; \cite{maywil}). However, these stars are at the top of the AGB in the pre-supernova stages (\cite{hashi}), which is inconsistent with the small radius of \d. Also, synthetic spectra obtained from the low mass models have shallow absorption lines compared to the observations. In conclusion, the progenitor of \d\ was probably as massive as 25 - 40 \M. Our 26 \M\ model is in this mass range, thus being consistent with the low $E/M$ and the small \ni\ mass. The presupernova radius depends on many parameters, and the reason for its being rather small in \d\/ is an open question. Spectral models appear to rule out the case of low metallicity and He enhancement. A possible scenario involves a close binary system in which the companion star spiraled in to make the envelope more massive (\cite{pod}). The ejection of a small mass of \ni\ provides a constraint on the mass of the collapsed remnant, $M_{\rm rem}$ (\cite{nom93}; \cite{thiel}). In our 26 \M\ model, is $M_{\rm rem} = 1.8$ \M. If $M_{\rm rem}$ were assumed to be smaller, more \ni\ would be ejected: e.g. if $M_{\rm rem} = 1.7$ \M, then M(\ni) = 0.1 \M. For more massive progenitor models, $M_{\rm rem}$ would be larger than 1.8 \M. If the equation of state of nuclear matter is relatively soft, the maximum (gravitational) mass of a cold neutron star is $\sim$ 1.5 \M\ (\cite{bb94}). If this is the case, the remnant of \d\ must be a small-mass black hole, and the low luminosity tail of \d\ might be a signature of the formation of such an object. \bigskip We would like to thank Piero Rosati for observing \d\/ and Rev. Robert Evans for providing his pre-discovery limits. P.A. Mazzali acknowledges receipt of a Foreign Research Fellowship at N.A.O., and is grateful to T. Kajino and the Dept. of Astronomy at the University of Tokyo for the hospitality. A. Piemonte aknowledges financial support from ESO during his stay in Chile. This work has been supported in part by the Grant-in-Aid for Scientific Research (05242102, 06233101) and COE research (07CE2002) of the Ministry of Education, Science, and Culture in Japan, and by the National Science Foundation in US under Grant No. PHY94-07194. Part of the observations have been obtained at the ESO 1.5m telescope operated under the agreement between ESO and Observatorio Nacional -- Brasil. \noindent | 98 | 3 | astro-ph9803216_arXiv.txt |
9803 | astro-ph9803166_arXiv.txt | The heavy elements present in the hot gas in galaxies are thought to be supplied both by stellar winds and by supernovae (SNe). Theoretical calculations of the nucleosynthesis of the SNe have been improved greatly in the past ten years (e.g., Thielemann, Nomoto, \& Hashimoto 1990; Thielemann, Nomoto, \& Hashimoto 1996, hereafter TNH). In contrast, the observational evidences of metal rich gas have been obtained for only a limited number of young supernova remnants (SNRs). Ku et al. (1984) first performed spatially-resolved spectroscopy on the Cygnus Loop with the Einstein observatory. They constructed an X-ray image with spatial resolution of $\sim$1$^\prime$ in the energy region of 0.1--4.0 keV. The X-ray image clearly showed limb-brightening structure. They noticed that kTe at the limb was generally lower than that at the center. Charles, Kahn, \& McKee (1985) found that kTe gradually increased toward the center from the vicinity of the limb. Tsunemi et al. (1988) observed the whole Cygnus Loop with the Gas Scintillation Proportional Counters on board the Tenma satellite, which possessed roughly twice better energy resolution than that of the proportional counter (Koyama et al. 1984). They detected emission lines from highly ionized Si and S. They fitted the X-ray spectrum in combination with the spectrum obtained with a previous rocket observation (Inoue et al. 1979, 1980). They found that a two-component NEI model could reproduce the X-ray spectra. Together with Charles et al. (1985), Tsunemi et al. (1988) suggested that tenuous high-kTe plasma was sitting in the interior region which was surrounded by dense low-kTe plasma. Such high-kTe component was confirmed by Hatsukade \& Tsunemi (1990) with the Ginga satellite. Due to its high sensitivity, high energy resolution, and wide energy band, the ASCA observatory has opened a new window for the physics of hot plasmas like SNRs. We can perform the plasma diagnostics in detail by using the X-ray CCD cameras. So far, the abundances of heavy elements have been well determined for many young SNRs with the ASCA observatory (see review by Tsunemi \& Miyata 1997). For evolved SNRs, it is difficult to determine the abundances due to low-kTe, low surface brightness, and large interstellar absorption. Furthermore, the swept up interstellar medium (ISM) dominates the ejecta mass from the progenitor star. This makes it difficult to detect the ejecta. The Cygnus Loop is an evolved SNR and one of the best studied SNRs. Since the Cygnus Loop is roughly 8 degrees away from the Galactic plane, the neutral hydrogen column (${\rm N_H}$) is only a few$\times {10}^{20}\ {\rm cm}^{-2}$. This enables us to detect emission lines from O and to determine kTe with high accuracy. The mass of the swept up ISM is estimated to be roughly 100 ${\rm M_{\odot}}$ since the average density of the ISM is 0.2 ${\rm cm}^{-3}$ and the mean radius is 18.8 pc (Ku et al. 1984). In comparison, the mass of the ejecta ranges from one to several tens ${\rm M_{\odot}}$. Thus, the ejecta are submerged under the sea of the ISM. However, due to its high surface brightness and large apparent size, we can perform spatially-resolved analysis for the Cygnus Loop. This enables us to search ejecta in detail if ejecta have not yet mixed well with the ISM. The first observation of the Cygnus Loop with the ASCA observatory was performed by Miyata et al. (1994; hereafter MTPK) on the NE limb. They found that kTe increased toward the center whereas log\hspace{0.2mm}($\tau$) decreased. They determined the abundances of heavy elements from O to Fe and found that metals were deficient at the NE limb. In this {\it paper}, we present the observational results of the center portion of the Cygnus Loop observed with the ASCA observatory. The X-ray spectrum obtained from the center portion was quite different from that obtained at the NE limb. We studied the plasma diagnostics of the interior of the Cygnus Loop and nucleosynthesis at the SN, apart from dense shell regions. | We observed the center portion of the Cygnus Loop supernova remnant with the X-ray CCD cameras on board the ASCA observatory. We confirmed a significant departure from a collisional ionization equilibrium condition at the center portion. Obtained abundances at the core region were larger than those of cosmic abundances for Si, S. Moreover, abundances of Si, S, and Fe were $\sim$40 times larger than those obtained at the NE limb. This strongly supports the hypothesis that the X-ray emitting plasma in the core is ejecta in origin. Although previous X-ray observations could not detect any signature of ejecta, we could find it by using both the imaging capability and the high energy resolving power of the ASCA observatory. Obtained abundances can be compared with theoretical calculations of nucleosynthesis by a SN explosion. Since we only observed the core region, we integrated the model calculations from the explosion center to some mass radius. For the explosion models for type Ia SN or for type II with the progenitor mass less than 20 ${\rm M_{\odot}}$, the abundance of Fe is much larger than that of Si or S even if we integrate to the outermost radius. On the other hand, the abundances of Si, S, and Fe can be explained with the type II model with 25${\rm M_{\odot}}$ if we integrate these abundances from the mass cut radius 1.77${\rm M_{\odot}}$ to $2 \sim 3.9 {\rm M_{\odot}}$. Therefore, our results suggest the massive progenitor origin of the Cygnus Loop and the existence of the stellar remnant associated with the Cygnus Loop. \par \vspace{1pc}\par We would like to acknowledge Prof. Nomoto and Drs. T. Suzuki and K. Iwamoto for fruitful discussions about the theoretical aspects of the SNRs at conference of Thermonuclear Supernovae. Dr. B. Aschenbach kindly gave us the whole X-ray image of the Cygnus Loop obtained with the ROSAT all-sky survey. We thank Dr. Thielemann to use his data on nucleosynthesis of SNe. We would like to thank the anonymous referee for her or his detailed comments and suggestions which greatly improved this paper. We are greatful to all the members of ASCA team for their contributions to the fabrication of the apparatus, the operation of ASCA, and the data acquisition. We thank the members of {\tt ASCA\_ANL/Sim ASCA} software team. | 98 | 3 | astro-ph9803166_arXiv.txt |
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9803 | astro-ph9803344_arXiv.txt | We present X-ray and optical data of the new\-ly discovered AM Her variable \rxj18. This X-ray source was observed in the \ros\ All-Sky-Survey in September 1990 and subsequently discovered as a highly variable and soft X-ray source. Follow-up pointed observations confirmed the strong variability and revealed a periodic flux modulation including a sharp and nearly complete eclipse of the X-ray emission. Based on the shape and duration of this eclipse as well as the lack of optical eclipses we favour an interpretation in terms of a self-eclipse by the accretion stream. The X-ray spectrum averaged over the full period is dominated by soft emission below 0.5 keV. The ratio of soft (blackbody) to hard (bremsstrahlung) bolometric energy flux is 1670, distinguishing this object as another example of a polar with an extreme strong excess in soft X-rays. | AM Her type variables are accreting binaries in which the strong magnetic field of the white dwarf controls the geometry of the material flow coming from the late-type companion star. The radial inflow of matter produces a shock front above the white dwarf surface which gives rise to hard X-rays and polarized cyclotron radiation in the IR to UV range. In the simplest physical model (Lamb 1985) half of this shock radiation intercepts the white dwarf surface and is reradiated as thermal emission from the heated accretion spot in the UV and soft X-ray band. X-ray observations, most notably with the \ros\ satellite, had and have a great impact on the study of magnetic cataclysmic variables. The rather high intensity, strongly variable and soft X-ray emission allows to set up selection criteria (e.g. using the \ros\ data base) with a high detection/identification rate which reflects in the fact that about 80\% of all magnetic cataclysmic variables have been discovered at X-ray wavelengths (Beuermann 1998). During the search for supersoft X-ray sources in the \ros\ (Tr\"umper 1983) All-Sky-Survey (RASS), we have found a bright X-ray source showing strong scan-to-scan variability in the X-ray count rate as measured in the PSPC. Subsequent spectroscopic observations re\-vealed a 15th magnitude cataclysmic variable located near the X-ray position. Additional pointed \ros\ PSPC observations strengthened the evidence of a strongly modulated X-ray light curve as well as of a very soft X-ray spectrum. These combined characteristics, and particularly the supersoft X-ray spectrum, are typical signatures of the AM Her subclass (Beuermann \& Schwope 1994). Here, we report on the results of our extensive \ros\ and optical observations performed to study the accretion geometry in \rxj18\ in more detail. The identification as an AM Her system (polar) and some first details have been reported in June 1994 at the Abano-Padova conference on Cataclysmic Variables (Greiner \etal\ 1995a) after which the object got the variable star name V884 Her. The results of quick follow-up observations of \rxj18\, (= WGA J1802.1+1804) have also been communicated in Singh \etal\, 1995 (note the flux error; see Greiner \etal\ 1995b) and Szkody \etal\ (1995). Recently, Shrader \etal\, (1997) reported the results of IUE observations. Due to its soft X-ray spectrum \rxj18\ has also been detected with the EUVE satellite (Lampton \etal\ 1997; = EUVE J1802+18.0). | \subsection{AM Her type classification} The coincidence of periods derived from the X-ray and optical data and the stability of the eclipse feature over more than 5 years suggest that the binary system is synchronized. The evidence of an X-ray eclipse without a corresponding optical eclipse together with the shape and variations with time of the X-ray eclipse imply that (1) the optical and X-ray radiation originate from different emission regions and (2) the X-ray emitting region is confined to a small region near the white dwarf. The localized X-ray emitting region strongly indicates a magnetic CV subclass, since the X-ray emission from an inner accretion disk would only show eclipses when the disk is occulted by the companion star -- and such eclipses are evident in both X-ray and optical light. The supersoft X-ray spectrum further suggests an AM Her / `polar' subtype (rather than `intermediate polar' type), since the association between AM Her objects and luminous emission in soft X-rays is well established. We note, however, that the discovery of three intermediate polars with strong soft X-ray emission has changed this clear association (Haberl and Motch 1995). However, the consistency of singular optical and X-ray periods in the case of \rxj18 indicates the likelihood of spin / orbital synchronous rotation, another classic property of AM Her binaries. Finally, the large equivalent widths of the optical emission lines also suggest a polar nature. An independent confirmation of the polar classification comes from polarimetric observations (Szkody \etal\ 1995). The circular polarization in the 5700--7700 \AA\ range reaches a maximum of 4\% in two peaks separated by a minimum which occurs at the time of photometric minimum. \subsection{Geometric configuration} The relatively short time scales observed for X-ray eclipse ingress, egress, and duration argue against the interpretation of eclipses as the disappearance of the X-ray emitting region behind the limb of the rotating white dwarf. We therefore interpret the eclipse as a self-eclipse of the X-ray emitting region by the accretion stream. X-ray spectral evidence for increased absorption column associated with the eclipse and the variable eclipse length provide further support for this conclusion. Although one would expect the optical light curve to be more complex than the X-ray light curve due to several reasons (cyclotron radiation is beam-modulated and can be self-eclipsing; the recombination component is modulated by projection of the partially optically thick stream; additional light may be seen from the illuminated side of the secondary), the observed optical light curve is surprisingly smooth. Despite the stream-eclipse of the accretion spot in X-rays, the optical light is not affected due to its probable origin from a more extended region, including a portion of the accretion column upstream from the X-ray emitting area. The lack of optical eclipses suggests an inclination of $i < 78$\degr\ (assuming a typical M5/6 companion). During mid eclipse, there is still detectable emission not only in the total flux, but also in the soft X-ray component. Thus, it seems possible that the accretion spot is not fully eclipsed or we see another, uneclipsed emission component. The presence of X-ray emission at all phases implies that $i + \beta <$ 90\degr\ (with $\beta$ being the colatitude of the accreting magnetic pole). Combining this with the condition of stream eclipse ($i > \beta$) we derive the following constraints on $i$ and $\beta$: $$ 45\degr < i < 78\degr\ \hspace{0.5cm} {\rm and} \hspace{0.5cm} 0\degr < \beta < 45\degr $$ The X-ray light curve is very complex, and there are large differences seen when comparing observations well separated in time. This might suggest that the density and/or the size of the accretion stream vary with time and thus cause changes in the size and/or vertical extent of the accretion spot. In this regard, we note that the eclipse duration is not constant. Also, there are large changes in the characteristics of a secondary minimum, as noted in Section 2.1.2 and shown in Fig. \ref{foldlc3}. Furthermore, the eclipse length seems to be correlated with the intensity of the X-ray emission before and after the eclipse. During the Oct. 1992 observation (cycle --4269 in Fig. \ref{unfoldlc}) with its X-ray low-state, the eclipse length is only 0.07 phase units, whereas it is 0.1 during the high-state in Sep. 1993 and even 0.12 in the last eclipse observed. A higher X-ray intensity may be caused by an increase in the rate of mass transfer, perhaps leading to a larger width of the accretion stream, which then increases the duration of the X-ray eclipse. A significant larger ingress (5--7 min.) than egress (2--2.5 min.) is measured. Such asymmetry is possible if the impact area is not a circle but an arc due to the coupling of the stream to nonpolar field lines (see e.g. Beuermann \etal\ 1987 and their notion of ``X-ray auroral oval''. Alternatively, the accretion stream can be imagined to impact the white dwarf surface not from a perpendicular direction thus forming an elongated footprint on the white dwarf surface. The slow fall of X-ray intensity into the eclipse minimum then can be understood as being due to two effects: first, the projected area of the footprint decreases due to the rotation of the white dwarf and secondly, the stream starts to occult the footprint area. This is consistent with the observed X-ray spectral property that the absorbing column does not change during the start of the slow fall into eclipse minimum (phase interval 0.0--0.1, see Fig. \ref{foldlc}), but only during the eclipse (phase interval 0.1--0.2). Such behaviour is also consistent with the interpretation of the variation of the polarization being due to varying aspects of an emission region extended in longitude (Szkody \etal\ 1995). A size estimate of this elongated emitting region beyond an axis ratio of about 2--3:1 is hard to evaluate since the relative sizes of the emitting region and the stream are important. The X-ray eclipse occurs about 35\degr\ after the blue-to-red crossing of the radial velocity curve. In our interpretation of a stream-eclipsing geometry this is consistent with the picture that the emission lines are produced in the ballistic part of the accretion stream. We note, however, that the radial velocity amplitude is lower than one would expect as the maximum free fall velocity. The best fit blackbody temperature and the normalization transform into a corresponding blackbody radius of the emission region of 750 (D/100 pc) km. Bearing in mind the above mentioned elongation this is understood as the radius of a circle which has the same area as the elongated emitting region. If true, this implies a stream diameter near the white dwarf surface of the order of 850--1050 km. Besides the herewith presented \rxj18\ several other polars have been found where the accretion stream was identified as the cause of narrow absorption dips in the EUV/X-ray light curves: EF Eri (Patterson \etal\ 1981, Beuermann \etal\ 1991), UZ For (Warren, Vallerga and Sirk 1995), EK UMa (Clayton \& Osborne 1994), QS Tel (Beuermann \& Thomas 1993, Buckley \etal\ 1993, Schwope \etal\ 1995), HU Aqr (Schwope \etal\ 1993, Hakala \etal\ 1993, Schwope \etal\ 1997), QQ Vul (Beardmore \etal\ 1995) and V2301 Oph (= 1H1752+081) (Hessman \etal\ 1997). Another stream-eclipsing polar may be AX\,J2315--592 which has recently been found from ASCA observations (Misaki \etal\ 1996, Thomas and Reinsch 1996). The EUVE light curves of UZ For show that both the phase and the amplitude of the dips vary over the 3-day observation suggesting variations in the density and position of the accretion stream similar to what we observed in the \ros\ light curves of \rxj18. Also, the X-ray light curve of QQ Vul as observed with ROSAT is surprisingly similar to that of \rxj18\ (as given in Fig. \ref{foldlc}). \subsection{Soft X-ray excess} As in several other AM Her binaries, the soft X-ray luminosity of \rxj18 is larger than the hard one, a fact known as `soft X-ray problem' (Rothschild \etal\ 1981). Kujpers \& Pringle (1982) proposed that non-stationary accretion of dense blobs can heat the photosphere from below. This implies that high magnetic field systems should have a weak bremsstrahlung component, which is proved by a comparison of soft-to-hard flux ratios with measured magnetic field strengths in different AM Her systems (Beuermann \& Schwope 1994). According to this correlation a magnetic field larger than 40 MG is suggested for \rxj18\ which has one of the highest ratios of soft/hard emission among polars. This would produce a very impressive Zeeman split in the H absorption lines coming from the white dwarf as well as strong cyclotron humps (depending on the viewing angle). However, since we observed the system ``only'' in its robust state of accretion, the optical spectrum is strongly dominated by the accretion stream making the determination of the magnetic field impossible. | 98 | 3 | astro-ph9803344_arXiv.txt |
9803 | gr-qc9803073_arXiv.txt | In the hyperbolic slicing of de Sitter space appropriate for open universe models, a curvature scale is present and supercurvature fluctuations are possible. In some cases, the expansion of a scalar field in the Bunch-Davies vacuum includes supercurvature modes, as shown by Sasaki, Tanaka and Yamamoto. We express the normalizable vacuum supercurvature modes for a massless scalar field in terms of the basis modes for the spatially-flat slicing of de Sitter space. | Scalar fields in de Sitter spacetime have long provided a testing ground for issues of quantum field theory in curved spacetime \cite{birdavies,sfulling}. Further motivation for their study stems from the central role they play in inflationary cosmology \cite{KolbTurner}. Several different coordinate systems can be used to cover de Sitter space, and subtleties in the quantization of fields can arise in some of the less familiar coordinate systems. These subtleties have been highlighted by recent models of open inflation \cite{bgt,openi}, in which two periods of inflation are separated by nucleation of a bubble. The bubble interior includes an open universe ($\Omega_0 <1$), where we could be living today, described by hyperbolic, spatially-curved coordinates \cite{colegott}. A key difference between these hyperbolic coordinates and the more familiar spatially-flat slicing of de Sitter space is the presence of a curvature scale. This in turn leads to the possibility of supercurvature \cite{lythwos} fluctuations, fluctuations with wavelength longer than the curvature scale. Unlike the continuum of modes familiar from the spatially-flat slicing of de Sitter space, a normalizable supercurvature mode may exist for an isolated, discrete eigenvalue of the spatial Laplacian, or not at all. Although such modes have no analogue in the spatially-flat slicing of de Sitter space, it has been shown \cite{sasaki95} (see also \cite{Mosch}) that the supercurvature modes must be included in the vacuum spectra of low-mass scalar fields in order to produce a complete set of states, and hence the proper Wightman function. For massless, minimally-coupled scalar fields in de Sitter space, there is in addition a well-known infrared divergence in the (coordinate-independent) Wightman function. The infrared divergence is related to a dynamical zero mode in the spectrum of a massless field \cite{desinfra,kg93}. Kirsten and Garriga \cite{kg93} covariantly quantized this zero mode in a spatially-closed slicing of de Sitter space. Extending their result from these closed coordinates to the coordinate system appropriate to open inflation has not yet been done. Here we identify the zero mode as one of the supercurvature modes when quantizing the minimally-coupled massless field in the open hyperbolic coordinates. It has already been noted that supercurvature modes can make significant contributions to the fluctuations in the cosmic microwave background (CMB) radiation, and several of their effects have been calculated for models of open inflation \cite{sccmb,garriga,sasaki96,gar-97}. (The massless field zero-mode subtleties, except for a variant studied in \cite{gar-97}, do not arise for these specific CMB calculations, which are sensitive to higher multipoles.) In addition to contributing to observable density fluctuations, such long wavelength, supercurvature modes might play a role \cite{dkpre} at the end of inflation, when recent advances \cite{newreh} in the theory of reheating are taken into consideration. \indent In summary, increased understanding of these supercurvature modes is motivated both by general questions of quantizing fields in curved backgrounds, and by recent inflationary model-building. We will focus here on the example of supercurvature modes for a massless, minimally-coupled scalar field, expressing them as a sum over the basis modes for a spatially-flat slicing of de Sitter space. This overlap gives a measure of \lq\lq where all the supercurvature modes go" in the familiar flat-space spectrum of such fields. The massless case is chosen for tractability, and questions about the zero mode are postponed for future work \cite{jdcdk}. Throughout this paper, we consider only an unperturbed de Sitter metric; a more complete treatment would include study of the backreaction of such fields on the background metric. Because the supercurvature modes stretch beyond the horizon, any such study of the coupled metric fluctuations would need to pay special attention to the gauge subtleties which always accompany superhorizon fluctuations \cite{rhb}, and such issues are not pursued here. In section II, the two covers of de Sitter space (including the flat and hyperbolic slicings) are given, and the field quantization pertinent to open inflation in both systems is reviewed. As these supercurvature modes have no analogue in the usual flat slicing of de Sitter space, this section gathers some previous work on supercurvature modes and provides notation and context for the rest of the paper. Section III specializes to the massless case. The explicit calculation of the overlap for some of these supercurvature modes and the more familiar flat space modes is given, and indicates a general form for the overlap between all the massless supercurvature and spatially flat modes. We verify this general form by integrating over the spatially flat modes, weighted by the overlap, to obtain the original supercurvature modes. Concluding remarks follow in Section IV. Three Appendices include the supercurvature mode normalization at fixed time in the flat coordinates, and details of the overlap calculation along different hypersurfaces within de Sitter space. | In conclusion, we have given the explicit form for the overlap between the flat basis functions (\ref{flatF}) and the massless supercurvature modes. As a result, the supercurvature modes within that patch of de Sitter space which would contain our open observable universe can be written \eq \uel (x_r(x_f)) = -i \sqrt{2\ell (\ell + 1)} \int_0^\infty dk \> k^{-1/2} \> j_\ell (k) \> \phi_{k\ell m} (x_f) . \en The long-wavelength supercurvature modes are distributed over the spatially-flat basis modes, oscillating over flat-space comoving wavenumber $k$ with decreasing amplitude. More quantitatively, the spherical bessel functions $j_\ell(k)$ have their first and largest maximum (\cite{Absteg} 10.1.59) near $k \sim (\ell + \frac{1}{2}) [1 + O(\ell^{- 2/3})]$ with the approximation improving as $\ell$ increases, but the damping envelope going as $k^{-1/2}$ lowers this peak for higher $\ell$. It may be possible to use the description \cite{gar-97} of $M^2 \ne 0$ supercurvature modes as small perturbations of the massless supercurvature modes to extend the above to small $M^2$. This expression for the supercurvature mode on fixed $\eta$ surfaces, extending into region $R$, may be useful for understanding the effects of supercurvature modes during reheating. Unlike the event of bubble nucleation, reheating occurs in the future of region $C$ and hence descriptions using fixed time $r_c$ surfaces are not appropriate. Rather, it is important to understand the dynamics of these modes within region $R$, corresponding to our observable, open universe. Having an expression for the normalizable supercurvature modes on slicings extending into region $R$ is a step in separating long wavelength properties of these modes from issues related to their non-normalizability at fixed time $t_r$ in region $R$. We showed by comparing Cauchy and non-Cauchy surface calculations that in some cases non-Cauchy surface calculations of norms (in the appendix) and overlaps (in the text) agree for supercurvature modes, up to an identifiable boundary term. These non-Cauchy surfaces (fixed time in the flat coordinates) extend into the open universe and thus could be used (with caution) to calculate other properties. The complementary calculations along different surfaces were required to verify that the non-Cauchy surface representation was indeed correct. In addition, we identify one specific supercurvature mode as responsible for the well-known infrared divergence for massless scalar fields in de Sitter space. Not only does this mode have divergent norm; it is demonstrated here to be a dynamical zero-mode. This identification is a necessary first step toward its eventual replacement with an appropriate collective coordinate, similar to what has been done in closed coordinates\cite{kg93}. Something similar has been done in \cite{gar-97} in the context of two field models and quasi-open inflation, here it is found more generally as a property of the massless supercurvature modes. | 98 | 3 | gr-qc9803073_arXiv.txt |
9803 | astro-ph9803034_arXiv.txt | It is shown that the new observed redshift distributions of various flux-limited samples of radio sources in general are consistent with the predictions of two basic evolutionary models published by Condon (1984) and Dunlop \& Peacock (1990), i.e. none of them can be rejected at the confidence level of about 95 per cent. However, the models allowing a free-form evolution and suggesting both density and luminosity evolution are more consistent with the observational data \underline{at lower redshifts}, while the 'pure luminosity evolution' model fits better the data \underline{at higher redshifts}. This leads to a suspicion that the {\it same evolution} governing {\it all} radio sources, suggested by Condon (1984), might not be the case. | 98 | 3 | astro-ph9803034_arXiv.txt |
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