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9804 | astro-ph9804340_arXiv.txt | In the process of searching the Hubble Space Telescope archive, we have serendipitously discovered three populous Large Magellanic Cloud (LMC) clusters with ages that place them in the LMC `age gap.'. These clusters - NGC 2155, SL663, and NGC 2121 - turn out to have $[Fe/H]$$\sim$ --1.0 and ages of $\sim$4 Gyr. This puts them in the age gap between the intermediate-age LMC clusters, the oldest of which are $\sim$2.5 Gyr old, and ESO121-SC03, which has an age of $\sim$ 9 Gyr. The addition of these three clusters to the LMC age - metallicity relation has reduced the discrepancy between the age distribution of the LMC clusters and the field stars. Furthermore, it indicates that searches to find more clusters older than $\sim$2.5 Gyr in the LMC are crucial to a better understanding of its global star formation history. | The chemical enrichment/star formation history (SFH) is an identifying feature of every self-gravitating stellar system. One manifestation of this is the relation between age and metallicity among the star clusters in a given galaxy. Empirical information on how cluster age and metal abundance correlate provides an important clue that will eventually allow us to understand how star formation (and hence chemical enrichment) proceeds under a variety of potentially influential circumstances. It is for this reason that we strive to better define the age - metallicity relations of the cluster and field populations in galaxies. Nearby galaxies are no exception, especially the Large and Small Magellanic Clouds (L/SMC) which have provided numerous puzzles and challenges for theorists. One of the most persistent of these has been the `age gap' seen between the intermediate-age clusters ($t \lea 2.5$ Gyr; see Sec. 4.1 for a justification of this limit) and the old clusters ($t \gea 13$ Gyr) in the LMC (Geisler et al. \cite{geisea1997}). The LMC age gap is also a metallicity gap in the sense that the intermediate-age clusters have $\langle$$[Fe/H]$$\rangle$$\sim$$-0.5$, while the old clusters are closer to $\langle$$[Fe/H]$$\rangle$$\sim$$-2$. There is only one cluster that is known to lie in the gap - ESO121-SC03 with an age of $\sim$9 Gyr and $[Fe/H]\sim-1$. In contrast, recent studies based on Hubble Space Telescope (HST) data of the LMC field stars tell a signficantly different story (see Geha et al. \cite{gehea1998} for a review). While there are variations that depend on position in the LMC (Vallenari et al. \cite{valea1996}), the global star formation rate has been fairly constant for {\it most} of the LMC's history. However, sometime between 2 and 4 Gyr ago, a burst of star formation occurred producing the present population of young to intermediate age field stars and clusters (Bertelli et al. \cite{bertea1992}; Vallenari et al. \cite{valea1996}; Gallagher et al. \cite{gallea1996}; Holtzman et al. \cite{holtzea1997}). Furthermore, the models produced by Geha et al. (\cite{gehea1998}) suggest that roughly half of the LMC field stars are older than 4 Gyr. What this implies is that there should be many star clusters with ages between $\sim$2.5 Gyr and $\sim$13 Gyr, which were formed contemporaneously with the field stars. The obvious question of course is: where are these `age gap' clusters? A number of investigators have undertaken surveys to find clusters in the LMC age gap (Da Costa \cite{dac1991}; Geisler et al. \cite{geisea1997}, and references therein). The overwhelming conclusion has been that the LMC contains only one cluster with an age between 2.5 Gyr and 13 Gyr (i.e. ESO121-SC03). There is the possibility that some clusters did exist in the age gap sometime in the past, and that these have either dissolved into the LMC field or been stripped off (Olszewski \cite{olsz1993}). However, it is difficult to see how this process can preferentially affect only clusters formed during a given epoch, unless the mass spectrum of density fluctuations producing the gap clusters was somehow different from that which produced the other clusters in the LMC. The possible reasons for this are not obvious. While understanding the SFH of the LMC is extremely important, this was not our original aim when we began this study. Initially, we were searching the Hubble Space Telescope (HST) archive looking for intermediate age clusters that display the red giant branch (RGB) Bump in their color-magnitude diagrams (CMD). We did find one such cluster (NGC 411; see Alves \& Sarajedini \cite{alsa1998}). However, more importantly for the present work, we discovered three clusters whose CMDs exhibit significant numbers of stars in the Hertzsprung gap, indicative of ages older than $\sim$2.5 Gyr. These clusters are NGC 2155, NGC 2121, and SL663, and the next section describes the observations of these clusters and the data reduction. Section 3 presents the CMDs while Section 4 discusses the ages yielded by these CMDs. The implications of these results are detailed in Section 5. | The resulting relation between age and metallicity for the LMC star clusters is shown in Fig. 11. The open symbols are the ages and abundances from the work of Geisler et al. (\cite{geisea1997}) supplemented by additional clusters from Bica et al. (\cite{bicaea1998}) and the values for NGC 2193, Hodge 4, and ESO121-SC03 from this paper. The filled square represents the location of our three `age gap' clusters. Clearly, the clusters NGC 2155, NGC 2121, and SL 663 do indeed fall in the age gap between $\sim$2.5 Gyr and $\sim$13 Gyr. The reader should keep in mind, however, that if the metallicities of these clusters are higher than the values we have adopted herein, their ages will be correspondingly younger (see Sec. 4.1). As such, they may eventually be considered as belonging to the old-age tail of the IAC distribution. Future spectroscopic abundance measurements will shed more light on this. In any event, the addition of these three clusters to the age - metallicity relation of the LMC has not eliminated the discrepancy between the cluster age distribution and that of the field stars. If there are no more clusters to be discovered in the gap, then we will require some explanation for why the clusters and the field stars exhibit such differing SFHs. However, what is {\it more} likely is that there are as yet unstudied clusters in the LMC that will further fill in the age gap. Future ground-based and HST photometric surveys may reveal more such clusters. For the present paper, we have utilized archival HST/WFPC2 images of LMC populous clusters to show that there are at least three clusters in the LMC age gap - NGC 2155, SL663, and NGC 2121. These clusters have $[Fe/H]$$\sim$ --1.0 and ages of $\sim$4 Gyr. The addition of these three clusters to the LMC age - metallicity relation is the first step in reducing the significant difference between the inferred SFHs of the LMC clusters and the field stars. This strongly indicates that searches to find more clusters older than $\sim$2.5 Gyr in the LMC are crucial to a better understanding of its global SFH. | 98 | 4 | astro-ph9804340_arXiv.txt |
9804 | astro-ph9804083_arXiv.txt | Defect models have recently been declared dead\cite{watson97}, because they predict microwave background and matter fluctuations grossly out of line with what we see. In this talk we apply the fact that many defects are automatically destroyed at the time of radiation-matter transition, thus resurrecting the defects model. Moreover, the resurrected version predicts a cosmological constant, explains the apparent excess of hot clusters and the non-Gaussianity observed in galaxy surveys. If this model is correct, then the MAP and PLANCK missions will not measure what people expect them to (oscillations); rather, they will measure a broad hump. | 98 | 4 | astro-ph9804083_arXiv.txt |
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9804 | astro-ph9804169_arXiv.txt | Surface brightness fluctuations (SBFs) are much brighter in the infrared than they are at optical wavelengths, making it possible to measure greater distances using IR SBFs. We report new \Kp\ (2.1\micron) SBF measurements of nine galaxies in the Fornax and Eridanus clusters using a 1024$^2$-pixel HgCdTe array. We used improved analysis techniques to remove contributions to the SBFs from globular clusters and background galaxies, and we assess the relative importance of other sources of residual variance. We applied the improved methodology to our Fornax and Eridanus images and to our previously published Virgo cluster data. Apparent fluctuation magnitudes were used in conjunction with Cepheid distances to M31 and the Virgo cluster to calibrate the \Kp\ SBF distance scale. We find the absolute fluctuation magnitude $\MKp\,{=}\,{-}5.61\,{\pm}\,0.12$, with an intrinsic scatter to the calibration of 0.06 mag. No statistically significant change in \MKp\ is detected as a function of \v-i. Our calibration is consistent with simple (constant age and metallicity) stellar population models. The lack of a correlation with \v-i\ in the context of the stellar population models implies that elliptical galaxies bluer than $\v-i\,{=}\,1.2$ have SBFs dominated by younger (5--8 Gyr) populations and metallicities comparable to redder ellipticals. Significant contributions to the SBFs from anomalous populations of asymptotic giant branch stars are apparently uncommon in giant ellipticals. \Kp\ SBFs prove to be a reliable distance indicator as long as the residual variance from globular clusters and background galaxies is properly removed. Also, it is important that a sufficiently high signal-to-noise ratio be achieved to allow reliable sky subtraction because residual spatial variance can bias the measurement of the SBF power spectrum. | Because the light from a distant galaxy comes from discrete but unresolved stars, Poisson statistics lead to mottling of the galaxy's otherwise smooth surface brightness profile. Surface brightness is independent of distance, but the amplitude of the surface brightness fluctuations (SBFs) is not. As the distance ($d$) to a galaxy increases, the number of stars ($n$) in a given resolution element increases as $d^2$, but the observed flux ($f$) from each star is reduced by $d^{-2}$, making surface brightness ($nf$) independent of distance. On the other hand, the rms variation in observed flux from region to region is $n^{1/2}f$, which scales as $d^{-1}$. The variance in surface brightness $nf^2$ normalized by the mean galaxy brightness $nf$ decreases with distance; distant galaxies appear smoother than nearby galaxies. Because the variance in surface brightness is proportional to the second moment of the stellar luminosity function, it is dominated by luminous red giant stars. Although individual stars are not resolved, measuring SBFs probes the stellar population of the galaxy directly. With good theoretical models for stellar populations, the absolute magnitude of SBFs can be calculated, allowing a direct determination of the distance that is independent of the global dynamics or the environment of the galaxy (to the extent that the stellar populations in old stellar systems are independent of these variables). Alternatively, measuring SBFs in galaxies with known distances provides insight into stellar populations, allowing comparison with stellar evolution models and providing an empirical calibration of the SBF distance scale. Good descriptions of the theory and practice of using SBFs as a distance measurement tool and stellar population probe can be found in several papers by Tonry and collaborators (Tonry et al. 1997; Jensen, Luppino, \& Tonry 1996, hereafter JLT; Tonry 1991; Tonry, Ajhar, \& Luppino 1990; Tonry \& Schneider 1988), in papers by Sodemann \& Thomsen (1995, 1996), and in the review by Jacoby et al. (1992). The first $K$-band SBF studies are described by JLT, Luppino \& Tonry (1993), and Pahre \& Mould (1994). J. Tonry and coworkers have completed an extensive survey of $I$-band SBF distances in a sample of several hundred early-type (dynamically hot elliptical and S0) galaxies which is more than 50\% complete out to $\sim$2800 \kms\ (Tonry et al. 1997). They find that the $I$-band absolute fluctuation magnitude \MI\ is a linear function of \v-i\ and has a universal zero point. Their $I$-band calibration, which is based empirically on Cepheid distances, is in good agreement with Worthey's (1993a,b; 1994) simple stellar populations models. Worthey's models predict that the effects of age and metallicity in the $I$ band are largely degenerate, so that \MI\ may be calibrated using a single parameter such as the \v-i\ color. The resulting intrinsic scatter of the $I$-band SBF distance scale is of order 0.07 mag. The purpose of the current study is to examine the behavior of the SBF calibration in the near-IR \Kp\ band, where Worthey's models predict a much weaker dependence on \v-i, but potentially larger scatter as the effects of age and metallicity are no longer degenerate. Stellar surface brightness fluctuations are very red, since they are dominated by luminous red giant stars. The advantages of observing IR SBFs are clear: fluctuations are ${\sim}33$ times brighter at $K$ than at $I$, making them observable to greater distances. The SBF amplitude is inversely proportional to the seeing FWHM, and the seeing is typically much better in the IR than at optical wavelengths. Fluctuations are also red compared to the globular cluster (GC) population, so the contrast between stellar SBFs and GCs is higher at $K$. Finally, stellar population models predict that $K$-band SBF magnitudes have a much weaker dependence on \v-i\ than at $I$, reducing or eliminating the need to accurately measure the color of the galaxy (Worthey 1993a). Several studies have demonstrated the feasibility of measuring IR SBFs, including Luppino \& Tonry (1993), Pahre \& Mould (1994), and JLT. These papers report results for a rather limited set of galaxies in the Local Group (LG) and the Virgo Cluster, but the consistency in the measured calibration of \MK\ is encouraging ($\MKp\,{=}\,{-}5.61\,{\pm}\,0.16$ Luppino \& Tonry, $M_{K_{sh}}\,{=}\,{-}5.77\,{\pm}\,0.18$ Pahre \& Mould, $\MKp\,{=}\,{-}5.62\,{\pm}\,0.29$ JLT). All three of these papers report absolute fluctuation magnitudes that are consistent with Worthey's (1994) predictions based on simple stellar population models. While the mean \MKp\ of JLT and Pahre \& Mould (1994) are quite consistent, differences in apparent fluctuation magnitudes for individual galaxies are larger than the stated uncertainties allow. JLT discuss several possible reasons for the disagreement in fluctuation magnitudes and showed that residual variance resulting from spatial variations in the detector sensitivity and dark current can contribute significantly to the power spectrum in low signal-to-noise ratio ($S/N$) observations. A residual pattern of only 0.1\% of the sky level can significantly change the fluctuation magnitude measured. To use IR SBFs as a reliable distance indicator, it is critical that the observations be of sufficiently high $S/N$ ratio to avoid biases from residual variances. JLT also demonstrated that properly subtracting the sky and carefully sampling the point spread function (PSF) are crucial to accurately measuring the fluctuation amplitude. An uncertainty in the background subtraction of 2\% of the sky level can dominate the uncertainty in the fluctuation magnitude. Our Virgo sample showed a large dispersion (0.29 mag) due to the depth of the cluster and observational errors. Pahre \& Mould (1994) also addressed the dispersion in fluctuation magnitude in a similarly-sized sample of Virgo ellipticals. Clearly a larger sample of galaxies is needed to calibrate the IR SBF distance scale. To better understand the effects stellar populations have on the fluctuation amplitude, we observed five galaxies in the Fornax cluster. Fornax is ideal for several reasons: first, it is much more centrally concentrated than the Virgo cluster. The reduced dispersion in distances allows us to better quantify and understand the uncertainty in our SBF measurements. Second, Fornax contains a large number of giant elliptical galaxies with a wide range of \v-i\ colors, metallicities, and globular cluster populations. The Virgo galaxies observed by JLT and Pahre \& Mould (1994) spanned a very limited range in \Mg2\ index to avoid stellar population variations which may affect the calibration of \MK. We now extend the sample to calibrate \Kp\ fluctuation magnitudes across a wider range of metallicities and \v-i\ colors to compare with theoretical stellar population models. Finally, Cepheid distances to several Virgo cluster galaxies and to the Fornax spiral NGC~1365 have been measured using the Hubble Space Telescope (HST). We can empirically anchor our \Kp-band SBF calibration both to the Cepheid and the $I$-band SBF distance scales. This paper reports results from our observations of four Eridanus cluster elliptical galaxies in addition to the Fornax galaxies. Eridanus is not as compact a cluster as Fornax, but we can still compare our IR SBF results with the $I$-band results from Tonry (1997). Finally, we return to our SBF measurements for the seven Virgo ellipticals discussed by JLT and apply the improved techniques we describe in this paper. We present the updated results for these galaxies in Section~\ref{virgorevisited}, along with new observations of NGC~4365. | $K$-band SBFs can be measured reliably and used to determine distances to early-type galaxies. Because SBFs are relatively bright in the IR and the seeing is typically very good, integration times can be quite modest. However, observations must be sufficiently deep to adequately sample the GC luminosity function and remove sources of residual variance. Low-$S/N$ ratio measurements are unreliable. Optical images can be used to identify and remove GCs and background galaxies that are not detected in the IR, improving the IR SBF measurement. We empirically calibrated the \Kp\ SBF distance scale using Cepheid distances to M31 and the Virgo cluster. The absolute \Kp\ fluctuation magnitude is $\MKp\,{=}\,{-}5.61\,{\pm}\,0.06$ (statistical error) with a total uncertainty of 0.12 mag. No significant change in \MKp\ with \v-i\ was observed over the range in color spanned by the galaxies in this sample. Accurate color measurements are not required to measure the \Kp\ SBF distance to a galaxy. \Kp\ SBF magnitudes are consistent with predictions from simple stellar population models. The lack of a correlation in \MKp\ with \v-i\ is best explained by a spread in ages among the galaxies observed. The redder ellipticals are consistent with 12 to 17 Gyr stellar population models, while the bluer galaxies in our sample must have younger 5 to 8 Gyr populations. Metallicities appear to vary less than ages, with the typical galaxy having [Fe/H]\,=\,${-}0.25$. The stellar population models show that the age-metallicity degeneracy is broken with $K$-band SBFs, allowing one to distinguish between old, metal-poor and young, metal-rich populations. Examining the radial variation in IR SBFs will help distinguish between different galaxy formation scenarios. Our observations do not agree with the relationship between \MKp\ and the \Mg2\ index predicted by the stellar population models. | 98 | 4 | astro-ph9804169_arXiv.txt |
9804 | astro-ph9804263_arXiv.txt | The high redshift radiogalaxy 1243+036 ($z=3.6$) presents an asymmetric \lya\ profile of FWHM 1550 \kms\ as measured by van~Ojik \etal\ We propose that the blue asymmetry in the \lya\ profile is not due to narrow absorption dips but consists of narrow emission peaks. We interpret the blueshifted peaks near $-1130, -850$ and $-550 \kms$ (relative to the peak of full profile) as being the result of Fermi acceleration of \lya\ produced by jet-induced star formation in the wake of a 300~\kms\ shock. This shock would be caused by the deflection of the radio-jet at the observed position of the radio bend which also coincides spatially with the excess \lya\ emission reported by van~Ojik \etal\ | \label{intro} Lyman $\alpha$ is the strongest line observed in very high redshift radio galaxies (HZRG). Although the brightness of \lya\ peaks near or at the nuclear position, most of the emission is spatially resolved with the fainter emission extending up to radii 40--130\,kpc. The \lya\ profile in HZRG is characterized by a FWHM in the range 700--1600\kms\ (van~Ojik 1995 and references therein). The intermediate resolution study of \lya\ profiles carried out by van~Ojik \etal\ (1997) has revealed the presence of troughs which are well explained by \hi\ gas absorption present in the environment of the parent radio galaxy. In their sample, however, the HZRG 1243+036 ($z=3.6$) may call for a different interpretation, namely that the broad profile presents true narrow emission features in between the narrow `dips' which van~Ojik \etal\ (1996: vO96) have interpreted as absorption features. The repeated scattering of the resonance \lya\ line across a shock discontinuity was shown by Neufeld \& McKee (1988: NM88) to result in a systematic blueshift of \lya. In this paper, we develop in more detail the Fermi acceleration model and propose that the narrow features observed by vO96 correspond to a small number of across-shock scatterings. | 98 | 4 | astro-ph9804263_arXiv.txt |
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9804 | astro-ph9804325_arXiv.txt | Recent discoveries by the {\it Rossi X-Ray Timing Explorer} indicate that most of the rapidly accreting ($\dot M \gtrsim 10^{-11} M_\odot \ {\rm yr}^{-1}$) weakly magnetic ($B\ll 10^{11} \ {\rm G}$) neutron stars in the Galaxy are rotating at spin frequencies $\nu_s \gtrsim 250 \ {\rm Hz}$. Remarkably, they all rotate in a narrow range of frequencies (no more than a factor of two, with many within 20\% of 300 Hz). I suggest that these stars rotate fast enough so that, on average, the angular momentum added by accretion is lost to gravitational radiation. The strong $\nu_s$ dependence of the angular momentum loss rate from gravitational radiation then provides a natural reason for similar spin frequencies. Provided that the interior temperature has a large scale asymmetry misaligned from the spin axis, then the temperature sensitive electron captures in the deep crust can provide the quadrupole needed ($\sim 10^{-7} M R^2$) to reach this limiting situation at $\nu_s\approx 300$ Hz. This quadrupole is only present during accretion and makes it difficult to form radio pulsars with $\nu_s>(600-800) \ {\rm Hz}$ by accreting at $\dot M \gtrsim 10^{-10} M_\odot \ {\rm yr^{-1}}$. The gravity wave strength is $h_c\sim (0.5-1) \times 10^{-26}$ from many of these neutron stars and $>2\times 10^{-26}$ for Sco X-1. Prior knowledge of the position, spin frequency and orbital periods will allow for deep searches for these periodic signals with gravitational wave interferometers (LIGO, VIRGO and the ``dual-recycled'' GEO 600 detector) and experimenters need to take such sources into account. Sco X-1 will most likely be detected first. | The launch of the {\it Rossi X-Ray Timing Explorer} (RXTE) has allowed for the discovery of fast quasi-periodic variability from many rapidly accreting ($ \dot M\gtrsim 10^{-11} M_\odot \ {\rm yr^{-1}}$) neutron stars. These observations strongly suggest that these neutron stars (NSs) are rapidly rotating, as predicted by those scenarios connecting the millisecond radio pulsars to this accreting population (see Bhattacharya 1995 for an overview). Strohmayer et al. (1996) were the first to detect nearly coherent $\nu_B=363$ Hz oscillations during type I X-ray bursts from the low accretion rate ($\dot M< 10^{-9} M_\odot \ {\rm yr^{-1}}$) NS 4U~1728-34. Pulsations were detected in six of the eight bursts analyzed at that time. In addition, two high frequency quasi-periodic oscillations (QPOs) were seen in the persistent emission. These changed with accretion rate, but maintained a fixed difference frequency of $\nu_d\approx 363 $ Hz, identical to the period seen during the bursts. The detection of two drifting QPO's (in the persistent emission) separated by a fixed frequency identical to that seen in the bursts naturally leads to beat frequency models (Strohmayer et al. 1996; Miller, Lamb, \& Psaltis 1998). The difference frequency is presumed to be the NS spin frequency, $\nu_s$, whereas the upper frequency has different origins in different models (see van der Klis 1998 for a summary). In addition, the temporal behavior of the periodic oscillations both during the rise of the bursts (Strohmayer, Zhang, \& Swank 1997b) and in the cooling tails (Strohmayer et al. 1997a) are most easily explained in terms of rotation. There are six NSs with measured periodicities during Type I X-ray bursts (see Table 1). Both the difference frequencies ($\nu_d$) and the burst frequencies ($\nu_B$) are in a narrow range, from 260 to 589 Hz. For two objects (KS 1731-260 and 4U 1636-53) the difference frequencies are one-half the burst values. Which value is $\nu_s$ is not resolved. There are also many NSs that accrete at higher rates and are not regular Type I X-ray bursters. Many of these objects, notably the ``Z'' sources, also show drifting QPO's at fixed separation, again with a similarly narrow frequency range (roughly 250-350 Hz). Beat-Frequency like models are also applied to these observations so as to infer $\nu_s$. The applicability of such a model is less clear when the difference frequency is not constant (Sco X-1, van der Klis et al. 1997; 4U 1608-52, Mendez et al. 1998). If accreting matter always arrives with the specific angular momentum of a particle orbiting at the NS radius ($R=10 R_6 {\rm km}$), then it only takes $\sim 10^7 \ {\rm yrs}$ of accretion at $\dot M\approx 10^{-9} M_\odot \ {\rm yr^{-1}}$ for a $M=1.4M_{1.4} M_\odot$ NS to reach $\nu_s=50$ Hz from an initially low frequency. It is thus remarkable that these NSs are all rotating at nearly the same rate. White and Zhang (1997) argued that this similarity arises because these NSs are magnetic and have reached an equilibrium where the magnetospheric radius equals the co-rotation radius. This requires an intrinsic relation between their magnetic dipoles, $\mu_b$, and $\dot M$ so that they all reach the same rotational equilibrium (most likely $\mu_b\propto \dot M^{1/2}$) and a way of hiding the persistent pulse typically seen from a magnetic accretor. My alternative explanation for these spin similarities is that gravitational wave (GW) emission has started to play an important dnrole. If the NS has a misaligned quadrupole moment, $Q$, then the strong spin frequency dependence of GW emission defines a critical frequency beyond which accretion can no longer spin-up the star. Such a NS will radiate energy via GW's at the rate $\dot E=32 GQ^2\omega^6/5 c^5$, where $\omega=2\pi \nu_s$, and lose angular momentum at the rate $N_{gw}=\dot E/\omega$. Balancing this spin-down torque with the characteristic spin-up torque from time-averaged accretion, $N_a\approx \dot M(GMR)^{1/2}$, gives the $Q$ needed so as to make the critical frequency 300 Hz, \begin{equation}\label{eq:qneed} Q\approx 4.5 \times 10^{37} \ {\rm g \ cm^2}\left(\dot M\over 10^{-9} \ {\rm M_\odot \ yr^{-1}}\right)^{1/2}\left(300 \ {\rm Hz}\over \nu_s\right)^{5/2}, \end{equation} or $<10^{-7}$ of the NS moment of inertia, $I\approx 10^{45} \ {\rm g \ cm^2}$. The similarities in $\nu_s$ may then arise because of the weak dependencies of the critical frequency on $Q$ and $\dot M$. What is the source of the misaligned quadrupole? Wagoner (1984) argued that accreting NSs would get hung-up at spin frequencies where the Chandrasekhar-Friedman-Schutz (CFS) instability sets in. However, Lindblom (1995) and Lindblom \& Mendell (1995) have shown that the star needs to be very near the breakup frequency ($\nu_s \gtrsim \ {\rm kHz}$) for such an instability to occur, even for the core temperatures $T_c=(1-3)\times 10^8 \ {\rm K}$ of rapidly accreting NSs (Ayasli \& Joss 1978; Brown \& Bildsten 1998). The spin frequencies for these NSs are too slow for such an instability. I present in \S 2 a new source for lateral density variations in an accreting NS; electron captures (hereafter EC) in the crust. The constant compression of the crust forces nuclei to undergo EC when the electron Fermi energy, $E_F$, is high enough to make a transition. However, the crust is hot enough in a rapidly accreting NSs to make the EC rates temperature sensitive. Hotter regions then capture at lower pressures, so that the density jump associated with the EC is at a higher altitude in the hotter parts of the crust. Moderate lateral temperature variations then lead to density variations large enough to generate the required $Q$. This outcome is independent of the particular source of the temperature variations. One possible cause for a $T$ asymmetry relative to the spin axis is a weak magnetic field. I conclude in \S 3 by finding the GW signal strength and estimating detection. | 98 | 4 | astro-ph9804325_arXiv.txt |
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9804 | astro-ph9804113_arXiv.txt | \vspace{-3ex} We present a numerical simulation of the dynamical collapse of a nonrotating, magnetic molecular cloud core and follow the core's evolution through the formation of a central point mass and its subsequent growth to a $1~\msol$ protostar. The epoch of point-mass formation (PMF) is investigated by a self-consistent extension of previously presented models of core formation and contraction in axisymmetric, self-gravitating, isothermal, magnetically supported interstellar molecular clouds. Prior to PMF, the core is dynamically contracting and is not well approximated by a quasistatic equilibrium model. Ambipolar diffusion, which plays a key role in the early evolution of the core, is unimportant during the dynamical pre-PMF collapse phase. However, the appearance of a central mass, through its effect on the gravitational field in the inner core regions, leads to a ``revitalization'' of ambipolar diffusion in the weakly ionized gas surrounding the central protostar. This process is so efficient that it leads to a decoupling of the field from the matter and results in an outward-propagating hydromagnetic C-type shock. The existence of an ambipolar diffusion-mediated shock of this type was predicted by Li \& McKee (1996), and we find that the basic shock structure given by their analytic model is well reproduced by our more accurate numerical results. Our calculation also demonstrates that ambipolar diffusion, rather than Ohmic diffusivity operating in the innermost core region, is the main field decoupling mechanism responsible for driving the shock after PMF. The passage of the shock leads to a substantial redistribution, by ambipolar diffusion but possibly also by magnetic interchange, of the mass contained within the magnetic flux tubes in the inner core. In particular, ambipolar diffusion reduces the flux initially threading a collapsing $\sim 1~\msol$ core by a factor $\simgt 10^3$ by the time this mass accumulates within the inner radius ($\simeq 7.3~{\rm AU}$) of our computational grid. This reduction, which occurs primarily during the post-PMF phase of the collapse, represents a significant step towards the resolution of the protostellar magnetic flux problem. Our calculations indicate that a $1~\msol$ protostar forms in $\sim 1.5 \times 10^5~\rm{yr}$ for typical cloud parameters. The mass accretion rate is time dependent, in part because of the C-shock that decelerates the infalling matter as it propagates outward: the accretion rate rises to $\simeq 9.4~\msol~\rm{Myr}^{-1}$ early on and decreases to $\simeq 5.6~\msol~{\rm{Myr}}^{-1}$ by the time a solar-mass protostar is formed. The infalling gas disk surrounding the protostar has a mass $\sim 10^{-2}~\msol$ at radii $r \simgt 500~\rm{AU}$. A distinguishing prediction of our model is that the rapid ambipolar diffusion after the formation of a protostar should give rise to large ($\simgt 1~{\rm{km}}~{\rm s}^{-1}$), and potentially measurable, ion--neutral drift speeds on scales $r \simlt 200~\rm{AU}$. The main features of our simulation, including the C-shock formation after PMF, are captured by a similarity solution that incorporates the effects of ambipolar diffusion (Contopoulos, Ciolek, \& K\"{o}nigl 1997). | It is generally accepted that most of the star-formation activity in our galaxy takes place through the gravitational collapse of molecular cloud cores (e.g., Mouschovias 1987; Shu, Adams, \& Lizano 1987). It is, furthermore, believed that interstellar magnetic fields play a central role in this process in that their stresses are the dominant agent that acts against gravity to prevent, or delay, cloud collapse (e.g., Mouschovias 1978; McKee et al. 1993). This is embodied in the concept of a critical mass $M_{\rm crit}$, which in general takes account of both the magnetic and the thermal pressure contributions to the support of the cloud, but which, in the case that magnetic stresses dominate, reduces to $M_{\rm crit} \approx 0.13 \phiB/G^{1/2}$, where $G$ is the gravitational constant and $\phiB$ is the magnetic flux that threads the cloud. Clouds whose mass $M$ exceeds $M_{\rm crit}$ are ``supercritical'': they collapse on the free-fall timescale. In contrast, ``subcritical'' clouds ($M < M_{\rm crit}$) can avoid collapse on the much longer ambipolar diffusion timescale. In the latter case, the neutrals gradually contract by diffusing inward through the ions and field, leaving behind a magnetically supported envelope and eventually forming a supercritical core that undergoes dynamical collapse (e.g., Mouschovias 1996). Because of the complexity of the problem --- it involves solving the full dynamical equations of a magnetized, multicomponent (neutrals, ions, electrons, as well as charged and neutral grains) fluid that evolves over many decades in size and density in a nonspherically symmetric manner (because of the presence of ordered magnetic fields and likely also rotation) --- much of the progress in this area has been accomplished through the use of numerical simulations. One of the main efforts to simulate core formation and contraction due to ambipolar diffusion in magnetically supported molecular clouds has been carried out by Mouschovias and coworkers (e.g., Fiedler \& Mouschovias 1992, 1993; Ciolek \& Mouschovias 1993, 1994, 1995, hereafter CM93, CM94, CM95, respectively; Basu \& Mouschovias 1994, 1995a, b). These studies followed the evolution of the core over six decades in density up to central densities $\sim 3 \times 10^9~\cc$, where the assumption of isothermality starts to break down because of radiative trapping (e.g., Gaustad 1963; Hayashi 1966). This assumption had been adopted in the interest of simplicity: sophisticated and computationally intensive numerical techniques are generally needed to calculate the thermal structure of the gas during the opaque phase of protostellar evolution (e.g., Larson 1969, 1972; Winkler \& Newman 1980; Stahler, Shu, \& Taam 1981; Boss 1984; Myhill \& Kaula 1992; Myhill \& Boss 1993). As a result of this restriction, the aforementioned calculations did not follow the collapse of the core to the time where a point mass -- a protostar -- is formed at the center, although they did obtain valuable information on the conditions leading to this critical event. In particular, by the time these simulations were terminated, the inner region of the core was collapsing dynamically and was characterized by neutral infall speeds $\sim C$ (the isothermal speed of sound) and inward accelerations $\simgt 0.3 |\gr|$ [where $\gr(r)$ is the local gravitational acceleration]. Furthermore, the thermal pressure, while remaining relatively unimportant in the envelope, came to exceed the magnetic pressure near the center. Basu (1997) derived a time-dependent, semianalytic solution that extended these ambipolar diffusion models up to the instant of point-mass formation (hereafter referred to as PMF \footnote{Creation of a central point mass was commonly referred to in previous papers as ``core formation.'' In this paper we use the term ``point-mass formation'' so as not to confuse this process with the formation of an extended, magnetically supercritical, molecular cloud core.}). He found that ambipolar diffusion continues to gradually erode the retarding magnetic forces in the inner core, making the collapse increasingly more dynamical (and the thermal-to-magnetic pressure ratio in the inner core progressively larger) as PMF is approached. The diminution of magnetic forces in the innermost regions of a collapsing core just prior to PMF suggests that one could gain some insight into the protostar formation process from previous studies of PMF in {\it nonmagnetic}, spherically symmetric, isothermal clouds. Analytic similarity solutions have uncovered two limiting behaviors: Penston (1969) and Larson (1969) found a solution in which, just before PMF, the infall speed approaches $\sim 3.3~C$ at all radii while the density scales with radius as $r^{-2}$, resulting in a spatially uniform mass inflow rate $\sim 29~C^3/G$ (where $G$ is the gravitational constant). Hunter (1977) extended this solution past PMF and showed that, immediately after the central mass is formed, the accretion rate onto the protostar increases to $\sim 47~C^3/G$. In the other limit, Shu (1977) obtained a solution that is static prior to PMF (with the density distribution of a singular isothermal sphere, which also scales as $r^{-2}$) and that takes on an expansion-wave character (with a constant mass accretion rate $\sim 0.98~C^3/G$ onto the protostar) following PMF. Hunter (1977) and Whitworth \& Summers (1985) demonstrated that there are, in fact, infinitely many similarity solutions that span the range between the Larson-Penston and Shu results, with the nature of any given solution being determined by the initial configuration of the cloud and the conditions at its boundary. Numerical simulations carried out by Hunter (1977) and by Foster \& Chevalier (1993) confirmed the dependence on the initial and boundary conditions. In particular, it was found that the behavior of the central regions of clouds that are initially marginally stable to collapse approximates that of the Larson-Penston solution at the PMF epoch, although it was determined that the mass accretion rate onto the protostar declines at later times. It was, however, also found that the post-PMF evolution of clouds that initially have more extended envelopes approximates that of the Shu solution at late times. Since the initial and boundary conditions of real clouds are expected to depend on the detailed configuration and evolution of the embedded magnetic field, it is clear that one needs to incorporate magnetic field effects into the collapse calculations to adequately model the formation of protostars. There have been several recent attempts to calculate PMF following the collapse of magnetic interstellar clouds. Although they have all contributed to our understanding of the processes involved, their results were hampered by the adopted assumptions or approximations. For example, Tomisaka (1996) modeled clouds that had equal thermal and magnetic energy densities, so that they were not primarily supported by magnetic fields. This means that his model clouds were magnetically supercritical. This assumption is at variance with H {\sc I} and OH Zeeman measurements of magnetic field strengths in molecular clouds, which are consistent with models of magnetically subcritical clouds (Crutcher et al. 1993, 1994, 1996). Li \& Shu (1997) modeled PMF in self-similar, magnetic cores. They assumed that cores immediately before PMF can be represented by hydrostatic configurations of singular isothermal disks and that the magnetic flux is frozen into the neutrals. These assumptions are inconsistent with the above-cited results of numerical simulations and semianalytic solutions of the collapse of magnetically supported molecular clouds that undergo ambipolar diffusion (as well as with the simulations of thermally supported spherical clouds), which have found that the inner core regions collapse dynamically as PMF is approached (see also \S 3.2). As we show below, ambipolar diffusion, which plays a key role in bringing about the dynamical collapse, is generally important also {\it after} PMF. Safier, McKee, \& Stahler (1997) did include ambipolar diffusion in the modeling of magnetic cloud collapse and post-PMF evolution. However, their model is strictly spherical, and they neglected the effect of thermal pressure gradients in comparison with magnetic stresses. Furthermore, their formulation is quasistatic and does not incorporate the magnetic induction equation for the time evolution of the magnetic field. Li (1998) extended the Safier et al. model by including thermal-pressure and time-dependent terms and by adding the induction equation. This enabled him to follow the time evolution of his model cores (for $r > 150~\rm{AU}$) even during the dynamical phases of the collapse. However, by retaining the spherical-symmetry assumption of Safier et al., his calculations were also unable to yield the geometry of the magnetic field lines. \footnote{The same is true for the spherically symmetric self-similar model of a collapsing magnetic cloud devised by Chiueh \& Chou (1994), which, however, does not include ambipolar diffusion. It should be noted that all models that assume spherical symmetry do not satisfy the solenoidal condition $\nabla \cdot \Bvec =0$ on the magnetic field.} The importance of ambipolar diffusion after PMF can be inferred by comparing the ambipolar diffusion timescale $\tad = r/\vd$ (where $\vd$ is the ion--neutral drift speed) and the gravitational contraction ($\simeq$ free-fall) timescale $\tgr = (r/|\gr|)^{1/2}$ before and after PMF. In axisymmetric geometry, $\tad/\tgr \simeq (\tgr/\tni) \, {\mu_ B}^2$ in the inner flux tubes of a supercritical core (e.g., Mouschovias 1991), where ${\mu_B}(r)$ is the total mass-to-flux ratio at radius $r$ (in units of the critical value for gravitational collapse) and \begin{equation} \label{tnieq} \tni = 1.4 \left[1 + 0.067 \frac{\left(\mHII /2~\rm{a.m.u.}\right)}{\left(\mi/30~\rm{a.m.u.}\right)}\right] \frac{1}{\nni \sigw} \end{equation} is the neutral--ion collision time. In equation (\ref{tnieq}), $\nni$ is the ion density and $\sigw$ is the average collisional rate between ions of mass $\mi$ and neutral $\HII$ molecules of mass $\mHII$ ($\simeq 1.7 \times 10^{-9}~{\rm{cm^3}}~{\rm s}^{-1}$ for collisions between neutrals and $\Mgp$ or $\HCOp$ ions; McDaniel \& Mason 1973); the factor 1.4 accounts for a 20\% helium abundance by number. CM94 and CM95 found that, to first order for the late pre-PMF evolution of cores in disk-like clouds, the magnetic field $B \approx 3\, (r_0/r)~\rm{mG}$ (where $r_0 = 40~\rm{AU}$), $\nni \approx 0.1~\cc$ (valid for neutral densities $\nn \simgt 10^7~\cc$; see Figs. $2b$ and $4b$ in CM94), the vertical column density $\sign \approx 5\, (r_0/r)~{\rm g}~\rm{cm}^{-2}$, and the total mass $M(r) \approx 6 \times 10^{-3} (r/r_0) \msol$. For a disk-like cloud, the critical mass-to-flux ratio $(M/\phiB)_{\rm{d,crit}}= (4 \pi^2 G)^{-1/2}$, where $G$ is the gravitational constant (Nakano \& Nakamura 1978). Therefore, one finds $\mu_B \approx 2.7$, $|\gr| \approx G M(r)/r^2 \approx 2 \times 10^{-6} (r_0/r) ~{\rm{cm}}~\rm{s}^{-2}$, and $\tni \approx 7 \times 10^9~\rm{s}$, which yields $\tgr \approx 2 \times 10^{10} (r/r_0) ~\rm{s}$ and $\tad/\tgr \approx 20\, (r/r_0)$. It follows that $\tad/\tgr \gg 1$ for $r \gg 2~\rm{AU}$. Hence, as has already been demonstrated in the past, ambipolar diffusion is ineffective as a dynamically collapsing core approaches PMF, and for $r \gg 2~\rm{AU}$ the magnetic flux can be considered frozen into the neutrals. (For reasons discussed in \S 2.2, we do not consider $r \simlt 5~\rm{AU}$.) Turning now to the post-PMF epoch, when the central point mass comes to dominate the gravitational field in the innermost flux tubes, we use the relation $\tad/\tgr \approx \left(1- |\an|/|\gr|\right)^{-1} (\tgr/\tni)$, where $\an$ is the inward acceleration of the neutrals and $\gr$ is the total gravitational acceleration (see eqs. [\ref{rforceeq}] and [\ref{tadeq}] in \S 3.3). In this case, normalizing $r$ as before, $|\gr| \approx 4 \times 10^{-4} (\mcent/\msol) (r_0/r)^2 {\rm{cm}}~\rm{s}^{-2}$. Substituting again $\tni \approx 7 \times 10^9~{\rm s}$ (corresponding to $\nni \approx 0.1~\cc$) , the above ratio becomes $\tad/\tgr \approx 0.2\, (r/r_0)^{3/2}(\msol/\mcent)^{1/2}\left(1-|\an|/|\gr|\right)^{-1}$. For $|\an|/|\gr|$ in the range 0.2 -- 0.9, which corresponds to the period after PMF when the collapse becomes progressively more dynamical, one infers $\tad/\tgr \approx (0.2 - 2) (r/r_0)^{3/2}(\msol/\mcent)^{1/2}$. Hence $\tad \simlt \tgr$ for $\mcent \simgt 0.1 \msol$ and $r \simlt 50~\rm{AU}$. [A similar result is obtained if one continues to use the pre-PMF relations and simply substitutes $\mcent$ for $M(r)$.] This estimate indicates that {\em decoupling of the gas and magnetic field by ambipolar diffusion should occur in the inner core regions after PMF}. The physical reason for this is that the strength of the gravitational field in the weakly ionized gas near the origin is greatly enhanced by the appearance and growth of a central point mass, causing the neutrals to fall in more rapidly while the plasma and magnetic field are left behind. (The same basic reason --- the appearance of a progressively growing free-fall zone around the origin following PMF --- is also the cause of the increase in the mass inflow rate into the center at that epoch first discovered in the above-referenced nonmagnetic collapse calculations.) The foregoing conclusion is verified by a detailed calculation in \S~3.3, where we show that ambipolar diffusion after PMF is, in fact, so efficient that it effectively decouples the neutrals and magnetic field in the innermost core region, with dramatic consequences for the subsequent dynamical evolution of the core. An alternative, yet equivalent, analysis of the effectiveness of ambipolar diffusion following PMF, based on the scaling of magnetic forces (particularly the magnetic tension force) after PMF, is given in Appendix C. In this paper we present a detailed numerical simulation of point-mass formation in a nonrotating, magnetic, dynamically collapsing protostellar core, properly accounting for the effect of ambipolar diffusion. Unlike earlier studies, we use an initial state that is consistent with the realistic models of core formation and collapse, as presented earlier by Mouschovias and coworkers. As we noted above, the simulations carried out by that group were terminated at densities where the isothermality assumption started to become invalid because of radiative trapping. Although a proper treatment of radiative trapping is indispensable for a complete treatment of the star-formation process, one can adopt a simpler approach that circumvents this difficulty by removing the region of radiative trapping from the active computation mesh. This is justified by the fact that the region of radiative trapping (typically a few AU) is several orders of magnitude smaller than the characteristic core size ($r_{\rm{core}} \approx 0.1~\rm{pc}$). This region can therefore be considered to be effectively point-like, and one can proceed to calculate the formation of a central point mass within a gravitationally collapsing core and its effect on the subsequent evolution of the core by retaining the isothermality assumption. In adopting this approach, we note that the assumption of isothermality was also employed in previous studies of protostar formation in nonmagnetic cloud cores as well as in the more recent attempts to model PMF in magnetic clouds. The plan of the paper is as follows. In \S~2 we review the main characteristics of the pre-PMF core-evolution calculations of CM93, CM94, and CM95 and outline the model modifications that we have implemented to extend the simulations beyond PMF. In \S~3 we present the results of our calculations. We consider a typical model and follow the evolution of all physical quantities of interest through the PMF epoch. We also describe the formation (due to ambipolar diffusion in the innermost core after PMF) of a C-type shock and consider its propagation through the collapsing core. The appearance of such a shock as a result of field--matter decoupling was first pointed out by Li \& McKee (1996), who proposed that the relevant field decoupling mechanism was Ohmic dissipation in the innermost regions of the core. Our study, however, reveals that ambipolar diffusion occurring outside the region of Ohmic dissipation is the main cause of magnetic flux decoupling after PMF. In \S~4 we discuss the structure of the ambipolar diffusion-mediated shock and show that our numerical calculations qualitatively reproduce the predictions of the simplified analytic model constructed by Li \& McKee (1996). We present a quantitative comparison with their results and also address the issue of interchange instability in the post-shock region, first raised in their work, in light of our detailed computations. In that section we also discuss the observational implications of our simulation and briefly comment on the magnetic flux problem during star formation. Our results are summarized in \S~5. | \subsection{Observational Comparisons and Predictions} We may compare our result for the protostellar accretion rate during the PMF epoch with observations of star-forming cores. As shown in \S~3.3, the accretion rate rises rapidly early on to $\mcentdot \simeq 9.4~\msol~\rm{Myr}^{-1}$ for $\delt \simlt 10^3~\rm{yr}$ (see Fig. $2b$) and stays at this value up to the formation of the hydromagnetic shock. For $\delt \simgt 4 \times 10^3~\rm{yr}$ the shock is able to decelerate the infalling matter, and the accretion rate decreases to $\mcentdot \simeq 5.6~\msol~\rm{Myr}^{-1}$ by $\delt \simeq 1.5 \times 10^5~\rm{yr}$ (the time when $\mcent = 1~\msol$). Therefore $\mcentdot$ decreases with increasing central mass (see Fig. $2c$). This is consistent with estimates of ages and accretion rates ($\propto t_{\rm{age}}^{-1}$) for young stellar objects, as deduced from evolutionary diagrams inferred from observations of Class 0 and Class I objects (e.g., Saraceno et al. 1996). In particular, the lifetimes of Class 0 objects were estimated in this way to be an order of magnitude shorter than those of Class I objects, providing evidence for a decrease in the protostellar accretion rate as an object evolves from a Class 0 source to a Class I source (e.g., Andr\'{e} 1995; Ward-Thompson 1996). Another argument for a time-dependent accretion rate was given by Bontemps et al. (1996), who analyzed the observed CO momentum flux of several young stellar objects and found a noticeable decline in the CO flux with decreasing circumstellar envelope mass. They suggested that this is indicative of a decrease in the protostellar accretion rate (which they assumed to be proportional to the mass outflow rate) as an object evolves from Class 0 to Class I. It is also of interest to note the ${}^{13}{\rm{CO}}(J=1-0)$ observations of infalling disks for the protostellar candidates HL Tauri (Hayashi, Ohashi, \& Miyama 1993) and L1551-IRS5 (Ohashi et al. 1996). HL Tauri has a mass $\sim 0.6~\msol$ and a surrounding disk with radius $\sim 1400~\rm{AU}$ and mass $\sim 0.03~\msol$. From the observed kinematics Hayashi et al. derive an accretion rate $\sim 9~\msol~\rm{Myr}^{-1}$ at $r \sim 700~\rm{AU}$. The embedded protostar L1551-IRS5 has a mass $\sim 0.5~\msol$ and is surrounded by a disk with radius $\sim 700~\rm{AU}$ and mass in the range $3.9 \times 10^{-2} - 8.1 \times 10^{-2}~\msol$. Ohashi et al. deduce an accretion rate in the range $13 - 26~\msol~{\rm{Myr}}^{-1}$ at $r \sim 600~\rm{AU}$. These values are comparable to our model results for $\delt_5 \simlt \delt \simlt \delt_6$ (corresponding to $0.2~\msol \simlt \mcent \simlt 1~\msol$; see Fig. $6a$). During this period, $ 6~ \msol~{\rm{Myr}}^{-1} \simlt \mdot \simlt 9~ \msol~{\rm{Myr}}^{-1}$ for $r \simgt 500~\rm{AU}$ (see Fig. $6g$). [Note, however, that the temperatures of HL Tauri and L1551-IRS5 are in the range $15 - 50~{\rm K}$, which is greater than our assumed value of 10 K and should lead to higher values of $\mdot$ and $\mcent(t)$; e.g., Shu et al. 1987. For a discussion of how quantities scale with temperature in our models, see Basu \& Mouschovias 1994.] In Figure 7 we show the mass ($M - \mcent$) of the gas surrounding the central point mass in our typical model as a function of $r/\rzero$ for the same seven times $\delt_j$ as in Figure 6. (Taken together, Figs. $6a$ and 7 may be taken to represent the evolution of a protostar from a Class 0 to a Class I object.) For times $\simgt \delt_5$, the surrounding disk mass spans the range $0.01 - 0.1~\msol$ for $r \simgt 500~\rm{AU}$, which agrees with the ${}^{13}{\rm{CO}}$ disk masses of HL Tauri and L1551-IRS5 cited above. Finally, we note that the age of the oldest part of the molecular outflow from L1551-IRS5 is estimated to be $\sim 10^5~\rm{yr}$ (Bachiller, Tafalla, \& Cernicharo 1994). This age is consistent with the time $\delt$ needed for the central mass in our typical model to become $\simgt 0.3~\msol$ (see Fig. $2e$). Another observational consequence of our model is the magnetic field structure in the core after PMF. Figures $6b$ and $6d$ show that $\BrZ \approx \Beq$ inside the core for the radius range $ 2 \times 10^{-4} \simlt r/\rzero \simlt 4 \times 10^{-3}$. Hence, there is significant curvature of field lines (though, as discussed in \S~3.3, there is less bending than there would be if the field had remained frozen into the neutrals), with bending angles $\theta_B \approx \arctan(\BrZ/\Beq)$ in the range $20^{\circ}- 50^{\circ}$ for $180~ {\rm AU} \simlt r \simlt 3.5 \times 10^3~{\rm{AU}}$. \footnote{In agreement with the results of CM94 and CM95, we find that the magnetic tension force is generally not negligible in comparison with the magnetic pressure force at any radius (in either the core or the envelope), both before and after PMF.} This type of field geometry may be described roughly as having an hourglass shape. In fact, sub-mm polarimetry of the cores of W3 IRS5 (Greaves, Murray, \& Holland 1994), Mon R2 (Greaves, Holland, \& Murray 1995), and OMC-1 (Schleuning 1998) find field geometries suggestive of an hourglass shape on sub-parsec scales. In contrast, field lines remain essentially straight and parallel in the magnetically supported envelope ($r > 0.1~\rm{pc}$), even after PMF. This agrees with polarimetric observations of molecular clouds that indicate well-ordered fields on these scales (e.g., Hildebrand, Dragovan, \& Novak 1984; Hildebrand 1989, 1996; Novak et al. 1989; Kane et al. 1993; Hildebrand et al. 1995). A unique prediction of our model is the large ion--neutral drift speed that occurs during the post-PMF epoch. As shown in Figure $6f$, effective ambipolar diffusion following PMF yields $\vd \approx |\vn| \gg C$ in the inner regions of the core. In our typical model we find $\vd \simgt 1~\rm{km}~{\rm s}^{-1}$ for $\delt \simgt 2 \times 10^4~\rm{yr}$ on scales $r \simlt 2 \times 10^{-4} \rzero \simeq 180~\rm{AU}$. Large drift speeds between neutrals and ions (such as $\HCOp$, $\rm{HCN}^+$, $\rm{DCO}^+$, to name but a few) on these scales are therefore expected in our model. Detection of such drift speeds (through high-resolution observations of HCN or $\HCOp$, say) could be used to observationally confirm our model results and to distinguish them from those of nonmagnetic collapse calculations (e.g., Shu 1977, Hunter 1977, Foster \& Chevalier 1993) or of magnetic collapse models that do not account for the effect of ambipolar diffusion (e.g., Tomisaka 1996; Li \& Shu 1997). \subsection{Features of the Hydromagnetic Shock} The properties of hydromagnetic shocks in partially ionized gases have been developed extensively by many other authors (e.g., Mullan 1971; Draine 1980; Chernoff 1987; Roberge \& Draine 1990; Draine \& McKee 1993; Smith \& Mac Low 1997). Because the ion {\Alf} speed $v_{\rm{A,i}} = \Beq/(4 \pi \mi \nni)^{1/2} = (\mn/\mi)^{1/2} (\nn/\nni)^{1/2} \van$ is much larger than $\van$, $|\vn|$, and $|\vi|$ in our model, we expect the outward-propagating shock that develops after PMF to have a magnetic precursor. This is indeed what we find in our simulation: the jump in the ion speed $\vi$ and the magnetic field strength $\Beq$ typically occurs at a distance of 1 to 3 computational mesh spacings further away from the symmetry axis than the jump in the neutral infall speed $\vn$ and the column density $\sign$. The displacement between the locations of the head of the disturbance in the neutral fluid and in the plasma and magnetic field decreases at later times. Examination of Figures $5b$, $5c$, and $6e$ reveals that, in the reference frame of the shock, the preshock infall speeds are supersonic, and the postshock speeds are also supersonic or just slightly subsonic. Similarly, the ion infall speeds (in the frame of the shock) are much less than the ion {\Alf} speed. Therefore the shock we observe in our model is probably best classified as being C-type. \footnote{A C-type shock is characterized by neutral velocities that (in the shock frame) remain supersonic throughout. Hence the shock in our simulation cannot be strictly of this type when the downstream neutral speed is subsonic. In that case a viscous (J-type) subshock may form (Draine \& McKee 1993), although, as noted by Li \& McKee (1996), turbulent diffusivity behind a real shock could plausibly keep the postshock flow supersonic and thereby obviate the need for such a subshock.} Making the approximation that in the vicinity of the shock the predominant magnetic stress is that due to the magnetic pressure gradient, the ion force equation becomes \begin{equation} \label{ionforceq} \frac{\sign}{\tni} \vd = - \frac{Z}{4 \pi} \frac{\partial \Beq^2}{\partial r} \end{equation} (see eqs. [28c] and [51] in CM93, which contain additional terms, involving in particular the magnetic tension force, that could be used to refine the following simple estimate). This yields an approximate shock width \leteq \begin{eqnarray} \label{shkapproxa} \delshk &\approx& \frac{\Bequ^2}{4 \pi \sign} \frac{\tni Z}{\vd} \left[\left(\frac{\Beqd}{\Bequ}\right)^2 -1 \right] = C \tni \frac{\left(\vanu/C \right)^2}{2 \left(\vd/C\right)} \left[\left(\frac{\Beqd}{\Bequ}\right)^2 -1 \right] \\ \label{shkapproxb} &=& 7.8 \times 10^{13} \frac{\left(T/10~{\rm K}\right)^{1/2}\left(\vanu/C\right)^2}{\left(\mn/2.33~{\rm{a.m.u.}}\right)^{1/2}\left(\nni/0.1~\cc\right)\left(\vd/C\right)} \left[1 + 0.067 \frac{\left(\mHII/2~{\rm{a.m.u.}}\right)}{\left(\mi/30~{\rm{a.m.u.}}\right)}\right] \nonumber \\ &&\hspace{5em}\times \left[\left(\frac{\Beqd}{\Bequ}\right)^2-1\right]~\rm{cm}\ , \end{eqnarray} \beq where $\Bequ$ and $\Beqd$ are the values of $\Beq$ upstream and downstream of the shock, and $\vanu$ is the upstream {\Alf} speed. In deriving the last equality of equation (\ref{shkapproxa}) we have used the relation $\sign= 2 \rhon Z$; equation (\ref{tnieq}) and the relation $C= \left(\kB T/\mn\right)^{1/2}$ have been used in deriving equation (\ref{shkapproxb}). At the time $\delt_6$ the shock front is located at $r \simeq 3.9 \times 10^{-3} \rzero = 5.2 \times 10^{16}~{\rm cm}$, and, in the vicinity of the front, $\nni \simeq 10^{-2}~\cc$, $\vanu \simeq C$, $\Beqd/\Bequ \simeq 4.6$ (see Fig. $6b$), and $\vd \simeq 0.4 C$ (see Fig. $6f$). For these values our rough estimate for the shock width given by equation (\ref{shkapproxb}) yields $\delshk \simeq 4.3 \times 10^{16}~{\rm{cm}}$. Examination of Figure $6e$ at the time $\delt_6$ reveals that the shock has an actual width $\delshk \simeq 2.1 \times 10^{-3} \rzero = 2.8 \times 10^{16}~\rm{cm}$. (Thus $\delshk/r_{\rm shk} \approx 0.5$ at that time, so the thin-shock approximation that underlies the estimate [17] is marginally satisfied.) As noted in \S~1, Li \& McKee (1996) proposed that a hydromagnetic shock would form in a collapsing core as a result of the decoupling of flux from the ion and neutral fluids because of Ohmic dissipation (a process that becomes important in regions of density $\nn \gg 10^{11}~\cc$; see \S~2.1). They argued that the accumulating flux diffuses outward to regions of lower density, where improved coupling with the matter causes it to present an obstacle to the infalling neutral gas --- thereby giving rise to a hydromagnetic shock. While our numerical results have confirmed Li \& McKee's basic shock-formation scenario, they have revealed that ambipolar diffusion, which mediates the shock, can also, following PMF, halt the inward advection of magnetic flux on scales ($r \simgt 5~\rm{AU}$) where Ohmic dissipation is not important. In other words, our results have shown that, during the post-PMF epoch, the field--matter decoupling that drives the hydromagnetic shock is due to ambipolar diffusion alone and does not depend on the effect of Ohmic dissipation at $r < \rinner$. Because of the similarity in the basic shock-formation mechanism, it is of interest to compare our detailed simulation results with the predictions of the (simplified and analytic) shock model of Li \& McKee (1996). From their requirement that the magnetic pressure of the shock balance the ram pressure of the neutrals, which were assumed to be freely-falling into the shock, Li \& McKee derived relations for the shocked magnetic field strength and the shock location (see their eqs. [7] and [8]) in terms of the accretion rate $\mdot$, the flux-to-mass ratio in units of the critical value for collapse (dubbed $\epsilon$ in their paper), a parameter related to the logarithmic gradient of the magnetic field (dubbed $\chi$), the ratio $Z/r$ of the local gas scale height and the radius (dubbed $h$), and the protostellar mass (denoted by $m_{\ast}$). At the time $\delt_6$ we have at the location of our shock $\dot{M} \approx 9~\msol~{\rm{Myr}}^{-1}$, $\epsilon \approx 0.9$, $\chi \approx 2$, $h \approx 0.3$, and $m_{\ast}=\mcent(\delt_6) \approx 1~\msol$. Inserting these values into their equations (7) and (8) yields a shocked magnetic field strength $\sim 630~\mu{\rm G}$ and a shock radius $\sim 2.1 \times 10^3 ~{\rm{AU}}$; by comparison, in our model the shocked magnetic field strength at that time is $\Beqd \approx 330~\mu\rm{G}$ and the shock radius is $r_{\rm shk} \approx 3.5 \times 10^3~\rm{AU}$ (see Figs. $6b$ and $6e$). The main reason why the analytic expression overestimates the numerically calculated magnetic field strength is that, contrary to the assumption of Li \& McKee, the preshock neutrals are not in free fall but, rather, are strongly decelerated by magnetic forces (in our simulation we find that the preshock acceleration of the neutrals is reduced to $\simeq 0.25 \gr$). The corresponding reduction in the preshock ram pressure leads to a lower value of the postshock field amplitude, with a further reduction in the calculated field strength brought about by the contribution of magnetic tension (ignored in the analytic estimate) to the total magnetic force. Since the analytic estimate of $r_{\rm shk}$ is based on relating the postshock field strength to the total magnetic flux inside the shock, the overestimate of the field strength naturally results in an underestimate of the shock radius. Despite these discrepancies, Li \& McKee's analytic representation of the shock parameters provides a decent approximation to the results of our numerical calculation. Comparison of the column density $\sign$ (see Fig. $6c$) and neutral infall speed $\vn$ (see Fig. $6e$) behind the shock shows qualitative agreement with Figures $2a$ and $2d$ of Li \& McKee (1996), including the free-fall behavior near the central protostar. In the pre-shock region, however, our infall speeds are smaller than theirs because of their assumption of free fall upstream of the shock: Li \& McKee typically overestimate $\vnu$ by a factor $\sim 2$. As a result, the Mach number of the shock relative to the upstream gas (and thus the shock strength) is greater in their model than in ours (see Fig. $5c$) by a similar factor. We have not compared our results for the magnetic field structure behind the shock with those of Li \& McKee on account of the fact that their system of MHD equations was not closed (it did not include the magnetic induction equation), so that they were unable to calculate the magnetic field with any accuracy (see their Fig. $2b$). \subsection{Stability of the Core Against Magnetic Interchange} Our simulation has revealed that rapid ambipolar diffusion occurs behind the outward-propagating HMD. The effect that this has on the mass in the flux tubes downstream of the HMD can be seen in Figure $8a$, which shows the local mass-to-flux ratio $d M/d \phiB = \sign/\Beq$ (normalized to the critical value for collapse) as a function of $r/\rzero$ for the same seven times $\delt_j$ as in Figure 6. For $\delt > \delt_1$ a local minimum in $\sign/\Beq$ appears after the passage of the HMD. Hence, there is a region behind the HMD for which $d (\sign/\Beq)/dr > 0$. This is a necessary condition for the onset of an interchange instability (e.g., Spruit \& Taam 1990; Lubow \& Spruit 1995; Spruit, Stehle, \& Papaloizou 1995). Li \& McKee (1996) speculated that such a situation could arise in the wake of a hydromagnetic shock in a collapsing core and suggested that it would act as source of turbulence in the postshock region of the inflow. \footnote{Li \& McKee (1996) also noted that the shock may be unstable to the Wardle instability, which involves ions collecting in magnetic field ``valleys'' (Wardle 1990). However, the shock will be immune to this instability if the ion density is determined by the local chemical reaction balance (as assumed in our calculation) rather than by the divergence of the ion mass flux.} Blaes \& Balbus (1994) found that the magnetic shearing instability in differentially rotating disks could be stabilized if the disk is weakly ionized. This will also be the case for the interchange instability in a weakly ionized disk: instability is possible only if the growth rate $\gaminst$ and the neutral--ion collision time $\tni$ satisfy the condition $\gaminst \tni < 1$. This condition reflects the fact that there has to be sufficient collisional coupling between the ion and neutral fluids for a magnetic interchange instability to grow in the neutrals; otherwise the instability is damped. Calculation of $\gaminst$ in our model is hampered by the fact that previous studies of interchange instability have been carried out only for disks that are in hydrostatic equilibrium, with exact balance between magnetic and gravitational forces. We can apply the results of these studies to our model only if the region behind the shock where $d (\sign/\Beq)/dr > 0$ is effectively in quasi equilibrium, with approximate balance between gravitational and magnetic forces, and with infall speeds $|\vn| \ll (r |\gr|)^{1/2}$ ($\simeq$ the free-fall speed). In our model, the magnitude of the acceleration of the neutrals in this region of the core does not become $\simlt 0.1 | \gr |$ until times $\sim \delt_6$; hence, approximate equilibrium between magnetic and gravitational forces is valid only at these later times. Spruit \& Taam (1990) found that the growth rate for the most unstable linear interchange modes is \begin{equation} \label{growtheq} \gaminst = \left( \frac{\Beq \BrZ}{2 \pi \sign} \frac{d}{dr} \ln \frac{\sign}{\Beq} \right)^{1/2} . \end{equation} The product $\gaminst \tni$ for the region of the core susceptible to interchange instability is shown in Figure $8b$ as a function of $r/\rzero$ at the time $\delt_6$. We also plot (Fig. $8c$) the product $\gaminst \tau_{\rm{kin}}$, where $\tau_{\rm{kin}} \equiv r/|\vn|$ is the kinematical timescale, as a function of $r/\rzero$ for the same region of the core and time. If this product is $ < 1$, the unstable modes will be ``swept up'' by the infalling gas before they have time to grow. It is evident from these figures that $\gaminst \tni < 1$ and $\gaminst \tau_{\rm{kin}} > 1$ for the region of the core susceptible to magnetic interchange. This means that there is sufficient collisional coupling between the ions and neutrals, and that the instability will grow before being swept along with the neutrals. Hence, this region of the core may be interchange unstable. An instability of this type would enhance the tansfer of gas with a high mass-to-flux ratio to the center (e.g., Spruit \& Taam 1990), and, as noted by Li \& McKee (1996), might also lead to the development of turbulence that could increase the field diffusivity in the postshock gas. However, the onset and development of this instability can only be studied by means of a fully 3-D simulation. \subsection{Implications to the Magnetic Flux Problem} The magnetic flux problem in star formation has to do with the fact that the magnetic flux of a $1\msol$ blob of matter in the diffuse interstellar medium is typically several orders of magnitude greater than the flux of a $1\msol$ protostar. Such a blob of matter would therefore have to get rid of most of its flux before becoming a star. Ambipolar diffusion has long been suggested as a means by which the magnetic flux problem could be resolved (e.g., Mestel \& Spitzer 1956; Mouschovias 1978; Paleologou \& Mouschovias 1983; Nakano 1984; Mouschovias, Paleologou, \& Fiedler 1985). In general, these earlier studies focused primarily on the role of ambipolar diffusion and the magnetic flux problem for the pre-PMF phase of protostellar evolution. \footnote{On the basis of a consideration of the timescales for ambipolar diffusion and Ohmic dissipation at high densities, Nakano \& Umebayashi (1986b) suggested that significant flux loss could only take place (primarily by Ohmic dissipation, according to their estimates) during the dynamical phase of core collapse. Lizano \& Shu (1989) similarly concluded that the resolution of the protostellar magnetic flux problem must occur during the dynamical stage of core evolution: using the quasi-static approximation (valid for $\nnc \simlt \rm{a~few} \times 10^4~\cc$; Fiedler \& Mouschovias 1993; CM94; Basu \& Mouschovias 1994) to calculate the contraction of a slightly subcritical molecular cloud, they found that only a small amount of flux is lost by ambipolar diffusion from the central flux tubes before runaway collapse is initiated.} While ambipolar diffusion does indeed reduce the flux-to-mass ratio during that phase, the flux contained within a $1\msol$ region of a molecular cloud core is still much larger at the time of PMF than typical protostellar fluxes. Specifically, CM94 and CM95 found that the central $1\msol$ flux tube within their cores had a total magnetic flux $\sim 10^{30}~\rm{Mx}$ (consistent with our results in \S~3.2) during the pre-PMF dynamical collapse phase of the typical model. This value represents a reduction by a factor $\sim 5.6$ of the flux associated with that mass before the onset of ambipolar diffusion. Nevertheless, it greatly exceeds the plausible upper limit ($\sim 6 \times 10^{26}~\rm{Mx}$) on the flux of a solar-mass protostar (estimated assuming an average surface field of 10 kG and a stellar radius of $10^{11}~{\rm cm}$; see Li \& McKee 1996). We have shown in this paper that the rate of ambipolar diffusion is strongly increased during the post-PMF epoch of star formation. It is therefore of interest to examine the implications of our simulation results to the magnetic flux problem. As discussed in \S~2.2, in our calculations we only consider the core regions at radii $r \simgt \rinner$ ($\simeq 7.3~\rm{AU}$ for our typical model), where ambipolar diffusion is the dominant mechanism of flux loss. Initially, the magnetic flux contained within $\rinner$ is $\phicent(\delt=0) = 6.7 \times 10^{26}~\rm{Mx}$. As shown in Figure $3a$, the central flux increases before the onset of rapid ambipolar diffusion. This continues to the time $\delt \approx 10^3~\rm{yr}$. For $\delt > 10^3~\rm{yr}$, ambipolar diffusion prevents further advection of flux from $r > \rinner$ into the central sink, and $\phicent$ changes very little after this time. By the time $\delt \approx 10^5~\rm{yr}$, when $\mcent \approx 1\msol$, $\phicent \approx 5 \times 10^{27}~\rm{Mx}$. This represents a decrease of over two orders of magnitude relative to the flux associated with this mass at the time of PMF. While this value is still about an order of magnitude higher than our adopted upper limit on the protostelar flux, the discrepancy is now much lower than previous estimates of ambipolar diffusion have indicated. The important conclusion from our work is thus that {\em ambipolar diffusion in contracting molecular cloud cores can in principle contribute significantly to the resolution of the magnetic flux problem} by reducing the magnetic flux brought into a solar-mass protostar by a factor $\simgt 10^3$. The new, and somewhat surprising, result is that {\em most of this reduction occurs after PMF.} The remainder of the protostellar magnetic flux could possibly be extracted from the infalling mass through Ohmic dissipation within $\rinner$, although refreezing of the magnetic field into the matter, brought about by collisional reionization at densities $\nn \simgt 10^{14}~\cc$ (e.g., Pneuman \& Mitchell 1965; Nakano \& Umebayashi 1986b; Li \& McKee 1996), as well as anomalous diffusivity (e.g., Norman \& Heyvaerts 1985) operating in the reionized gas, could complicate the issue. Another complicating factor is the strong likelihood that much of the mass and flux carried into the protostar pass through a rotationally supported, circumstellar accretion disk of size $\gg \rinner$ (e.g., Lubow, Papaloizou, \& Pringle 1994; Reyes-Ruiz \& Stepinski 1996; Li 1996). It is also conceivable that magnetic flux is brought to the vicinity of the protostar but excluded from its interior by turbulent diffusivity associated with convection. Since the region within $\rinner$ was excluded from our calculation, we do not pursue this topic any further in this paper. | 98 | 4 | astro-ph9804113_arXiv.txt |
9804 | hep-ex9804007_arXiv.txt | With an effective telescope area of order $10^4$~m$^2$, a threshold of $\sim$50~GeV and a pointing accuracy of 2.5~degrees, the AMANDA detector represents the first of a new generation of high energy neutrino telescopes, reaching a scale envisaged over 25 years ago. We describe its performance, focussing on the capability to detect halo dark matter particles via their annihilation into neutrinos. | \unskip High energy neutrino telescopes are multi-purpose instruments; their science mission covers particle physics, astronomy and astrophysics, cosmology and cosmic ray physics. Their deployment creates new opportunities for glaciology and oceanography, possibly geology of the earth's core\cite{pr}. Astronomy with neutrinos does have definite advantages. They can reach us, essentially without attenuation in flux, from the largest red-shifts. The sky is, in contrast, partially opaque to high energy photons and protons because of energy-loss suffered in interactions with infrared light, CMBR photons and radio waves\cite{cronin}. They do not reach us from distances much larger than tens of megaparsecs once their energy exceeds thresholds of 10~TeV for photons and $5\times 10^7$~TeV for protons. (Below this energy charged protons do not point back to their sources.) The drawback is that neutrinos are difficult to detect: the small interaction cross sections that enable them to travel without attenuation over a Hubble radius, are also the reason why kilometer-scale detectors are required in order to capture them in sufficient numbers to do astronomy\cite{halzenkm}. Some opportunities may, however, be unique. If, for instance, the sources of the highest energy cosmic rays are beyond $10^2$~Mpc, conventional astronomy is unlikely to discover them. Some science missions do not require a detector of kilometer size. The best opportunities to search for halo dark matter are, in fact, associated with the present instrument which, while smaller in telescope area than the planned extension of AMANDA to ICE3 (ICECUBE), has a lower threshold. At this meeting, the capability of neutrino telescopes to discover the particles that constitute the dominant, cold component of the dark matter is of special interest. The existence of the weakly interacting massive particles (WIMPs) is inferred from observation of their annihilation products. Cold dark matter particles annihilate into neutrinos; {\it massive} ones will annihilate into {\it high-energy} neutrinos which can be detected in high-energy neutrino telescopes. This so-called indirect detection is greatly facilitated by the fact that the earth and the sun represent dense, nearby sources of accumulated cold dark matter particles. Galactic WIMPs, scattering off nuclei in the sun, lose energy. They may fall below escape velocity and be gravitationally trapped. Trapped WIMPs eventually come to equilibrium and accumulate near the center of the sun. While the WIMP density builds up, their annihilation rate into lighter particles increases until equilibrium is achieved where the annihilation rate equals half of the capture rate. The sun has thus become a reservoir of WIMPs which we expect to annihilate mostly into heavy quarks and, for the heavier WIMPs, into weak bosons. The leptonic decays of the heavy quark and weak boson annihilation products turn the sun and earth into nearby sources of high-energy neutrinos with energies in the GeV to TeV range. Figure~1 displays the neutrino flux from the center of the earth calculated in the context of supersymmetric dark matter theories\cite{scopel}. The direct capture rate of the WIMPs in germanium detectors is shown for comparison. Contours indicate the parameter space favored by grand unified theories. Most of this parameter space can be covered by improving the capabilities of existing detectors by 2 orders of magnitude. Existing neutrino detectors have already excluded fluxes of neutrinos from the earth's center of order 1~event per $1000 \rm~m^2$ per year. The best limits have been obtained by the Baksan experiment\cite{suvorova}. They are already excluding relevant parameter space of supersymmetric models. We will show that, with data already on tape, the AMANDA detector will have an unmatched discovery reach for WIMP masses in excess of 100~GeV. \begin{figure}[t] \centering \hspace{0in}\epsfxsize=4.5in\epsffile{fig1.eps} \caption{Direct and indirect detection rates (for neutrinos from the center of the earth in the figure shown) of cold dark matter particles predicted by supersymmetric theory. Grand unified theories favor the parameter space indicated. Part of it is already excluded by present experiments as indicated by the horizontal and vertical lines.} \end{figure} The potential of neutrino telescopes as dark matter detectors has been documented in detail\cite{kamionkowski}. With a sensitivity which increases with the WIMP mass, they are complementary to direct, cryogenic detectors. They can detect WIMPS beyond the kinematic limits of the LHC: about 500~GeV for neutralinos. A striking way to illustrate their potential is to use the possible detection\cite{belli} in the DAMA NaI detector in the Gran Sasso tunnel as an example. If their seasonal variation is indeed evidence for WIMPS, observation of a signal in an exposure of 4500 kg\,days requires a WIMP-nucleon cross section of $10^{-42}{\sim}10^{-41}$~cm$^2$ for a WIMP mass of $50{\sim} 150$~GeV. This information is sufficient to calculate their trapping and annihilation rate in the sun and earth. Both will be a source of, on average, 100 neutrinos per year of WIMP origin in the existing AMANDA detector with an effective area of $10^4$~m$^2$. The exact rate varies with the mass of the WIMP. | 98 | 4 | hep-ex9804007_arXiv.txt |
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9804 | astro-ph9804107_arXiv.txt | We have computed line profiles from self-gravitating toroids around black holes. The specific angular momentum of the toroids is assumed to be constant in space. The images of the toroids show peculiar feature in the rear side of the black holes. Concerning the line profiles, the red wing extends to the very small frequency region because the location of the inner edge is rather near the event horizon of the black hole and consequently the velocity of the inner edge of toroids can be faster than that of the Kepler disks. | Recently, strong evidence for existence of supermassive black holes in the central regions of galaxies has been found from observations of electromagnetic waves with various wavelengths ranging from the radio waves to X-rays (see e.g. Miyoshi et al. 1995 for the radio observation; Ford et al. 1994, Harms et al. 1994 for the optical observation; Tanaka et al. 1995 for the X-ray observation). Radio and optical observations have shown that there exist very rapidly rotating gaseous disks in the central regions of galaxies. Since the high velocity of a disk implies the existence of a large amount of mass within a region of a very small size, it is widely considered that there are supermassive black holes with masses of $10^7 \sim 10^9 M_{\sun}$ at the centers of galaxies. However, such observations do not reveal the nature of black holes because the size of the observed region is still too large to get detailed information about the black holes. Contrary to these optical and radio analyses, recent X-ray observations have brought us important information about gravitational fields very near the massive objects or black holes. By using the ASCA satellite, broad iron emission lines have been detected in active galaxies (see e.g. Fabian et al. 1994; Mushotzky et al. 1995). In particular, Tanaka et al.~(1995) observed the Seyfert 1 galaxy, MCG--6--30--15, and discovered a very broad and skewed iron emission line. The broadness and skewness of the line profile can be explained only by assuming that the inner edge of the accretion disk is located very near the event horizon of the central black hole (Tanaka et al. 1995; Fabian et al. 1995). It implies that X-ray observations can be used to understand the very vicinity of the black holes. It is important to know the gravitational field near the event horizon because it is the key feature to distinguish a Schwarzschild black hole from a Kerr black hole. Consequently we have reached a stage to investigate very strong gravitational fields of black holes and to be able to test the validity of general relativity. However, since the emission line profiles depend not only on the gravitational fields but also on the structures of the accretion disks, it is not easy to determine the type of the black hole, i.e. whether the central black hole is of a Schwarzschild type or of a Kerr type. In fact, concerning the nature of the black hole at the center of MCG--6--30--15, various discussions have not settled down to a unique interpretation yet (see e.g. Tanaka et al. 1995; Iwasawa et al. 1996; Dabrowski et al. 1997; Bromley et al. 1997; Reynolds \& Begelman 1997; Bromley et al. 1998). Therefore, in order to get a consistent picture of black hole -- accretion disk systems, we have to pursue much more investigations theoretically as well as observationally. Concerning the theory about the spectra of accretion disks, Cunningham~(1975) was the first to formulate the problem for Kerr black holes and obtained theoretical spectra from accretion disks around black holes (see also Cunningham 1976). Gerbal \& Pelat~(1981) investigated lines from a ring around a black hole and found double-peaked asymmetric profiles. From the end of 80's, observations of lines in the X-ray spectrum stimulated many authors to study line profiles of accretion disks as well as disk structures (e.g. Nandra et al. 1989; Fabian et al. 1989; Kunieda et al. 1990; Kojima 1991; Laor 1991; Chen \& Halpern 1989, 1990). However, in almost all theoretical studies mentioned above, investigations of the emission line profiles have been done by assuming that disks are geometrically thin and that only direct photons are observed. Some authors have studied the effect of multiple images (Luminet 1978; Bao, Hadrava \& {\O}stgaard 1994) and that of self-eclipse due to toroidal configurations (Bao \& Stuchlik 1992; Kojima \& Fukue 1992). For toroidal configurations the rotation law of the toroid is not always that of the Kepler rotation because of the presence of the pressure within the toroid. In fact, Kojima \& Fukue~(1992) employed a variety of rotation laws. However, their analysis was done in the framework of Newtonian gravity. Therefore, quantitative treatment of non-Keplerian toroids in general relativity has not been carried out yet. Furthermore, in some situations, self-gravity of disks or toroids plays an important role for the structures of disks or toroids. In particular, massive neutron toroids around neutron stars or black holes have been proposed as possible sources of $\gamma$-ray bursts (e.g. Paczy\'nski 1991; Narayan, Paczy\'nski \& Piran 1992; Jaroszy\'nski 1993; Witt et al. 1994). Although there might be little chance to observe line profiles from such systems even if exist, it would be interesting to study the effect of self-gravity of toroids on the line profiles. Such self-gravitating disks were investigated by Karas, Lanza \& Vokrouhlicky~(1995). However, they have studied very light thin disks whose mass is less than several percent of the mass of the black hole. From the theoretical point of view, it would be interesting to investigate more massive disks or toroids as well as the effect of geometrical thickness. By treating massive toroids, the central objects are no more of Schwarzschild type nor Kerr type black holes because of the gravitational effect of the self-gravitating massive toroids on the black holes (Nishida \& Eriguchi 1994). Moreover, gravitational effect of massive toroids may bring some differences to line profiles. Therefore, in this paper, we will consider self-gravitating toroid -- black hole systems and study their effect on emission line profiles. | In this paper we have constructed toroid -- black hole systems in general relativity. We have computed photon trajectories in the numerically obtained gravitational fields and investigated the images of the toroids as well as the line profiles from the toroids. In these computations, we have assumed that the observer is located very near the black hole because the metric in the whole space has not been calculated. Therefore, the quantitative values in this paper will be changed a little if the observer is located at infinity. However, the important thing is not to obtain exact quantitative values but to know characteristic feature which appears only by introducing thickness and self-gravity numerically exactly. In this sense, our results serve as representative ones for self-gravitating toroids around black holes. As discussed in Introduction, the observation of MCG--6--30--15 has given us time varying line profiles whose interpretation has not been clarified yet (Tanaka et al. 1995; Iwasawa et al. 1996; Dabrowski et al. 1997). Some authors have proposed models for the system which would explain the observations (Dabrowski et al. 1997; Bromley et al. 1997; Bromley et al. 1998). However, since there are many parameters about the disk structures and the X-ray sources, it is very difficult to obtain a unique solution for the system. This can be also seen from our result. For our Model B with the emissivity indices $s = -1$ and $s = -4$, the line profiles are shown in Fig.~\ref{comparison} together with the observational data. As seen from this figure, the tendency of the observational data seems to be roughly explained by the change of the emissivity index of Model B. Thus, at the present stage, we can only say that since the model cannot be uniquely determined, we have to get more accurate observational data to clarify the spacetime of the central region of galaxies. \begin{figure*} \begin{tabular}{cc} (a)&(b)\\ \epsfile{file=tanaka.eps,width=0.5\textwidth} &\epsfile{file=iwasawa.eps,width=0.5\textwidth}\\ \end{tabular} \caption{ Same as Fig. 4 but for the model with $\theta_0 = 30 \degr$ and emissivity indices (a) $s=-1$ (solid line) and (b) $s=-4$ (solid line). The observational data of the ASCA are plotted by points with error bars for the data of Tanaka et al.~(1995) (a) and that of Iwasawa et al.~(1996) (b). } \label{comparison} \end{figure*} | 98 | 4 | astro-ph9804107_arXiv.txt |
9804 | astro-ph9804331_arXiv.txt | We report the detection of the GRB 971214 counterpart in the near-infrared by means of two images in the K$^{\prime}$-band taken at Calar Alto only $\sim$3.5 and $\sim$5 hours after the gamma-ray event. We detect the transient at K$^{\prime}$=$18.03\pm0.18$ and K$^{\prime}$=$18.00\pm0.22$ respectively. Our data seem to indicate the existence of a plateau with duration $ 1.5 \leq T \le 6.7 $ hours (between 3.5 and 10.2 hours after the high-energy event). Moreover the power-law decline should be steeper than the one given by the index $\alpha_{\rm K^{\prime}}=0.45$. There is also a change in the slope of the broad-band spectrum at some wavelength between the J and K$^{\prime}$ bands (possibly around the H-band). | After 31 years since the discovery of gamma-ray bursts (GRBs), the origin of such brief gamma-ray flashes remains unknown. The observed isotropy of GRBs in the sky, could only be explained by theoretical models where GRBs originate either in a extended halo around the Galaxy or arise from sources at cosmological distances. Before the launch of the BeppoSAX satellite, the poor localization capability of the GRB detectors made the searches at other wavelengths unfruitful. The breakthrough took place in 1997 when the first X-ray afterglows were observed by the BeppoSAX, RXTE, ROSAT and ASCA satellites (Costa et al. 1997, Heise et al. 1997a, Marshall et al. 1997, Greiner et al. 1998, Murakami 1998). They were able to localize the fading X-ray emission that followed the more energetic gamma-ray photons once the GRB event had ended. This emission (the afterglow) extends to longer wavelengths, and the good accuracy in the position determination by BeppoSAX (typically $1^{\prime}$ radius error boxes) has led to the discovery of the first optical counterparts for GRB 970228 (van Paradijs et al. 1997, Guarnieri et al. 1997), and GRB 970508 (Bond 1997, Djorgovski et al. 1997, Castro-Tirado et al. 1998), greatly improving our understanding of these puzzling sources. The measurement of the redshift for the GRB 970508 optical counterpart (Metzger et al. 1997) has established that one GRB, maybe all, lie at cosmological distances. GRB 971214 is the third GRB with a known optical counterpart. It was detected by the BeppoSAX Gamma-ray Burst Monitor (GBM, Frontera et al. 1997) on Dec 14.97 1997, as a 25 s long-structured gamma-ray burst. Simultaneous to the detection of the GBM, the Wide Field Cameras (WFC, Jager et al. 1997) on board BeppoSAX provided an accurate position (a $3^{\prime}.9$ radius error box at a 3$\sigma$ confidence level, Heise et al. 1997b) that allowed deep optical, infrared and radio observations. The position was also consistent with the one given by the all-sky monitor on RXTE (Bradt et al. 1993) and by the BATSE/Ulysses annulus (Kippen et al. 1997). When BeppoSAX pointed its Narrow-Field Instruments (NFI) to the GRB position, on Dec 15.25 ($\sim$ 6.5 hours after the burst), a previously unknown variable X-ray source was found inside the WFC error box (Antonelli et al. 1997) which was identified as the X-ray afterglow of GRB 971214. Soon after, Halpern et al. (1997) reported the presence of a fading object inside the WFC GRB error box, based on two I-band images separated 24 hours. The object was afterwards confirmed as the counterpart of GRB 971214 by means of additional observations at other wavelengths: R-band (Castander et al. 1997, Diercks et al. 1998), I-band (Rhoads 1997) and J-band (Tanvir et al. 1997). No detections in the K-band were reported in the literature, although observations performed on Dec 15.54 imposed an upper limit of K $>$ 18.5 (Garcia et al. 1997). As it will be explained later, this upper limit will be used to constraint the power-law index $\alpha_{\rm K^{\prime}}$ and the position of the possible maximum of the light curve. We report here the detection of the GRB 971214 counterpart in the near infrared (IR) by means of two K$^{\prime}$-band images taken at Calar Alto on Dec 15.12 and 15.18 (mean observing time, only $\sim$3.5 and $\sim$5 hours after the gamma-ray event). The second image is almost simultaneous to the beginning of the observations performed by the BeppoSAX narrow-field instruments. We discuss whether our observations are in agreement with the extrapolation of the power-law seen at other bands in later epochs. | We have detected the GRB 971214 near-IR counterpart $\sim$3.5 hours and $\sim$5 hours after the gamma-ray event which enables to conclude that: i) a magnitude difference $\Delta$K${^{\prime}}=-0.03\pm 0.28$ is derived from our measurements, whereas $\Delta$K${^{\prime}}=0.464$ would be expected assuming a power-law decay with index $\alpha_{\rm K^{\prime}}=1.2$ (similar to the one observed at optical wavelengths). This implies a deviation of $1.7\sigma$. If the assumed power-law index $\alpha_{\rm K^{\prime}}$ were 1.4, then the rejection level would be $2.0\sigma$. Thus, our measurements suggest a rising or a flat light curve segment with a duration $ 1.5 \leq T \le 6.7 $ hours (between 3.5 and 10.2 hours after the burst). This conclusion must be taken with care since the above-mentioned rejection levels are not stringent enough to assure the result with total confidence; ii) the power-law decline in the near-IR should be steeper than the one given by $\alpha_{\rm K^{\prime}}=0.45$; iii) for the observations carried out on Dec 15.44-15.51, there is a change in the slope of the measured energy distribution at some wavelength between the J and K$^{\prime}$ bands (possibly around H). | 98 | 4 | astro-ph9804331_arXiv.txt |
9804 | astro-ph9804041_arXiv.txt | The measurements of CMB anisotropy have opened up a window for probing the global topology of the universe on length scales comparable to and beyond the Hubble radius. For compact topologies, the two main effects on the CMB are: (1) the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the manifold and (2) an infrared cutoff in the power spectrum of perturbations imposed by the finite spatial extent. We present a completely general scheme using the {\em regularized method of images} for calculating CMB anisotropy in models with nontrivial topology, and apply it to the computationally challenging compact hyperbolic topologies. This new technique eliminates the need for the difficult task of spatial eigenmode decomposition on these spaces. We estimate a Bayesian probability for a selection of models by confronting the theoretical pixel-pixel temperature correlation function with the {\sc cobe--dmr} data. Our results demonstrate that strong constraints on compactness arise: if the universe is small compared to the `horizon' size, correlations appear in the maps that are irreconcilable with the observations. If the universe is of comparable size, the likelihood function is very dependent upon orientation of the manifold {\it wrt} the sky. While most orientations may be strongly ruled out, it sometimes happens that for a specific orientation the predicted correlation patterns are preferred over the conventional infinite models. | The remarkable degree of isotropy of the cosmic microwave background (CMB) points to homogeneous and isotropic Freidmann-Robertson-Walker (FRW) models for the universe. The underlying Einstein's equations of gravitation are purely local, completely unaffected by the global topological structure of space-time. In fact, in the absence of spatially inhomogeneous perturbations, a FRW model predicts an isotropic CMB regardless of the global topology. The observed large scale structure in the universe implies spatially inhomogeneous primordial perturbations exist which gave rise to the observed anisotropy of the CMB. The global topology of the universe does affect the local observable properties of the CMB anisotropy. In compact universe models, the finite spatial size usually precludes the existence of primordial fluctuations with wavelengths above a characteristic scale related to the size of the universe. As a result, the power in the CMB anisotropy is suppressed on large angular scales. Another consequence is the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the manifold. One can search for such patterns statistically in the COBE maps, and to the extent that they are not there, one can constrain the size of the universe and its topology. Much recent astrophysical data suggest the cosmological matter density parameter, $\Omega_0$, is subcritical~\cite{opencase}. If a (possibly varying) cosmological constant is absent or insufficient to bring the total density to the critical value, this would imply a hyperbolic spatial geometry for the universe (commonly referred to as the `open' universe in the cosmological literature). The topologically trivial (simply connected) hyperbolic 3-space, $\hm$, is non-compact and has infinite size. There are numerous theoretical motivations, however, to favor a spatially compact universe~\cite{motive}. To reconcile a compact universe with a flat or hyperbolic geometry, consideration of spaces with non trivial topology (non simply connected spaces) is required. A compact cosmological model can be constructed by identifying points on the standard infinite flat or hyperbolic FRW spaces by the action of a suitable discrete subgroup, $\Gamma$, of the full isometry group, $G$, of the FRW space. The infinite FRW spatial hypersurface is the {\em universal cover}, tiled by copies of the compact space (most appropriately represented as the {\em Dirichlet domain} with the observer at its {\em basepoint}). Any point ${{\bf x}}$ of the compact space has an image ${{\bf x}}_i = \gamma_i {{\bf x}}$ in each copy of the Dirichlet domain on the universal cover, where $\gamma_i \in \Gamma$. The hyperbolic manifold, $\hm$, can be viewed as a hyperbolic section embedded in four dimensional flat Lorentzian space. The isometry group of $\hm$ is the group of rotations in the four space -- the proper Lorentz group, $SO(3,1)$. A compact hyperbolic (CH) manifold is then completely described by a discrete subgroup, $\Gamma$, of the proper Lorentz group, $SO(3,1)$. The Geometry Centre at the University of Minnesota has a large census of CH manifolds and public domain software SnapPea~\cite{Minn}. We have adapted this software to tile $\hm$ under a given topology using a set of generators of $\Gamma$. The tiling routine uses the generator product method and ensures that all distinct tiles within a specified tiling radius are obtained. A CH manifold, ${\cal M}$, is characterized by a dimensionless number, ${\cal V}_{\!\cal M}\equiv V_{\!\cal M}/d_c^3$, where $V_{\!\cal M}$ is the volume of the space and $d_c$ is the curvature radius \cite{Thur7984}. There are a countably infinite number of CH manifolds with no upper bound on ${\cal V}_{\!\cal M}$. The smallest CH manifold discovered so far has ${\cal V}_{\!\cal M} =0.94$~\cite{smallestCH}. ~\footnote{ The volume of CH manifolds is bounded from below and the present theoretical lower bound stands at ${\cal V}_{\!\cal M} \ge 0.167$~\cite{gab_mey96}. There exist sharper lower bounds within subclasses of CH manifolds under restrictions on topological invariants~\cite{minvol}. It has been conjectured that the smallest known manifold is in fact the smallest possible~\cite{smallestCH}.} The Minnesota census lists several thousands of these manifolds with ${\cal V}_{\!\cal M}$ up to $\sim 7$. In the cosmological context, the physical size of the curvature radius $d_c$ is determined by the density parameter and the Hubble constant $H_0$: $d_c=(c/H_0)/\sqrt{1-\Omega_0}$. The physical volume of the CH manifold with a given topology, \ie a fixed value of $V_{\!\cal M}/d_c^3$, is smaller for smaller values of $\Omega_0$. Two quantities which characterize linear dimensions of the Dirichlet domain are $R_>$ and $R_<$, the radii of circumscribing and inscribing spheres, respectively. In the standard picture, the CMB that we observe is a Planckian distribution of relic photons which decoupled from matter at a redshift $\approx 1100$. These photons have freely propagated over a distance $ R_{\sc ls} \approx 2 d_c\, {\rm arctanh} \sqrt{1-\Omega_0} $, comparable to the ``horizon'' size. For the adiabatic fluctuations we consider here, the dominant contribution to the anisotropy in the CMB temperature measured with wide-angle beams ($\theta_{\sc fwhm} \gta 2^\circ \Omega_0^{1/2}$) comes from the cosmological metric perturbations through the Sachs-Wolfe effect. The adiabatic cosmological metric perturbations can be expressed in terms of a scalar gravitational potential, $\Phi({\bf x},\tau)$. The dynamical equation for the gravitational potential allows for separation of the spatial and temporal dependence,\footnote{At the scales appropriate to CMB anisotropies, damping effects on $\Phi$ can be neglected.} $\Phi({\bf x},\tau) = F(\tau) \Phi({\bf x},\tauls)$, where $F(\tau)$ encodes the time dependence of the metric perturbations and $\Phi({\bf x},\tauls)$ is the field configuration on the three-hypersurface of constant time $\tau=\tauls$ when the last scattering of CMB photons took place. We shall study open, $\Omega_0<1$, models with zero cosmological constant where in the matter dominated phase,~\cite{mukh92} \begin{equation} F(\tau) = \frac{5(\sinh^2\tau -3 \tau\sinh\tau +4\cosh\tau -4)} {(\cosh\tau -1)^3}\,. \lbl{Feta} \end{equation} Here and further on we use dimensionless conformal time $\tau$ expressed in units of the curvature radius. A non-zero cosmological constant can be trivially incorporated in our analysis by using the appropriate solution for $F(\tau)$. We write the Sachs-Wolfe formula for the CMB temperature fluctuation, $\Delta T(\hat q)$, in a direction $\hat q$, in the form \begin{equation} \fl \dT(\hat q) = \frac{1}{3} \Phi(\hat q\chiH,\tauls) + 2 \int_{0}^\chiH d\chi f(\chi) \Phi(\hat q\chi,\tauls)\,, \quad f(\chi)=\frac{d}{d\tau} F(\tau)\bigg|_{\tau=\chiH-\chi}\,, \lbl{dTSW} \end{equation} where $\chi$ is the affine parameter along the photon path from $\chi=0$ at the observer position to $\chiH=R_{\sc ls}/d_c$. The first term is called the {\em surface} or ``naive'' Sachs-Wolfe effect (NSW). The second term, which is nonzero only if $\Phi$ varies with time between $\tauls$ and now, is the {\em integrated} Sachs-Wolfe effect (ISW). The angular correlation between the CMB temperature fluctuations in two directions in the sky is then given by \begin{eqnarray} \fl C(\hat q,\hat q^\prime)\equiv \left\langle\dT(\hat q)\dT(\hat q^\prime)\right\rangle = \frac{1}{9} \langle\Phi(\hat q\chiH,\tauls)\Phi(\hat q^\prime\chiH,\tauls)\rangle \nonumber\\ \lo +\frac{2}{3}\int_{0}^\chiH d\chi~f(\chi)~\left[ \langle\Phi(\hat q\chi,\tauls)\Phi(\hat q^\prime\chiH,\tauls)\rangle +\langle\Phi(\hat q^\prime\chi,\tauls)\Phi(\hat q\chiH,\tauls)\rangle\right] \nonumber\\ \lo +4\int_{0}^\chiH d\chix{1} ~f(\chix{1})~ \int_{0}^\chiH d\chix{2}~f(\chix{2})~ \langle\Phi(\hat q\chix{1},\tauls)\Phi(\hat q^\prime\chix{2},\tauls)\rangle\,. \lbl{cthetaSW} \end{eqnarray} The main point to be noted is that $C(\hat q,\hat q^\prime)$ depends on the spatial two point correlation function, $\xi_\Phi \equiv \langle\Phi({\bf x},\tauls)\Phi({\bf x^\prime},\tauls)\rangle $ of $\Phi$ on the three-hypersurface of last scattering. This is due to the fact that the equation of motion for $\Phi$ allows a separation of spatial and temporal dependence. Although in this work we restrict our attention to the Sachs-Wolfe effect which dominates when the beam size is large, we should point out that other effects which contribute to the CMB anisotropy at finer resolution can also be approximated in terms of spatial correlation of quantities defined on the hypersurface of last scattering~\cite{us_inprep}. | 98 | 4 | astro-ph9804041_arXiv.txt |
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9804 | astro-ph9804088_arXiv.txt | We have imaged the inner square arcminute of the well known lensing and cooling flow cluster A2390 ($z = 0.23$) down to a sensitivity of 65 and 130 $\mu$Jy at 6.75 and 15 $\mu$m , respectively. We report the first evidence of an active star-forming region in a cooling flow (at those wavelengths) and strong emission in the mid-IR from lensed galaxies located at $z=0.9$. | The cluster of galaxies A2390 (z=0.228) possesses remarkable properties which makes its study particularly attractive: presence of a ``straight'' giant gravitational arc (z=0.913), numerous arclets, an elongated galaxy distribution (\cite{Mellier} ; \cite{Pello}) and a large velocity dispersion (1093 km s$^{-1}$, Carlberg et al., 1996) as well as a high X-ray luminosity ($\sim1.5 ~10^{45}$ erg s$^{-1}$ in the [0.1--2.4] keV band). A deep HRI ROSAT pointing revealed an elongated X-ray morphology, the existence of a secondary maximum responsible for the observed gravitational shear in the optical and a strong cooling flow of $\sim800$ M$_{\odot}$yr$^{-1}$ (Pierre et al., 1996). All this indicates that A2390, and its underlying gravitational potential, is especially relevant for our understanding of massive cluster formation, which is, in a hierarchical scenario, closely related to the history of galaxy/star formation. This has motivated deep ISOCAM observations of the cluster core during the guaranteed time programme DEEPXSRC. We present here the observations and results of the cD galaxy and the $z=0.9$ lensed system. Throughout the paper we assume $H_{o} = 50$ km s$^{-1}$Mpc$^{-1}$ and $q_{o}$ = 0.5. | In this letter, we restrict the discussion to the 4 sources found in the maximum sensitivity area of ISOCAM rasters and seen both at 6.7 $\mu$m and 15 $\mu$m, i.e. \#1--4. Their photometric properties are summarized in Table \ref{sources}. The visible and near-IR spectral energy distributions (SED) computed for these 4 sources are given in Table \ref{SED}, and compared to those obtained for a typical cluster galaxy. Details on these photometric data can be found in Pell\'o et al. (1998) and the references therein. In addition, about 20 objects are identified at 6.7 $\mu$m and 10 at 15 $\mu$m, which will be discussed in a forthcoming paper. All sources are point-like for ISOCAM, except the cD galaxy which extends over two times the PSF FWHM at 6.7 $\mu$m (i.e $\sim 20$ kpc), and then allow us to exclude a pure AGN emission. \subsection{The cD galaxy and its cooling flow} The cD galaxy is detected both at 6.7 $\mu$m and 15 $\mu$m, with a flux of $300^{+50}_{-40} \mu$Jy and $500^{+80}_{-70} \mu$Jy respectively. VLA observations (\cite{Arnouts}) show a point-like source with decreasing radio fluxes of 140, 16 and 5.5 mJy at 6, 2 and 1.3 cm, respectively. Assuming a power law spectrum, the mid-IR flux would be some $10^5$ fainter than observed, which excludes a jet-like synchrotron contribution to the observed mid-IR emission. In galaxies where the mid-IR emission is dominated by an old stellar population, the ratio 6.7 $\mu$m/15 $\mu$m is $> 1$. An excess of 15 $\mu$m emission in field galaxies indicates the presence of dust heated by UV photons from star-forming regions. Compared to the 6.7 $\mu$m/15 $\mu$m ratios observed in other nearby early-type galaxies (\cite{Suzanne}) or in distant cD galaxies in clusters (\cite{next}), the ratio of $\sim 0.6$ found here for the cD is exceptional. This ratio is compatible with the colors of the disk component of Centaurus A (\cite{Felix}), a nearby giant early-type galaxy exhibiting active star-forming regions in dust lanes, due to a merge with a spiral galaxy. Thus, the cD in A2390 is probably also undergoing active star formation. However, our cD galaxy looks notably different in other wavelengths. Cen A shows a jet plus an extended radio emission but our cD does not. In addition, a B image reveals the existence of a filament extended along the main axis of the cD, while V and I HST images show the presence, within the filament, of very blue globules possibly associated with the 15 $\mu$m maximum (Fig. \ref{zoom}). Strong emission lines are present across the long-slit spectrum of the filament (Fig. \ref{Halpha}) with ratios indicative of massive star formation associated with shocks and incompatible with an active nucleus (\cite{Allen}, \cite{Baldwin}). Finally the SED of the filament exhibits a clear excess in the V and B bands with respect to what is expected for a typical elliptical. Assuming that the V flux is mainly produced by forming stars, we derive $M_V = -20.8 \pm 0.2$ and a SFR of $8 \pm 4$ h$_{50}^{-2}$ M$_{\odot}$ yr$^{-1}$ for the optical filament, in agreement with the values obtained from the B flux, corrected for the $[OII]3727$ emission. No absorption has been considered in this calculation, so this estimated SFR has to be taken as a lower limit. According to the results derived from V and B band, and from IR, the optical light is probably coming from the most external regions of the star-forming system, whereas part of the star-formation activity remains shrouded within more dense and dusty clouds, as in the Antennae Galaxies, where absorption is ten times higher when derived from mid-IR than from J, H, K bands, as most massives stars are not visible at optical wavelengths (\cite{antenne}). This implies that we can not exclude a SFR as high as ten times what is derived from the optical for the cD of A2390. Those differences explain why, with a 6.7 $\mu$m/15 $\mu$m ratio of 0.6, we can not just consider the cD of A2390 as an early-type galaxy undergoing simple star-formation in dust lanes as Cen A, but most probably as the place of one or several massive star-bursts which may be located in the central globules (see Fig \ref{zoom}). Indeed, the study of the X-ray image of this cluster demonstrated the presence of one of the strongest cooling flows known ($\sim 800$ M$_{\odot}$ yr$^{-1}$ within a cooling radius of 200 kpc), surrounding the cD galaxy (\cite{marguerite}). Giving the size of the mid-IR emitting region in the cD, $\sim 20$ kpc, we derive a mass flow of 80 M$_{\odot}$ yr$^{-1}$, assuming that $\dot{M} \sim r$ (\cite{fabian}). Note that 20 kpc is also about the size of the optical filament (Fig. \ref{zoom}). However, our present understanding of the relationship between mid-IR dust emission and star formation is still too preliminary to infer quantitative constraints on the IMF or even on the heating processes involved in this complex medium. Finally, the total star-formation rate of $\sim 10$ M$_{\odot}$yr$^{-1}$ deduced from the optical in the filament is clearly a lower limit, and the huge quantity of gas needed could be provided by the cooling flow. However, despite the fact that spiral galaxies are very rare in the core of rich clusters, the hypothesis of a past merge with a late-type galaxy cannot be formally excluded here, which would also provide gas for some $10^{7-8}$ years. \begin{figure} \psfig{file=lemonon.center3.ps,width=8.5cm} \caption []{The V-I color HST image, obtained from F555W and F814W, of the core of the cD galaxy of A2390. White and black crosses indicate 6.7 $\mu$m and 15 $\mu$m maximum emission respectively with the associated uncertainties (see text). Some very blue globules suspected to be active star forming regions could be responsible for the 15 $\mu$m emission which coincides with the most luminous one, and associated with strong cooling flow ($\sim800$ M$_{\odot}$yr$^{-1}$) surrounding this cD galaxy. They are part of a NW--SE filament, aligned with the cD main axis, and the overall cluster X-ray elongation on large scale.} \label{zoom} \end{figure} \begin{figure} \psfig{file=lemonon.halpha.ps,width=9.8cm} \caption []{Mean long-slit spectrum of the filament within the cD galaxy (without reddening correction), from CFHT (\cite{Arnouts}).} \label{Halpha} \end{figure} \subsection{What is new in the arc system of Abell 2390 ?} After the detection of the giant arc at $z=0.724$ in Abell 370 (\cite{Metcalfe}), observation of the complex arc system of Abell 2390 confirms the capability of ISOCAM to point up very distant lensed objects. The giant arc consists of three parts, A at $z=1.033$ (\cite{brenda}), and B--C, at $z=0.913$ (\cite{Pello}). Near IR imaging already distinguished A from B--C, as A was not detected in the K band (\cite{Smail}). HST images revealed that B and C are likely two interacting galaxies. The present ISOCAM images are in full agreement with this picture. Although it was not possible to estimate properly the 6.7 $\mu$m flux because of blending, the 15 $\mu$m/6.7 $\mu$m ratio for the B--C component is well larger than unity which is indicative of the presence of an active star forming region in agreement with the strong [OII] line detected in the optical spectrum (\cite{Pello}). Except for its lower amplification factor, the case of object D is very similar. Its morphology in the HST images is complex with probable signs of interaction and low surface brightness extensions. The existence of starburts in the two interacting galaxies is then not a surprise. The optical and near-IR SEDs of objects B--C and D appear brighter in the near-IR and fainter in the blue bands compared to A. These SEDs can be fitted by different synthetic spectra at $z = 0.913$, using the GISSEL96 code (Bruzual \& Charlot, 1998) to approximately constrain the parameters, and a single stellar population (instantaneous burst), an extinction curve of SMC type (Pr\'evot et al. 1984), and assuming the Scalo IMF (1986). The best fits of the sources B--C and D are obtained with a rest-frame $A_{V} \sim3$ in both cases, a stable result with respect to metallicity changes. The corrected magnitudes for objects B--C and D (lensing and absorption) are very similar ($M_B = -20.8$) , the total mass involved in the burst being $\sim10^{10} M_{\odot}$ in both cases. Despite uncertainties on burst age, a constant star-forming model gives similar results and a mean corrected SFR of $40$ to $50\, M_{\odot} h_{50}^{-2} yr^{-1}$. According to these results, the two lensed sources detected by ISOCAM at $z=0.913$ are strongly reddened star-forming galaxies. In the case of A, there is no need for a reddening correction to fit the SED. Finally, the ISOCAM source \#4 detected in both channels may be associated with a very faint source in the HST image (I = 23.5), with a fuzzy shape. Its 15 $\mu$m/6.7 $\mu$m ratio is very high ($\sim 3$). A photometric redshift of $z=0.4^{+0.2}_{-0.08}$ is proposed for this object by techniques described by Miralles \& Pell\`o (1998). Even if the results are much more uncertain in this case ($75\%$ confidence), the best fit of the SED gives $A_{V} \sim$3.5--4.2 in order to explain the high J and K$^\prime$ emission compared to the optical bands. The corrected SFR is relatively low, $\sim1 M_{\odot} h_{50}^{-2} yr^{-1}$. Taking the photometric redshift into account, the SED of this object, with strong mid-IR emission with respect to its optical counterpart, is probably dominated by the so-called unidentified infrared band emitters, and its colors are similar to those of the post-starburst companion of M51 (Boulade et al., 1996). \subsection{Summary and conclusion} From deep and high-resolution ISOCAM images of the core of Abell 2390 we discovered active star forming regions in the two most distant lensed galaxies ever seen in a cluster by ISO. This allowed us to support the scenario of two interacting galaxies at $z=0.913$ in the ``straight arc'' of A2390, as well as in the other galaxy observed at the same redshift. More interesting, we detect a very faint emission from the cD galaxy at 6.75 $\mu$m, compared to other cluster dominant galaxies at similar redshift (\cite{next}). But, for the first time, the strong 15 $\mu$m/6.75 $\mu$m emission ratio flags the presence of a large amount of warm dust in the cD, probably associated with a very active star forming region located within the envelope of the galaxy. Thus, our observation may further elucidate the fate of part of accumulating gas in the complex cooling flow radio core environment. | 98 | 4 | astro-ph9804088_arXiv.txt |
9804 | astro-ph9804277_arXiv.txt | Using an Eulerian perturbative calculation, we show that the distribution of relative pairwise velocities which arises from gravitational instability of Gaussian density fluctuations has asymmetric (skewed) exponential tails. The negative skewness is induced by the negative mean streaming velocity of pairs (the infall prevails over expansion), while the exponential tails arise because the relative pairwise velocity is a {\it number}, not volume weighted statistic. The derived probability distribution is compared with N-body simulations and shown to provide a reasonable fit. | \label{sec-intro} Redshift surveys present a distorted picture of the world because peculiar motions displace galaxies from their true spatial positions. This phenomenon, which would make redshift surveys useless for intergalactic spaceship navigators, is extremely useful for cosmologists. It can serve as a probe of the dynamics of gravitational clustering and the cosmological mass density parameter, $\Omega$ (\cite{sar77}; \cite{pee80}, hereafter LSS; \cite{kai87}; \cite{ham92}; \cite{pee93}, hereafter PPC; \cite{reg95}). A convenient statistical measure of the distortion effect is the galaxy two-point correlation function in redshift space. Under certain assumptions it can be expressed as a convolution of the true spatial correlation function, $\xi(r)$, with the distribution of the relative line-of-sight velocities of pairs of galaxies, $p(w|r,\theta)$ . Here $r$ and $w$ are respectively, the spatial separation and relative radial velocity of a pair of galaxies, while $\theta$ is the angle between the separation vector ${\bf r}$ and observer's line of sight (cf. LSS; \cite{fis95}, hereafter F95). The purpose of this {\it Letter} is to derive $p(w|r,\theta)$, using weakly nonlinear gravitational instability theory. This distribution was measured from N-body simulations and estimated indirectly from redshift surveys. At $r \simlt$ 1\mpc, where\footnote{We use the standard parametrisation for the Hubble constant, $H = 100 \, h$ \kms Mpc$^{-1}$.} the galaxies are strongly clustered ($\xi \simgt 20$), the observations are consistent with an exponential distribution (\cite{pee76}; \cite{dav83}; \cite{fis94}, hereafter F94; \cite{mar95}; \cite{lsd97}). The fact that $p(w)$ at small separations differs strongly from its initial, Gaussian character, is not surprising: after all, the small-scale velocity field has been `processed' by strongly-nonlinear dynamics in clusters, and exponential distributions were recently derived from the \cite{pre74} theory (\cite{she96}, \cite{dia96}). On larger scales, where the fluctuations have small amplitudes, one na{\"\i}vely expects to see the `unprocessed' initial conditions. However, N-body experiments suggest that $p(w|r,\theta)$ retains its exponential character even at separations $r \simgt$ 10\mpc, where $\xi \simlt 0.1$, despite the fact that the initial density and velocity fields in those experiments were drawn from a Gaussian distribution (\cite{efs88}, hereafter EFWD; \cite{zur94}, hereafter ZQSW; F94). At similar separations, an exponential $p(w|r,\theta)$ has also been inferred from observations (F94; \cite{lov96}). The simulations also show that the radial component of the distribution, $p(w|r,0^{\circ})$ is significantly skewed, in particular at large separations (EFWD; ZQSW; F94). The physical origin of the skewness and exponential shape of $p(w)$ at large separations has until now remained unexplained. We provide the explanation below. | 98 | 4 | astro-ph9804277_arXiv.txt |
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9804 | astro-ph9804310_arXiv.txt | We have examined the hypothesis that the majority of the diffuse EUV flux in the Coma cluster is due to inverse Compton scattering of low energy cosmic ray electrons ($0.16 < \epsilon < 0.31$\,GeV) against the 3$^{\circ}$\,K black-body background. We present data on the two-dimensional spatial distribution of the EUV flux and show that these data provide strong support for a non-thermal origin for the EUV flux. However, we show that this emission cannot be produced by an extrapolation to lower energies of the observed synchrotron radio emitting electrons and an additional component of low energy cosmic ray electrons is required. | Diffuse EUV emission has been detected in five clusters of galaxies: Virgo (Lieu et al., 1996a), Coma (Lieu et al., 1996b), Abell 1795 (Mittaz et al., 1997), and Abell 2199 and Abell 4038 (Bowyer et al., 1997). These clusters were detected with a statistical significance varying from 8 to 50 standard deviations. The diameter of the diffuse emission in these clusters ranges from 20$\arcmin$ to 40$\arcmin$. Some diffuse EUV emission in clusters of galaxies would be expected from the well-studied X-ray cluster emission. However, in all cases examined to date, the EUV emission is far greater than the expected emission from the X-ray-emitting gas. Marginal signatures of this ``soft excess'' are sometimes present in the lowest energy resolution band of the ROSAT PSPC, where they produce less than a 20\% enhancement over the emission expected from the X-ray cluster gas. In contrast, the excesses found with EUVE range from 70\% to 600\% above that expected from the X-ray gas. A variety of instrumental and Galactic interstellar medium absorption effects have been suggested as alternative explanations for the EUVE data, but these have all been found wanting (For a discussion of these issues, see Bowyer et al., 1997). In the original reports of diffuse EUV cluster emission, the data were interpreted in terms of additional thermal gas components in the clusters. In these analyses, the known X-ray emission from the hot cluster gas was first fitted to the EUV data, and the excess EUV emission was computed. This excess was then fitted by additional components of thermal gas. Because of the low energy of the EUV emission, much lower temperature thermal gases ($\sim 10^6$\,K) are required. The concept that additional components of lower temperature gas are present in these clusters has not received wide support, primarily because gas at these temperatures is near the peak of the radiative cooling curve and hence cools rapidly, requiring a substantial energy input to sustain the gas at these temperatures. In addition, it is difficult to understand how different components of gas at grossly different temperatures could retain their separate identity. For example, the Coma cluster has been shown to be formed by merging of distinct subunits, in both X-ray (White et al., 1993), and optical (Colless and Dunn, 1996) studies. This produces variations less than a factor of 3 in temperatures of the X-ray emitting gas (Honda et al. 1996). Deiss and Just (1996) have shown in a general analysis, and with a specific application to the Coma cluster, that turbulent mixing time scales are only a few $10^9$\,years, which argues against the co-existence of major quantities of gas at two vastly different temperatures. However, Cen et al. (1995) have argued that a warm ($\simeq 10^6$\,K) thermal gas is widely distributed throughout the universe, as a direct product of the growth of structure leading, eventually, to clusters of galaxies. In their scenario, the energy required to sustain the warm gas is provided by gravitation. Hwang (1997) has examined the hypothesis that the source of the diffuse EUV flux is inverse Compton scattering by electrons that are a low energy extrapolation of electrons producing the observed synchrotron emission; these electrons are scattered against the 3$^{\circ}$\,K black-body background radiation. The magnetic field he derived for the cluster was 0.2 to 0.4\,$\mu$G which is consistent with the range of estimates for the cluster field. In this work he only considered the constraints imposed by the total EUV flux. En{\ss}lin and Biermann (1998) also considered this mechanism as the source of the EUV flux in the Coma cluster. They assumed that the relativistic energy density of the synchrotron emitting electrons scale radially with the same profile as the X-ray producing gas. This assumption can be questioned since the non-thermal relativistic electrons may well be independent of the thermal X-ray gas. These authors cite the best support for this assumption is given by the data in Figure\ 3 of Deiss et al. (1997) which compares the X-ray and radio radial emission profiles. Unfortunately, the results in this figure are incorrect as has been confirmed by Deiss (private communication). These authors find a magnetic field of 1.2\,$\mu$G is required given their assumptions; this field is also consistent with the range of estimates for the cluster. Sarazin and Lieu (1998) have explored the possibility that EUV radiation in clusters of galaxies could be produced by inverse Compton scattering by a population of very low energy cosmic ray electrons. They showed that the one dimensional EUV spatial profile for the cluster A1795, a radio quiet cluster, was consistent with this hypothesis. A potential problem with this hypothesis as a universal explanation for the EUV emission in clusters of galaxies is that the electrons they proposed have an energy density and pressure which are 1 to 10\% of that of the thermal gas in clusters. If one includes the pressure of cosmic ray ions to the pressure of the electrons proposed by Sarazin and Lieu, using the ratio expected on theoretical grounds (Bell 1978) and measured at Earth orbit (see. e.g., Weber 1983), the total cosmic ray pressure is substantially larger than that of the X-ray emitting gas. In this work we reconsider the hypothesis that the EUV emission in the Coma cluster is the result of inverse Compton emission. We first consider the constrains imposed by the total EUV flux. We review the existing radio data and obtain a different spectral index than that employed by Hwang (1997) and En{\ss}lin and Biermann (1998), which we argue is more appropriate. We then derive results which are generally consistent with the inverse Compton hypothesis. We then consider the two dimensional spatial distribution of the EUV flux; this data provides substantial support for a non-thermal origin for this flux. We find that the spatial distribution of the magnetic field required by En{\ss}lin and Biermann to produce the EUV emission profile is unrealistic. We show that the difference in spatial extent between the EUV and radio halos cannot be explained using an electron distribution which is an extrapolation of the known synchrotron emitting electrons and that an additional population of low energy cosmic ray electrons are required to explain these data. | We have examined the hypothesis that the EUV radiation from the Coma cluster is due to inverse Compton scattering of low energy cosmic ray electrons against the 3$^{\circ}$\,K black-body background radiation. The total integrated EUV emission produced by cosmic ray electrons which are a low energy extrapolation of higher energy electrons, known to be present from their synchrotron emission, gives results which are consistent with the range of estimates of the magnetic field in the cluster. We next consider the two dimensional spatial distribution of the EUV emission. This emission does not follow the distribution of the gravitationally bound X-ray gas, but rather exhibits an asymmetric distribution similar to that exhibited by the radio emission. This suggests a non-thermal origin for the EUV emission rather than a gravitationally constrained thermal gas. We show from a comparison of the size of the EUV halo and the radio halo that the EUV emission cannot be produced by inverse Compton radiation from electrons which are an extrapolation of the distribution which produces the observed radio emission. We develop a model for the EUV emission which is self-consistent and fits the existing data. This model requires an additional component of low energy cosmic rays. Inverse Compton EUV emission is surely present at some level in clusters of galaxies with radio halos. However, it may well be masked by emission from some other more dominant source mechanism. A test of the inverse Compton scattering hypothesis as the source of the EUV flux in the Coma cluster would be provided by a measurement of the size of the radio halo at $\sim 1$\,MHz. Unfortunately, this is a challenging measurement because of instrumental limitations and ionospheric effects. In addition, at these low frequencies self-absorption could affect the surface brightness profile by reducing the flux near the cluster center while increasing the halo size; this effect would have to carefully be considered when interpreting such a measurement. | 98 | 4 | astro-ph9804310_arXiv.txt |
9804 | astro-ph9804126_arXiv.txt | Cosmological implications of clusters of galaxies are discussed with particular attention to their importance in probing the cosmological parameters. More specifically we compute the number counts of clusters of galaxies, \ns relation, in X-ray and submm bands on the basis of the Press--Schechter theory. We pay particular attention to a set of theoretical models which well reproduce the {\it ROSAT} 0.5-2 keV band \ns, and explore possibilities to break the degeneracy among the viable cosmological models. | There are several reasons why clusters of galaxies are regarded as useful probes of cosmology including (i) since dynamical time-scale of clusters is comparable to the age of the universe, they should retain the cosmological initial condition fairly faithfully. (ii) clusters can be observed in various bands including optical, X-ray, radio, mm and submm bands, and in fact recent and future big projects (e.g., SDSS, AXAF, PLANCK) aim to make extensive surveys and detailed imaging/spectroscopic observations of clusters. (iii) to the first order, clusters are well approximated as a system of dark matter, gas and galaxies, and thus theoretically well-defined and relatively well-understood, at least compared with galaxies themselves, and (iv) on average one can observe a higher-z universe with clusters than with galaxies. In particular X-ray observations are well-suited for the study of clusters since the X-ray emissivity is proportional to $n_e^2$ and thus less sensitive to the projection contamination which has been known to be a serious problem in their identifications with the optical data. In fact, various statistics related to the abundances of clusters has been extensively studied to constrain theories of structure formation, including mass function (Bahcall \& Cen 1993; Ueda, Itoh, \& Suto 1993), velocity function (Shimasaku 1993; Ueda, Shimasaku, Suginohara, \& Suto 1994), X-ray Temperature function (hereafter XTF, Henry \& Arnaud 1991; White, Efstathiou, \& Frenk 1993; Kitayama \& Suto 1996 ; Viana \& Liddle 1996; Eke, Cole, \& Frenk 1996; Pen 1996). Previous authors have focused on cosmological implications of cluster XTF mainly because theoretical predictions are relatively easier although the observational data are statistically limited. In addition, the conversion to the number density at high z becomes very sensitive to the adopted cosmological parameters. On the other hand, \ns which we discuss in details below is observationally more robust (Ebeling et al. 1997; Rosati \& Della Ceca 1997) while its theoretical prediction is more model-dependent (Oukbir, Bartlett, \& Blanchard 1996; Kitayama \& Suto 1997; Kitayama, Sasaki \& Suto 1998). In this respect, both statistics are complementary. \begin{figure}[t] \begin{center} \leavevmode \psfig{figure=logns1.cps,height=6.8cm} \psfig{figure=chi2ns2.cps,height=6.8cm} \end{center} \vspace*{-0.5cm} \caption{ {\it Left:} Theoretical predictions for \ns of X-ray clusters in CDM models with different cosmological parameters; (a) $\sigma_8=1.04$ models with different $\Omega_0$, $\lambda_0$ and $h$, (b) $\Omega_0=1$ and $0.45$ models with different $\sigma_8$. Denoted by (COBE) are the models normalized according to the {\it COBE} 4 year data (Bunn \& White 1997). Data points with error bars at $S\simlt 10^{-12}$ \unit are from the {\it ROSAT} Deep Cluster Survey (RDCS, Rosati et al. 1995; Rosati \& Della Ceca 1997), and the error box at $S\simgt 2 \times 10^{-12}$ represents a power-law fitted region from the {\it ROSAT} Brightest Cluster Sample (BCS, Ebeling et al. 1997). For the BCS data at $S= 2 \times 10^{-12}$, $1 \times 10^{-11}$ and $6 \times 10^{-11}$\unit, we also plot the corresponding Poisson errors. {\it Right:} Limits on $\Omega_0$ and $\sigma_8$ in CDM models ($n=1$, $h=0.7$) with (a) $\lambda_0=1-\Omega_0$, and (b) $\lambda_0=0$. Constraints from cluster \ns (solid) and XTF (dotted) are plotted as contours at $1 \sigma$(68\%), $2\sigma$(95\%) and $3\sigma$(99.7\%) confidence levels. Dashed lines indicate the {\it COBE} 4 year results from Bunn \& White (1997). } \label{fig:ns1chi2} \vspace*{-0.5cm} \end{figure} | Let us summarize the conclusions of the present talk. \begin{description} \item{(1)} There exist { several theoretical models} which successfully reproduce the observed \ns relation of galaxy clusters over almost four orders of magnitude in X-ray flux. \item{(2)} The resulting $\sigma_8$ is given by the following empirical fit (95\% confidence limit): \begin{equation} \sigma_8 = (0.54 \pm 0.02 \pm 0.1) \times \Omega_0^{-0.35-0.80\Omega_0+0.55\Omega_0^2} \end{equation} for $\lambda_0=1-\Omega_0$ CDM, and \begin{equation} \sigma_8 = (0.54 \pm 0.02 \pm 0.1) \times \Omega_0^{-0.28-0.91\Omega_0+0.68\Omega_0^2} \end{equation} for $\lambda_0=0$ CDM. \item{(3)} Low-density CDM models ($n=1$) with $(\Omega_0,\lambda_0,h,\sigma_8) = (0.3,0.7,0.7,1)$ and $(0.45, 0, 0.7, 0.8)$ simultaneously account for the cluster \ns, XTF, the {\it COBE} 4 year normalization. \end{description} Maybe the most important point is that many cosmological models are more or less successful in reproducing the structure at redshift $z\sim0$ {\it by construction}. This is because the models have still several degrees of freedom or {\it cosmological parameters} which can be appropriately {\it adjusted} to the observations at $z\sim0$ ($\Omega_0$, $\sigma_8$, $h$, $\lambda_0$, $b(r,z)$). We have shown that, given a complete flux limited cluster sample with redshift and/or temperature information, one can further constrain the cosmological models. In fact, our tentative comparison indicates that our predictions reproduce well the evolutionary features of the XBACs and that the results, although preliminary, seem to favor low density ($\Omega_0 \sim 0.3$) universes. As indicated by this preliminary result, surveys of objects at at high redshifts in several different bands (X-ray, radio and submm) are the most efficient and rewarding to break the degeneracy among the viable cosmological models. | 98 | 4 | astro-ph9804126_arXiv.txt |
9804 | astro-ph9804256_arXiv.txt | Steep soft X-ray (0.1-2 keV) quasars share several unusual properties: narrow Balmer lines, strong FeII emission, large and fast X-ray variability, rather steep 2-10 keV spectrum. These intriguing objects have been suggested to be the analogs of Galactic black hole candidates in the high, soft state. We present here results from ASCA observations for two of these quasars: NAB0205+024 and PG1244+026. Both objects show similar variations (factor of $\sim2$ in 10 ks), despite a factor of about ten difference in the 0.5-10 keV luminosity ($7.3\times 10^{43}$ erg s$^{-1}$ for PG1244+026 and $6.4\times10^{44}$ erg s$^{-1}$ for NAB0205+024, assuming isotropic emission, $H_0$ = 50.0 and $q_0$ = 0.0). The X-ray continuum of the two quasars flattens by 0.5-1 going from the 0.1-2 keV band toward higher energies, strengthening recent results on another half dozen steep soft X-ray AGN. PG1244+026 shows a significant feature in the `1 keV' region, which can be described by either as a broad emission line centered at 0.95 keV (quasar frame) or as edge or line absorption at 1.17 (1.22) keV. The line emission could be due to reflection from an highly ionized accretion disk, in line with the view that steep soft X-ray quasars are emitting close to the Eddington luminosity. Photoelectric edge absorption or resonant line absorption could be produced by gas outflowing at a large velocity (0.3-0.6 c). | The ROSAT PSPC has found a large spread in the energy spectral indices of low-z quasars\footnote{We use ``quasars'' to describe broad line emission objects, regardless of luminosity.} : $0.5<\alpha_{0.1-2keV}<3.5$. In about 10$\%$ of cases $\alpha_{0.1-2keV}\gs 2$ (e.g. Laor et al. 1994, 1997, Walter \& Fink 1993, Fiore et al. 1994). The large spread in $\alpha_{0.1-2keV}$ favoured the discovery of its correlation with other properties. In fact, the steep soft X-ray quasars have then been realized to share a cluster of unusual properties: \begin{itemize} \item narrow Balmer lines \footnote{the permitted lines have FWHM$\ls$2000~km~s$^{-1}$, yet still are clearly broader than the forbidden lines.} ( Laor et al 1994, 1997, Boller et al. 1995); \item strong FeII emission (Laor et al 1994, 1997, Lawrence et al. 1997) \item Rapid, large amplitude variability (factor of 2-50 on timescales from minutes to months, Boller et al., 1995, Brandt et al., 1995, Otani 1995, Boller et al. 1997) \item Somewhat steep hard X-ray spectra (2$>\alpha_{2-10keV} >0.6$, Pounds et al. 1995, Brandt et al., 1997); \end{itemize} Pounds et al. (1995), suggest the latter to be a close physical analogy with the X-ray power-law produced by Comptonization in a hot accretion disk corona in Galactic black hole candidates (BHC) in their `soft-high' state. This is not the only analogy between BHC and steep X-ray spectrum quasars. Laor et al. (1994, 1997) explained the correlation with H$\beta$ FWHM as due to the larger size of a virialized broad emission line region for an AGN in a high $L/L_{Edd}$ state. Ebisawa (1991) found that while the soft component of 6 BHC observed by Ginga is roughly stable on time scales of 1 day or less, the hard component exhibits large variations down to msec time scales. These timescales translates to $10^4$ years and 0.1 day for quasars, if they scale with the mass of the compact object. The soft component of BHC extends up to $\sim 10$ keV in BHC in `soft-high' states, and it is often associated with optically thick emission from an accretion disk. If this is the case, the temperature should scale with the mass of the compact object as $M_{BH}^{-1/4}$, and the above energy translates to 0.1-0.4 keV for quasars. The rapid large amplitude variability shown by a few narrow line Seyfert 1 galaxies (NLSy1) at about 1 keV on timescales of hours to days (Otani 1995, Brandt et al. 1995, Boller et al. 1997) can then be analogous to the above BHC hard component flickering. A steep X-ray spectrum quasar with 10-100 times the luminosity of NLSy1s, should be larger and so should vary no more rapidly than several days. Instead Fiore et al. (1998a) find that steep X-ray spectrum PG quasars commonly vary by a factor 2 in 1~day. Variability seems therefore correlated with X-ray spectral slope and Balmer line width (and therefore possibly with the accretion rate) rather than with the luminosity. Evidence for spectral features in the `1 keV' region in many steep soft X-ray quasars is building up (Turner et al., 1991, Brandt et al., 1994, Otani et al., 1995, Comastri et al., 1995, Leighly et al., 1997, 1998a,b). Instead, `normal' Seyfert 1 galaxies (having broad Balmer lines and flatter soft X-ray spectra) usually have their strongest absorption features at lower energies (in the 0.6-0.9 keV `oxygen' band). An intriguing possibility is that the appearance of these features at different energies also depend on $L/L_{\rm Edd}$. Detailed high energy X-ray spectra of luminous quasars with steep soft X-ray spectra are essential to understand the `narrow-broad line' phenomenon in AGN, in particular whether the peculiar X-ray properties depend on optical luminosity, optical-to-X-ray ratio ($\alpha_{OX}$), or on their Eddington ratio. To this end we selected two bright quasars with $\alpha_{0.1-2keV}>$2.0 (Fiore et al., 1994) at the extreme values of optical luminosity, both with low Galactic $N_H$ (Table 1) of $1.9\times 10^{20}$ cm$^{-2}$ for PG~1244+026, and of $3.0\times 10^{20}$ cm$^{-2}$ for NAB0205+024, Elvis et al., 1989) and observed them with ASCA. We report the results in this paper. | ASCA observations of two steep soft X-ray quasars have shown that: \begin {enumerate} \item The X-ray continuum of the two quasars flattens by $\Delta\alpha=0.5-1$ going toward high energies. Similar results were obtained by authors on some half dozen steep soft X-ray quasars. \item PG1244+026 shows a significant feature in the `1 keV' region. Similar features were again reported in other steep soft X-ray quasars. The data are not good enough to discriminate between a broad emission line centered at 0.95 keV (quasar frame) or an absorption edge at 1.17 keV, or an absorption notch at 1.22 keV. Line emission could be due to reflection from an highly ionized accretion disk, in line with the view that steep soft X-ray quasars are emitting close to the Eddington luminosity. Photoelectric edge absorption or resonant line absorption could be produced by gas outflowing at a large velocity (0.3-0.6 c). In these absorption models significant cold (i.e. oxygen less ionized than OVI) absorption in excess of the Galactic is required. This would imply an increase by a factor 2-3 of the cold column with respect to a previous PSPC observation or a peculiar ionization structure. In neither the emission or absorption cases the SIS resolution is good enough to identify unambiguously the ions responsible for the feature. The high resolution and high throughput of the low energy gratings and spectrometers of AXAF and XMM are clearly needed to shed light on this puzzling case. \item The two quasars show similar variability properties (flux variations up to a factor of 2 in 10 ks) despite a factor of ten difference in the X-ray observed luminosity. This agrees with the Fiore et al (1998a) finding that the variability properties of radio-quiet quasars are correlated with the shape of the X-ray spectrum, the width of the Balmer lines and so possibly with the accretion rate. \end{enumerate} \bigskip F.F acknowledges support from NASA grants NAG 5-2476 and NAG 5-3039, B.J.W. acknowledges support from ASC contract NAS8-39073. | 98 | 4 | astro-ph9804256_arXiv.txt |
9804 | astro-ph9804060_arXiv.txt | Since it has become possible to discovery planets orbiting nearby solar-type stars through very precise Doppler-shift measurements, the role of methods used to analyze such observations has grown significantly. The widely employed model-dependent approach based on the least-squares fit of the Keplerian motion to the radial-velocity variations can be, as we show, unsatisfactory. Thus, in this paper, we propose a new method that may be easily and successfully applied to the Doppler-shift measurements. This method allows us to analyze the data without assuming any specific model and yet to extract all significant features of the observations. This very simple idea, based on the subsequent subtraction of all harmonic components from the data, can be easily implemented. We show that our method can be used to analyze real 16 Cygni B Doppler-shift observations with a surprising but correct result which is substantially different from that based on the least-squares fit of a Keplerian orbit. Namely, using frequency analysis we show that with the current accuracy of this star's observations it is not possible to determine the value of the orbital eccentricity which is claimed to be as high as 0.6. | Recent improvements in the long-term precision of Doppler-shift measurements \cite[]{Marcy:96c::} resulted in several spectacular detections of planetary companions to solar-type stars \cite[for review see the paper of][] {Marcy:97::}. As such discoveries supply indirect evidence of the existence of extra-solar planets, other explanations of observed radial-velocity variations appeared, e.g., stellar pulsations \cite[]{Gray:97::}. The most recent results, however, show that only a planetary hypothesis is acceptable \cite[]{Marcy:98::,Gray:98::}. The usual procedure showing that there exists a planet around s star consits of a direct least-square fit of Keplerian model to the observations. It always gives certain values of orbital parameters and their formal errors. In the case of `good' data this is the best and the quickest way to obtain relable results. However, in the case of spare data with big errors one has to prove that the least-squares method can be used and that the obtained parameter values and their errors are good estimates of the real values. This is a difficult and time consuming task. Without doubt, we have this situation with the Doppler observation of extra-solar planets. We present an analysis of this problem and we show that the eccentricity of the fitted orbit is a very sensitive parameter and, in some cases, its value and error given by the least-squares method are not correct. The aim of this paper is to show how the mentioned problem can be solved in practice. Namely, we propose a method that can be very useful for analyzing radial-velocity variations. It is based on a simple idea involving the subsequent subtraction of periodic components from the data. This approach allows us to analyze the observations without assuming any specific model describing the system behavior (like Keplerian motion or stellar pulsations). After the determination of all significant components of the data, it remains to be decided which process is responsible for what we observe and whether it is possible to choose only one. The plan of this paper is as follows. In Section 2 we analyse observations of 16 Cygni B and we explain why the standard least-squere fit does not give reliable estimates of parameters and their errors. In Section 3 we analyze analytically the Keplerian motion of the system `a star with one planet' in order to learn how its motion modulates the observed star radial-velocities. We investigate mainly the spectral properties of the motion which are essential for our method. In Section 4 we develop a simple numerical technique which can be used to extract all the information we need to compare with the results from Section 3. In Section 5, we perform a numerical test of the method using simulated radial-velocity variations with the orbital parameters of 70 Vir \cite[]{Marcy:96b::}. In section 6, we discuss the application of the method to finding the eccentricity of 16 Cygni B \cite[]{Cochran:97::}. | In this paper we have shown that results obtained from least-squares fits of Keplerian orbits to real Doppler-shift measurements may lead to incorrect interpretations. Specifically, they may give unrealistic or even entirely false values of parameters and their uncertainities. In order to solve these problems we have proposed a new method, frequency analysis, which efficiently provides an independent test of the reliability of determined orbital parameters. This method may deliver a substantial revision of the current values of planets' high eccentricities that are essential for our understanding of the formation and evolution of planetary systems. It might even lead to hints that some of the observed high eccentric planets are in fact planetary systems consisting of more than one planet or at least provide an independent point of view on the same data. These facts, together with the ease of applicability of frequency analysis, make our method worth trying on future observations if not for the data already gathered. | 98 | 4 | astro-ph9804060_arXiv.txt |
9804 | physics9804027_arXiv.txt | It is shown that the plasma, generated during an impact of a meteoroid with an artificial satellite, can produce electromagnetic radiation below the microwave frequency range. This interference is shown to exceed local noise sources and might disturb regular satellite operations. | At the end of 1998, the first modules of the \emph{International Space Station} will be put in orbit around the Earth and this should open new frontiers for life in space. The intensive use of the space makes necessary to know the potential risks. The threat from meteoroids is today well known and several authors have underlined the risks connected with the impact on a spacecraft (for a review, see \cite{REV}). However, the \emph{Olympus} end-of-life anomaly \cite{CASWELL} and the recent work of McDonnell \emph{et al.} \cite{MCDONNELL} put a new light on these issues. The \emph{Olympus} failure is a paradigmatic example: in that case, the impact with a Perseid meteoroid may have generated electrical failures, leading to a chain reaction which culminated with an early end of the mission \cite{CASWELL}. On the other hand, McDonnell \emph{et al.} \cite{MCDONNELL} showed that, if the plasma charge and current production during an impact are considered, meteoroid streams can be very dangerous, even during normal conditions. It should be noted that they considered only damages by direct discharges or current injection in circuits (\emph{e.g.} via the umbilical) \cite{MCDONNELL}. However, there are several other ways by which the plasma could interact with the spacecraft electronics. For example, it is useful to recall the work of Cerroni and Martelli \cite{CERRONI}, in which they showed that thermal forces in impact-produced plasmas could explain the magnetisation observed in the neighbourhood of lunar craters. Even if Cerroni and Martelli studied experimentally hypervelocity impacts of aluminium projectiles on basalt targets, it is possible to extend their work to general hypervelocity impacts. Here, we show that a plasma cloud, generated during a hypervelocity impact of a meteoroid with an artificial satellite, can radiate electromagnetic energy below the microwave frequency range and, therefore, may disturb regular satellite operations. | After the \emph{Olympus} end-of-life anomaly \cite{CASWELL} and the work of McDonnell \emph{et al.} \cite{MCDONNELL}, it seems clear that the meteoroids hazard is not restricted to a mechanical damage. Here it is suggested a new interference path, that is electromagnetic radiation emitted from the impact-produced plasma cloud. Even if the radiated power is not sufficient to destroy anything, it may disturb regular satellite operations. Further investigations should be made on specific satellite, because they require detailed information about onboard electronics, in order to calculate possible couplings and non-linearities. \stars Author wishes to thank Paolo Farinella, of Department of Mathematics of the University of Pisa, for constructive review. \newpage | 98 | 4 | physics9804027_arXiv.txt |
9804 | astro-ph9804195_arXiv.txt | The Hubble Deep Field (HDF) provides accurate multi-band photometry of galaxies to very faint magnitudes, with 10$\sigma$ $AB$ magnitude limit of $m_{814}=27.60$ (\cite{HDF}). The faint limit of the HDF makes it difficult to obtain spectroscopic redshifts for the majority of the galaxies in the field. It is therefore useful to derive estimated redshifts of these galaxies using the available multi-band photometry. Several groups (\cite{Lanze96}, \cite{GwHa96}, \cite{Saw97}) obtained photometric redshifts for the HDF by comparing the observed $UBVI$ fluxes of each object with a set of galaxy spectral templates of different galaxy types redshifted to evenly spaced redshifts. Since spectroscopic redshifts have been measured and published for $\sim$ 100 galaxies in the HDF (\cite{Cohen96}, \cite{Hogg98}; \cite{Stei}; \cite{DEEP}), it is possible to fit analytic expressions for photometric redshifts. In this paper, we explore a simple empirical approach to estimating redshifts of galaxies based on their colors (see \cite{Conno97}, \cite{Brun97}, \cite{Conno95} for an alternative empirical approach); this method has the advantage of being simple, model independent (i.e., it does not depend on the assumption of any particular set of galaxy spectral templates), and easy to use in determining approximate redshifts of $z\la 4$ galaxies. We determine the empirical analytic relations for color redshifts in \S 2. We compare our estimated redshifts for the HDF galaxies with those obtained by the template-fitting method in \S 3. We describe our estimated redshift catalog of HDF galaxies to $z\la 4$ in \S 4 (the Web site address of the catalog is given). We investigate the redshift clustering of HDF galaxies in \S 5, and summarize our results in \S 6. | Using HDF photometric and spectroscopic data, we have determined a set of simple analytic formulae that yield estimated galaxy redshifts to $z\la 4$ in terms of linear combinations of three measured colors, $U-B$, $B-V$, and $V-I$ (Eqs.(\ref{eq:Lz3,1})-(\ref{eq:Hz2,2})). The derived analytic formulae in five color ranges exhibit small dispersions between the estimated and spectroscopic redshifts. For $z\la 2$ galaxies, the redshift dispersion ranges from $\sigma_z=0.034$ to $\sigma_z=0.097$ for different color ranges. For $z\ga 2$ galaxies, we find $\sigma_z=0.14$ and $\sigma_z=0.36$ for two color ranges which typically represent $z\ga 3$ and $z\la 3 $ galaxies respectively. These color-redshift relations apply to about 90% in the sample. The smallest dispersion between the color and the spectroscopic redshifts, $\sigma_z=0.034$, occurs for the $z\la 2$ galaxies satisfying $(U-B)< (B-V) -0.1$; 28 galaxies with measured redshifts are used in deriving the relation for the estimated redshift, with only 4 free parameters (the coefficients in Eq.(\ref{eq:Lz3,1})). There are 230 HDF galaxies with $I< 27$ and measured UBVI magnitudes that belong to this color range; we investigate the large-scale redshift distribution of these galaxies and find evidence for peaks in the redshift distribution that suggest large-scale clustering to at least $z\sim 1$. These results are consistent with those of Cohen et. al. (1996) using observed spectroscopic redshifts of a smaller number of galaxies. We have applied our color redshift formulae to the entire HDF photometric catalog and find that the derived redshifts are consistent with those obtained from spectral template-fitting techniques. The analytic relations, by design, yield lower dispersion than the template-fitting method. The color-redshift relations have the advantage of being simple, model independent, and easy to use. They can be further improved with additional data. These analytic color-redshift estimators are useful in providing empirical estimates of galaxy redshifts to $z\la 4$ using multiband photometry. Our Estimated Redshift Catalog of HDF Galaxies, based on our color redshift formulae for all 848 HDF galaxies with $I<27$ and measured $UBVI$ fluxes, is available by anonymous ftp in the elt/:HDF subdirectory of astro.princeton.edu. Note that our color-redshift relations (Eqs.(\ref{eq:Lz3,1})-(\ref{eq:Hz2,2})) are derived using AB magnitudes and for the HDF filters. For application to other photometric catalogs, the appropriate spectroscopic training set should be used; when such a training set is not available, Eqs.(\ref{eq:Lz3,1})-(\ref{eq:Hz2,2}) may provide useful estimates after appropriate photometric transformation has been performed between the different filter systems. Also note that these color-redshift relations should not be applied to galaxies which lie close to the boundaries of the color ranges. Finally, we note that our color-redshift relations are limited by the absence of measured spectroscopic redshifts for galaxies in the range of $1.4 \le z \le 2.2$ (see Fig.1 and Fig.2). It is very important to obtain spectroscopic redshifts in this range, because it will not only enable better calibration of photometric redshifts, it will also help us understand the nature of galaxies in the intermediate redshift range. | 98 | 4 | astro-ph9804195_arXiv.txt |
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9804 | astro-ph9804148_arXiv.txt | Catalysis of nucleon decay in white dwarfs is used to constrain the abundance of magnetic monopoles arising from Grand Unified Theories. Recent discoveries of the dimmest white dwarf ever observed, WD 1136-286 with $L = 10^{-4.94} L_{\odot}$, place limits on the monopole flux. An abundance of monopoles greater than the new bound would heat this star to a luminosity higher than what is observed. The new bound is $(F/$cm $^{-2}$ s$^{-1}$ sr$^{-1}$) $(\sigma \upsilon/10^{-28} {\rm cm}^2) < 1.3 \times 10^{-20} (\upsilon/ 10^{-3}c)^2$, where $\upsilon $ is the monopole velocity. The limit is improved by including the monopoles captured by the main-sequence progenitor of the white dwarf: $(F/$cm $^{-2}$ s$^{-1}$ sr$^{-1}$ ) $(\sigma \upsilon /10^{-28} {\rm cm}^2) < 3.5(26) \times 10^{-21}$ for $10^{17}$ ($10^{16}$) GeV monopoles. We also note that the dependence on monopole mass of flux bounds due to catalysis in neutron stars with main sequence accretion has previously been calculated incorrectly (previously the bound has been stated as $F (\sigma \upsilon/10^{-28} {\rm cm}^2) < 10^{-28} $ cm $^{-2}$ s$^{-1}$ sr$^{-1}$). We show that the correct bounds are somewhat weaker for monopole mass other than $10^{17}$ GeV. | The question of whether or not magnetic monopoles exist has intrigued theorists and experimentalists for a long time \cite{[1]}. In 1974, t'Hooft \cite{[2]} and Polyakov \cite{[3]} independently showed that magnetic monopoles always appear as stable topological entities in any Grand Unified theory (GUT) that breaks down to electromagnetism. Hence, if Grand Unified theories are shown to be correct, monopoles of mass in the range $10^{15}$ - $10^{19}$ GeV should exist. Rubakov \cite{[4]} and Callan \cite{[5]} calculated that these monopoles catalyze nucleon decay with a cross section characteristic of strong interactions, $\sigma \upsilon \approx 10^{-28}$ cm$^2$. The abundance of these monopoles is an open question. The Kibble mechanism predicts roughly one monopole per horizon volume at the time of the Grand Unified phase transition. However, this estimate provides a severe overabundance of the number of monopoles: monopoles overclose the Universe by many orders of magnitude. Instead an inflationary epoch \cite{[6]} may reduce their density in the Universe. Then the present abundance is difficult to estimate. A clue for experimentalists about what monopole flux to expect can be provided by astrophysics. The Parker bound \cite{[7]} on the flux of monopoles was obtained by requiring survival of $\mu$G magnetic fields observed in our Galaxy and gave $F \leq 10^{-16}$ cm$^{-2}$ sr$^{-1}$ sec$^{-1}$. Subsequent improvements on this work include consideration of the monopole velocities \cite{[8]} due to acceleration by the galactic magnetic field. Another improvement is the extended Parker bound, which required survival of a smaller seed magnetic field in the early period of the Galaxy \cite{[9]}: $F \leq 1.2 \times 10^{-16} (\frac {m}{10^{17}GeV})$ cm$^{-2}$ s$^{-1}$ sr$^{-1}$. Another class of methods for determination of the monopole flux is based on the hypothesis that GUT monopoles give rise to the catalysis of nucleon decay. The basic idea is that monopoles traveling through the Galaxy lose enough energy to be captured in an object (e.q. white dwarfs, neutron stars, etc.) where they subsequently catalyze nucleon decay. The energy produced by the nucleon decay heats up the object and results in a flux of photons from the surface of the object. One can then compare this predicted luminosity with what is actually observed. One must ensure that the monopoles would not make the object brighter than what is seen. The coolest star (or other object) seen provides the tightest limit on the monopole flux. If there were more monopoles than allowed by the bound, then the dimmest star observed could not exist. Several authors have carried out this kind of analysis in neutron stars \cite{[10]}, nearby pulsar and white dwarfs. The strongest bound was obtained from consideration of the catalysis process in PSR 1929+10, an old pulsar \cite{[11]}. From this pulsar, the bound on the product of monopole flux times cross section for catalysis is $(F/$cm$^{-2}$sr$^{-1}$sec$^{-1})(\sigma\upsilon / 10^{-28}$ cm$^{2}) \leq 7 \times 10^{-22}$. If one includes the monopoles captured by the main sequence progenitor of the white dwarf, this bound becomes even tighter \cite{[12]}, \hfill\break $(F/$cm$^{-2}$ sr$^{-1}$ sec$^{-1})(\sigma\upsilon / 10^{-28}$ cm$^2$ $) \leq 10^{-28}$. The consideration of monopole dynamics inside superconducting neutron-star cores leads to a bound $5 \times 10^{-24} \tau_{10}^{-2}$ cm$^{-2}$ sr$^{-1}$ s$^{-1}$ \cite{[13]}, where $\tau_{10}$ is the age (in $10^{10}$ years) of the pulsar's present magnetic field. As neutron stars are the densest astrophysical objects observed, they give rise to the tightest catalysis bounds. However, there is a certain amount of uncertainty due to the fact that the interiors of neutron stars are not well understood. For example, neutron stars can have very large magnetic fields $\sim 10^{12}$G of unknown topology, and the motion of magnetic monopoles inside the neutron star would undoubtedly be affected by these magnetic fields. In addition neutron star interiors may contain pion condensates, again with uncertain effects on the monopoles. Because of the uncertainties with neutron star interiors, we turn to the next densest astrophysical objects in the Universe, white dwarfs. These stellar remnants are far better understood. The flux limits obtained from consideration of the catalysis process in white dwarfs are therefore important. Previously Freese \cite{[14]} considered monopole catalyzed nucleon decay in white dwarfs. By comparing with the lowest luminosity white dwarf that had been seen at that time, she obtained a limit \begin{equation} (F/{\rm cm}^{-2}{\rm s}^{-1}{\rm sr}^{-1}) (\sigma\upsilon /10^{-28}{\rm cm}^2) \leq 2\times 10^{-18}. \end{equation} The present work is motivated by new observational data of cool white dwarfs \cite{[15]}. In particular, Bergeron, Ruiz, and Leggett found a white dwarf 1136-286 (ESO 439-26) with luminosity $10^{-4.94}L_{\odot}$; this is the dimmest white dwarf observed to date. We use the measured luminosities of old white dwarfs to constrain the radiation due to monopole-catalyzed nucleon decay and thus to obtain an upper limit to the average flux of monopoles in the Galaxy. Since a white dwarf with luminosity $10^{-4.94}L_{\odot}$ is observed today, we know that the monopole-induced contribution to the white dwarf luminosity cannot exceed this value. These new data improve the limit on the monopole abundance due to catalysis in white dwarfs \cite{[14]} by roughly two orders of magnitude. Of course, as dimmer white dwarfs are found, the bound will continue to get more restrictive. A monopole flux saturating this bound would keep the white dwarfs at luminosities at least this great and would lead to the prediction that no cooler white dwarfs will be found. As we will discuss, if it were indeed true that monopoles are keeping dwarfs hot, one would expect a different dependence of white dwarf luminosity on mass than expected in the standard model. We shall explicitly indicate the dependence of our results on various parameters. We will parametrize the properties of the white dwarf in terms of typical values from observations: for the mass, $M = M_{0.6} 0.6 M_{\odot}$, for the radius $R = R_{9} 9\times 10^8$cm, and for the average density $\bar \rho = 4\times 10^5$ g cm$^{-3} M_{0.6} R_{9}^{-3}$. The central density is about an order of magnitude higher, $\rho_{c} = 3 \times 10^6$ g cm$^{-3} M_{0.6} R_{9}^{-3}$. Rubakov \cite{[4]} estimated the product of cross section for catalysis and relative velocity $\upsilon$ of the monopole and nucleon to be constant: $\sigma \upsilon = \sigma_{0} = 10^{-28} ({\sigma\upsilon}_{-28})$ cm$^2$. (Throughout, we take $\hbar = k_{B} = c = 1$.) For the thermal nucleon velocities expected inside a carbon and oxygen white dwarfs, $\upsilon \approx 10^{-3}c$, suppression effects may reduce the cross section by a factor of $10^{-2} s_{-2}$ \cite{[16]}, and so we include this factor. In white dwarfs made of helium the suppression effects would be less effective $(s_{-2} = 10)$, and all the monopole flux bounds would be an order of magnitude stronger. | Figure 1 shows a plot of several monopole bounds: the Parker bound, the extended Parker bound, neutron star bounds, and the new white dwarf bound with and without main sequence capture. In the plots we have used the Wood cooling curves to be conservative. We have found that consideration of newly observed white dwarf 1136-286 with luminosity $10^{-4.94} L_\odot$ and with new calculations of white dwarf cooling curves leads to a bound on the monopole flux that is two orders of magnitude lower than previous bounds due to catalysis in white dwarfs. The new bound is $F (\sigma \upsilon/10^{-28} {\rm cm}^2) < 1.3 (1.9) \times 10^{-20} (v/10^{-3}c)^2$ cm $^{-2}$ s$^{-1}$ sr$^{-1}$ for the Segretain \cite{[19]} (Wood \cite{[20]}) cooling curves respectively, where $\upsilon $ is the monopole velocity. The limit is improved by including the monopoles captured by the main-sequence progenitor of the white dwarf: $F (\sigma \upsilon /10^{-28} {\rm cm}^2) < 3.5(26) \times 10^{-21}$ cm$^{-2}$ s$^{-1}$ sr$^{-1}$ for $10^{17}$ ($10^{16}$) GeV monopoles with $g=g_D$. Flux bounds for other monopole masses and parameters are given in Table 1. If cooler white dwarfs are discovered, a stricter bound on the monopole flux will result. We also showed that the dependence on monopole mass of flux bounds due to catalysis in neutron stars with main sequence accretion has previously been calculated incorrectly. Previously the bound due to catalysis in PSR 1929+10 with main sequence accretion has been stated as \cite{[11]} $F (\sigma \upsilon/10^{-28} {\rm cm}^2) < 10^{-28} $ cm $^{-2}$ s$^{-1}$ sr$^{-1}$. Instead, as can be seen in Table 1 and Figure 1, the correct bounds are somewhat weaker for monopole mass other than $10^{17}$ GeV. {\bf Figure Caption} Bounds on the monopole flux as a function of monopole mass. The Parker bound \cite{[7]} due to survival of the galactic magnetic field is plotted, as is the extended Parker bound \cite{[9]} due to survival of the magnetic field early in the history of the Galaxy. Mass density limits ($\Omega h^2 <1$) are plotted for a uniform density of monopoles in the universe. The bounds due to catalysis in white dwarf WD1136-286 as discussed in this paper are plotted; the plots assume the cooling curves of Wood \cite{[20]}, and are very similar to those obtained using cooling curves of Segretain ${\it et al.}$ In addition, the bounds from this white dwarf with main sequence accretion (WD/MS) are plotted for $g=g_D$ (solid line) and for $g=2 g_D$ (dotted line). The bounds due to calaysis in neutron star PSR 1929+10 are plotted, as are bounds due to this neutron star with main sequence accretion. Again the solid line is for $g=g_D$ and the dotted line is for $g=2g_D$. Note that the neutron star bounds with main sequence accretion have dependence on the monopole mass. \vfill\eject \begin{center} \begin{tabular}{|c|c|c|c|c|c|c|c|}\hline &&&&&&$WD {\rm w}/MS$&$NS {\rm w}/MS$\\ \cline{7-8} $M_m$ (GeV)&$\beta$&${b_{\rm crit}\over R}$&$g/g_D$ &$N_{MS}/10^{38}F$&$N_{WD}/10^{38}F$&$F(\sigma\upsilon)_{-28}$&$ F(\sigma\upsilon)_{-28}$ \\[5pt] \hline \lower1.5ex\hbox{$10^{15}$}&\lower1.5ex\hbox{$10^{-2}$} &{0.4}&{1}&2.5 &\lower1.5ex\hbox{0.17} &$8.2\times 10^{-20}$& $6.2\times 10^{-27}$ \\ \cline{3-5}\cline{7-8} &&0.56&{2}&4.9 && $4.3\times 10^{-20}$& $3.2\times 10^{-27}$ \\\hline \lower1.5ex\hbox{$10^{16}$}&\lower1.5ex\hbox{$3\times 10^{-3}$}&0.48&1&7.4&\lower1.5ex\hbox{1.8}&$2.4\times 10^{-20}$& $2.1\times 10^{-27}$ \\ \cline{3-5}\cline{7-8} &&0.62&2&12.3 && $1.6\times 10^{-20}$& $1.3\times 10^{-27}$ \\\hline \lower1.5ex\hbox{$10^{17}$}&\lower1.5ex\hbox{$10^{-3}$}&0.54&1&52&\lower1.5ex\hbox{17}& $3.2\times 10^{-21}$& $3.0\times 10^{-28}$ \\ \cline{3-5}\cline{7-8} &&0.68&2&82 &&$2.2 \times 10^{-21}$& $1.9\times 10^{-28}$ \\ \hline \lower1.5ex\hbox{$10^{18}$}&\lower1.5ex\hbox{$10^{-3}$}&---&1&---& \lower1.5ex\hbox{17 }&---&---\\ \cline{3-5}\cline{7-8}&&0.24&2&10 &&$8.1 \times 10^{-21}$& $1.6\times 10^{-27}$ \\\hline \end{tabular} \end{center} \bigskip \begin{center} \parbox{6in}{\small Table 1: Bounds on the flux $F$ of magnetic monopoles in cm$^{-2}$s$^{-1}$sr$^{-1}$. Monopoles captured by white dwarfs (WD) or neutron stars (NS) and their main sequence (MS) progenitors are included. The white dwarf cooling time is taken to be $\tau = 9.63$ Gyr. $M_m$ is the monopole mass in GeV, $\beta$ is the monopole velocity, ${b_{\rm crit} \over R}$ is the ratio of the critical impact parameter for a monopole in units of the radius of the main sequence star, and the monopole charge is $g= 69 e (g/g_D)$ in units of the Dirac charge $g_D$. The number of monopoles captured by the MS progenitor and by the white dwarf are $N_{MS}$ and $N_{WD}$ respectively. The second to last column is the flux bound due to catalysis in WD 1136-286 (with MS monopoles included). The last column is the (corrected) flux bound due to catalysis in neutron star PSR 1929+10 (with MS monopoles included).} \end{center} \vspace{10mm} \begin{center} \begin{tabular}{|c|c|c|c|c|c|c|}\hline $M_m$ (GeV)&$\beta$&${b_{\rm crit}\over R}$&$g/g_D$&$N_{MS}/10^{38}F$&$N_{WD}/10^{38}F$&$F(\sigma\upsilon)_{-28}/ {\rm cm}^{-2} {\rm s}^{-1} {\rm sr}^{-1}$ \\[5pt] \hline \lower1.5ex\hbox{$10^{15}$}&\lower1.5ex\hbox{$10^{-2}$} &{0.4}&{1}&2.5 &\lower1.5ex\hbox{0.11} &$8.4\times 10^{-20}$ \\ \cline{3-5}\cline{7-7} &&0.56&2&4.9 && $4.4\times 10^{-20}$ \\\hline \lower1.5ex\hbox{$10^{16}$}&\lower1.5ex\hbox{$3\times 10^{-3}$}&0.48&1&7.4 &\lower1.5ex\hbox{1.2}&$2.6\times 10^{-20}$ \\ \cline{3-5}\cline{7-7} &&0.62&2&12.3 && $1.6\times 10^{-20}$ \\\hline \lower1.5ex\hbox{$10^{17}$}&\lower1.5ex\hbox{$10^{-3}$}&0.54&1&52 &\lower1.5ex\hbox{11 }& $3.5\times 10^{-21}$ \\ \cline{3-5}\cline{7-7} &&0.68&2&82 &&$2.4 \times 10^{-21}$ \\ \hline \lower1.5ex\hbox{$10^{18}$}&\lower1.5ex\hbox{$10^{-3}$}&---&1&---& \lower1.5ex\hbox{11 }&--- \\ \cline{3-5}\cline{7-7}&&0.24&2&10 &&$1.0 \times 10^{-20}$ \\\hline \end{tabular} \end{center} \bigskip \begin{center} {\small Table 2: Same as table 1 for white dwarfs, but for cooling time $\tau = 6.47$ Gyr.} \end{center} {\bf Acknowledgements:} We thank J. Allyn Smith, D. Graff, G. Laughlin, G. Tarle, and V.D. Ivanov for helpful conversations. We acknowledge support from the DOE at the University of Michigan. | 98 | 4 | astro-ph9804148_arXiv.txt |
9804 | astro-ph9804238_arXiv.txt | s{ We use numerical simulations of ray tracing through N-body simulations to investigate weak lensing by large-scale structure. These are needed for testing the analytic predictions of two-point correlators, to set error estimates on them and to investigate nonlinear gravitational effects in the weak lensing maps. On scales larger than 1 degree gaussian statistics suffice and can be used to estimate the sampling, noise and aliasing errors on the measured power spectrum. For this case we describe a minimum variance inversion procedure from the 2-d to 3-d power spectrum and discuss a sparse sampling strategy which optimizes the signal to noise on the power spectrum. On degree scales and smaller the shear and convergence statistics lie in the nonlinear regime and have a non-gaussian distribution. For this regime ray tracing simulations are useful to provide reliable error estimates and calibration of the measurements. We show how the skewness and kurtosis can in principle be used to probe the mean density in the universe, but are sensitive to sampling errors and require large observed areas. The probability distribution function is likely to be more useful as a tool to investigate nonlinear effects. In particular, it shows striking differences between models with different values of the mean density $\Omega_m$. } | Weak lensing by large-scale structure (LSS) shears the images of distant galaxies. The first calculations of weak lensing by LSS (Blandford et al. 1991; Miralda-Escude 1991; Kaiser 1992), based on the pioneering work of Gunn (1967), showed that lensing would induce coherent ellipticities of order 1$\%$ over regions of order one degree on the sky. Recently several authors have extended this work to probe semi-analytically the possibility of measuring the mass power spectrum and cosmological parameters from the second and third moments of the induced ellipticity or convergence (e.g. Bernardeau et al. 1997; Jain and Seljak 1997; Schneider et al. 1997). The analytical work cited above suggested that nonlinear evolution of the density perturbations that provide the lensing effect can significantly alter the predicted signal. It enhances the power spectrum on scales below one degree and makes the probability distribution function (pdf) of the ellipticity and convergence non-Gaussian. We have carried out numerical simulations of ray tracing through N-body simulation data to compute the fully nonlinear moments and pdf. Details of the method and results are presented in a forthcoming paper; here we summarize the method and present some highlights of the results in Figures 1-5. We also discuss reconstruction of the dark matter power spectrum and error estimation using the gaussian approximation. | 98 | 4 | astro-ph9804238_arXiv.txt |
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9804 | astro-ph9804242_arXiv.txt | The spatial distribution of the Ly$\alpha$ forest is studied using new HST data for the quasar pair Q~1026--0045 A and B at $z_{\rm em}$~=~1.438 and 1.520 respectively. The angular separation is 36~arcsec and corresponds to transverse linear separations between lines of sight of $\sim$300$h^{-1}_{50}$~kpc ($q_{\rm o}$~=~0.5) over the redshift range 0.833~$<$~$z$~$<$~1.438. From the observed numbers of coincident and anti-coincident Ly$\alpha$ absorption lines, we conclude that, at this redshift, the Ly$\alpha$ structures have typical dimensions of $\sim$500$h^{-1}_{50}$~kpc, larger than the mean separation of the two lines of sight. The velocity difference, $\Delta V$, between coincident lines is surprisingly small (4 and 8 pairs with $\Delta V$~$<$~50 and 200~km~s$^{-1}$ respectively). \\ Metal line systems are present at $z_{\rm abs}$~=~1.2651 and 1.2969 in A, $z_{\rm abs}$~=~0.6320, 0.7090, 1.2651 and 1.4844 in B. In addition we tentatively identify a weak Mg~{\sc ii} system at $z_{\rm abs}$~=~0.11 in B. It is remarkable that the $z_{\rm abs}$~=~1.2651 system is common to both lines of sight. The system at $z_{\rm abs}$~=~1.4844 has strong O~{\sc vi} absorption.\\ There is a metal-poor associated system at $z_{\rm abs}$~=~1.4420 along the line of sight to A with complex velocity profile. We detect a strong Ly$\alpha$ absorption along the line of sight to B redshifted by only 300~km~s$^{-1}$ relatively to the associated system. It is tempting to interpret this as the presence of a disk of radius larger than 300$h^{-1}_{50}$~kpc surrounding quasar A. | \label{intr} One way to probe the transverse extension of the gaseous structures giving rise to the Ly$\alpha$ forest seen in the spectrum of quasars is to observe multiple lines of sight to quasars with small angular separations on the sky and search the spectra for absorptions coincident in redshift. This technique originated with a suggestion by Oort (1981) to test the possibility that the Ly$\alpha$ forest clouds originate in large pancake structures. The first discoveries of common and associated absorption using pairs of distinct quasars (with separations $\sim$~1~arcmin) were made by Shaver et al. (1982) and Shaver \& Robertson (1983). These already indicated the possible existence of very large absorber sizes (hundreds of kpc), even for the Ly$\alpha$ clouds. At about the same time Sargent et al. (1982) found no detectable tendency for Ly$\alpha$ lines to correlate in QSO pairs separated by a few arcmin. Spectra of pairs of gravitational lens images revealed common absorptions on smaller scales (Weyman \& Foltz 1983, Foltz et al. 1984). The idea that Ly$\alpha$ clouds might have large sizes remained controversial untill the analysis by Smette et al. (1992), later confirmed by Dinshaw et al. (1994), Bechtold et al. (1994), Crotts et al. (1994), Bechtold \& Yee (1994), Smette et al. (1995), D'Odorico et al. (1998). Recently, Dinshaw et al. (1995) derived a radius of 330$h^{-1}_{50}$ kpc at $z$~$\sim$~0.7 for spherical clouds from observation of Q0107--0232 and Q0107--0235 separated by 86~arcsec. Larger separations have been investigated by Crotts \& Fang (1997) and Williger et al. (1997). Both studies conclude that the clouds should be correlated on scales larger than 500~kpc.\par\noindent Here we present observations of Q1026--005~A ($m_{\rm r}$~=~18.4, $z_{\rm em}$~=~1.438) and B ($m_{\rm r}$~=~18.5, $z_{\rm em}$~=~1.520), two distinct quasars separated on the sky by 36~arcsec or 300~$h^{-1}_{50}$~kpc ($q_{\rm o}$~=~0.5) at $z$~$\sim$~1. | 98 | 4 | astro-ph9804242_arXiv.txt |
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9804 | astro-ph9804304_arXiv.txt | How has the ``fluffy'' core of the Sgr dwarf galaxy survived multiple strong shocks from the tidal force of the Galactic halo and disc since the formation of the core a Hubble time ago? A scenario that Sgr was deflected to its current orbit by the Magellanic Clouds after a rendezvous on the north Galactic pole $2-3$ Gyrs ago is examined. It is shown that the conditions of the collision fix both the sense of circulation of Sgr and the LMC around the Galaxy and the slope of the Galactic rotation curve. The model argues that the two orthogonal polar circles traced by a dozen or so Galactic halo dwarf galaxies and globular clusters (LMC-SMC-Magellanic Stream-Draco-Ursa Minor along $l \approx 270^o$ and M54-Ter 7-Ter 8-Arp 2-NGC 2419-Pal 15 along $l \approx 0^o$) are streams of tidal relics from two ancient galaxies which was captured on two intersecting polar rosette orbits by the Galaxy. Our results favor the interpretation of microlensing towards the LMC being due to source or lens stars in tidal features of the Magellanic Clouds. We discuss direct and indirect observations to test the collision scenario. | The recently discovered dwarf galaxy at about $25$ kpc from the Sun in the direction of the Sagittarius constellation (Ibata, Gilmore \& Irwin 1994) is the closest galaxy known to us. It is traced by two long trailing/leading tails on the sky (together more than $8^o\times 22^o$ in solid angle) with most of its stars still clustered around a low density luminous core (roughly $0.001 L_\odot{\rm pc}^{-3}$ with semi-axes $1:1:3$ kpc). It is puzzling why this fluffy core of the dwarf galaxy has not been fully ``digested'' by the Galaxy, in the sense that stars have not fully dispersed out of the core despite the severe shocks at pericentric passage from the tidal force of the Galactic halo (about $10-100$ times stronger than that experienced by satellites in the outer halo, the Magellanic Clouds and the Fornax dwarf galaxy included) and shocks when crossing the disc of the Galaxy. The best fit to Sgr's morphology, radial velocity (Ibata, Gilmore \& Irwin 1995) and proper motion (Ibata, Wyse, Gilmore \& Suntzeff 1997) yields an orbit with a pericenter-to-pericenter period of about $0.8$ Gyr and a peri and apo-center at about 10 and 50 kpc respectively (Vel\'azquez \& White 1995). Simulations show that if a typical Galactic dwarf galaxy (such as Fornax) were replaced on Sgr's orbit, it would dissolve in no more than two peri-centric passages by the strong peri-centric tidal shock of the Galaxy near 10 kpc (Vel\'azquez \& White 1995; Johnston, Spergel \& Hernquist 1995; Johnston, Hernquist \& Bolte 1996; Edelsohn \& Elmegreen 1997). This apparently contradicts the observation that the dominant stellar population in the core is older than 10 Gyrs (Mateo et al. 1995, Fahlman et al. 1996), implying that Sgr has survived $10-20$ peri-centric tidal shocks of the Galaxy. To circumvent this dilemma we need to abandon either or both of the following hidden assumptions: (i) the light distribution of Sgr traces its mass, and (ii) Sgr has always been on the same low-pericentric orbit in a rigid Galactic potential for the past 5 to 10 Gyrs. Ibata, Wyse, Gilmore \& Suntzeff (1997) postulate a dense dark halo of Sgr surrounding the luminous part to hold the system together; they require Sgr's mass density to be uniform inside about 3 kpc of its core with a value ($\sim 0.03 M_\odot{\rm pc}^{-3}$) several times the mean Galactic halo density inside $10$ kpc ($0.013M_\odot{\rm pc}^{-3}$). An inspection of Sgr's rosette-like orbit in relation to that of the Magellanic Clouds (MCs) offers a completely different line of thought. They are on nearly orthogonal planes intersecting along the poles with their Galactocentric radii overlapping at about 50 kpc. So an encounter at the north or south pole some time in the past or future is quite inevitable. A recent preprint by Ibata \& Lewis (1998), shortly after the completion of the work reported in this {\it Letter}, also remarked on a small chance of an interaction after noticing in their simulations a weak perturbation to Sgr's orbit when they turned on the moving gravitational field of the massive MCs. Unfortunately the effect was in the end neglected on grounds of low probability without thoroughly exploring the parameter space (of satellite velocities and the Galactic potential) and the important consequences of a rare strong interaction. So same as Ibata et al. (1997) they were thus left with no alternative but to conclude a massive dark halo of Sgr to be the only explanation for Sgr's presence on a low-pericentric orbit after a Hubble time. In this {\it Letter} we examine the encounter scenario, as illustrated in Fig. 1, where Sgr has been pulled back from an originally high angular momentum/energy orbit to the present low angular momentum/energy orbit by the massive MCs. A recent encounter would have the advantage to allow Sgr to spend most of its lifetime on a ``safe'' orbit with a pericenter (say, $20$ kpc) too high to be harassed by the sharply declining tidal force of the Galaxy (e.g., Kroupa 1997, Oh, Lin \& Aarseth 1995); in a halo with an $r^{-2}$ density profile the pericentric shock would drop by a factor of $4$ from $10$ kpc to $20$ kpc. \onecolumn \begin{figure} \epsfysize=15cm \centerline{\epsfbox{sgrmc3d.ps}} \caption{ A 3D view and $x-y$, $y-z$ projections of the orbits of the LMC/SMC (thick red/dashed green lines near the $x=0$ plane) and Sgr (thin blue curve near the $y=0$ plane); the Sun is be at $(x,y,z)=(-8,0,0)$ kpc. The three systems are integrated backward for $3$ Gyrs (with ellipses marking steps of $0.5$ Gyr) from the present epoch (with velocity vectors marked by arrows) inside the Galactic potential.\label{sgrmc3d} } \end{figure} \twocolumn Various interesting aspects of this scenario will be discussed at the end. But the aim of this {\it Letter} is to report an independent constraint on the rotation curve of the Galaxy as imposed purely by {\it timing the collision}. The essence is the following. The random chance for the LMC and Sgr to meet each other is obviously low, about 1\% for a 10 kpc closest approach in the past 3 Gyrs for a general set of Galactic potentials and initial conditions of the satellites. So the same argument could be inverted: once we accept the deflection by the MCs as a plausible way out of Sgr's dilemma a stringent set of conditions on the potential of the halo and the proper motions of the satellites must follow. | 98 | 4 | astro-ph9804304_arXiv.txt |
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9804 | astro-ph9804214_arXiv.txt | A new approach to the study of the large-scale stellar cluster distribution in the Galaxy based on two-point correlation techniques is presented. The basic formalism for this method is outlined and its applications are then investigated by the use of a simple model of cluster distribution in the Galaxy. This provides an estimate of the potentials of the two-point correlation function for indicating clustering in the measured star positions, which can be related to the presence of star clusters in the observed field. This technique is then applied to several areas of the Two Micron Galactic Survey catalogue, from which information is obtained on the distribution of clusters according to position in the Galaxy, as well as about age, density of stars, etc. | Open star clusters provide valuable information on the evolution of the Galaxy. In this paper two-point correlation techniques are used to analyse the distribution of open clusters in order to gain an insight into the structure and evolution of the Galaxy. Open-cluster distributions have been widely studied at optical wavelengths as a means of studying Galactic structure and evolution (see, for example, Lyng\aa \ 1987b; Janes \& Phelps 1994). The Lyng\aa\ catalogue of open clusters (Lyng\aa \ 1987a) lists about 1200 clusters, which represent nearly all the open clusters accessible in the visible. Knowledge of the positions and ages of these clusters (a method of age determination for clusters is given by Carraro \& Chiosi 1994) enables the scale length and scale height of the disc to be derived for both young and old clusters (for a review of old open clusters, see Friel 1995) and theories to be developed on their formation and destruction throughout the history of the Galaxy. The limitations on these studies are imposed by the maximum distance at which open clusters can be detected. Most of the cataloged open clusters are in the solar neighbourhood, and very few have distances greater than 3 kpc (Payne-Gaposchkin 1979). Hence, information is obtained only for a small region of the Milky Way. The problem is caused by interstellar extinction. An excellent tool for studying star clusters and star formation regions is to observe in the infrared (Wynn-Williams 1977), where the effect of the extinction is far less. To date, however, the infrared has been little used in this field owing the absence of suitable databases. The $K$ band is probably the best region of the spectrum for tracing the stellar distribution of the Galaxy. The $K$ radiation is a mass tracer in spiral galaxies because it follows the old stellar population (Rhoads 1995). Furthermore, the $K$ light is dominated by high-mass stars in star formation regions, i.e. in open clusters, so it is especially appropriate in the search for young clusters, which are rich in massive stars. As explained in more detail by Garz\'on et al. (1993), the Two micron Galactic survey (TMGS) is a $K$-band survey of various regions along the Galactic equator between $-5^o<l<35^o$, $|b|\le 15^\circ$ and $35^\circ<l<180^\circ$, $|b|\le 5^\circ$. The TMGS catalogue has a positional accuracy of about 4 arcsec in right ascension and 7 arcsec in declination. These errors have been estimated after cross-correlating the original TMGS source positions with {\it Guide Star Catalogue} (GSC) counterparts. The larger error in declination comes from the orientation of the array with respect to the survey direction. Due to the dead spaces between detectors, the $K$ limiting magnitude for completeness has to be set conservatively within a range from 9 to 9.5 mag, although the detection limit magnitude of the survey is well in excess of 10 mag. In this paper a new method is presented for automatically determining the level of clustering in catalogues, the TMGS being used as an example. A set of criteria are defined for an automatic search for correlations among stars by means of the two-point correlation function and the two-point angular correlation function. A simple model which assumes a regular distribution of clusters with a constant star density is developed in Section 3. Predictions from this model are then compared with the TMGS and the discrepancies analysed. The causes of clustering are then discussed. The use of the tools described in Section 2 and their application to the TMGS catalogue are dealt with in Section 5, and the results for several regions of our Galaxy are given in Section 6. Finally, a summary of the results is given, and suggestions are made for future developments of the methodology described here. \section {The two-point correlation function and the two-point angular correlation function} Occasionally, when the scale length of a system is changed, certain aspects of the system remain invariable, as in the case for the distribution of matter in space. For example, there are mathematical methods for handling the spatial distribution of atoms in solids, gases and (particularly) liquids. Cosmologists face the same kind of mathematical problem when working with the distribution of galaxies and clusters of galaxies in the context of the large-scale structure of the Universe. They treat the Universe as a fluid whose `particles' are galaxies. Our aim in this paper is to develop the use of similar mathematical methods on an intermediate scale, i.e. in examining the distribution of the stars that make up our Galaxy. Correlation functions describe how points are distributed (e.g. Peebles 1980; Borgani 1995). Suppose a local density of objects $\delta N/{\delta V}=\rho(\vec{r})$ and an average density $\langle\rho\rangle$ (hereafter, the symbol $\langle\cdots\rangle$ indicates a local volume average). Note that $\rho $ corresponds to the exact distribution of objects, i.e. it is a Dirac delta function with zero value where there are no objects, and $\langle\rho\rangle$ is the local average density, i.e. the number of objects per unit volume, and provides no information concerning their distribution. The two-point correlation function (TPCF) is defined as \begin{equation}\xi (|\vec{r}-\vec{r}'|)=\frac{\langle\rho(\vec{r}) \rho(\vec{r}')\rangle} {\langle\rho(\vec{r})\rangle^2} -1 .\label{xi}\end{equation} The function $\xi (r)$ expresses the excess over the random probability of finding objects at separation $r$. ($\xi (r)=0$ means that the probability is totally random; $\xi (r)>0$ that the probability is greater than random, i.e. that there is clustering; and $\xi (r)<0$ that the probability is less than random, i.e. that there is relative avoidance). In the same way, the two-point correlation function can be defined for two dimensions on the surface on to which the distribution is projected (the celestial sphere in the case considered here). This is called the two-point angular correlation function (TPACF) and is defined as \begin{equation}\omega (|\vec{\theta}-\vec{\theta}'|)=\frac{\langle\sigma (\vec{\theta}) \sigma(\vec{\theta}')\rangle} {\langle\sigma(\vec{\theta})\rangle^2} -1 \label{omega},\end{equation} where $\sigma $ is the surface density per unit solid angle. Another mathematical technique for deciding whether a distribution is non-Poissonian is area tessellation, as was used by Bal\'azs (1995) to test the grouping tendency of H$_\alpha$-emission stars in the Orion molecular clouds without giving a quantitative measure of the departure from the Poissonian distribution. See also P\'asztor et al. (1993), P\'asztor \& T\'oth (1995) and references therein for other astronomical applications of spatial statistics. \subsection{Relationship between the TPCF and the TPACF for stars} When applying the above definitions to stars in the Galaxy, the luminosity function and space density have to be taken into account. By generalizing the result of the Limber (1953) equation for constant density, the relationship between the TPCF (which is non-zero for distances less than $\Delta r$) and the TPACF (for small $\theta$) for any density distribution is \[ \omega _{\rm t}(\theta)\approx \frac{1}{\langle\sigma _{\rm t}\rangle^2} \int _0^{\infty }dr\ r^4 \langle\rho \rangle^2(r) \int_{r-\Delta r}^{r+\Delta r}dr' \]\[ \times \int _{M_{\rm min}(r)}^{M_{\rm max}(r)} dM\ \phi(M)\int_{M_{\rm min}(r)}^{M_{\rm max}(r)}dM'\ \phi(M') \]\begin{equation} \times \xi \left(\sqrt{ (r\theta )^2+(r-r')^2};r,M,M'\right) ,\label{omegat2}\end{equation} where $r$ is the distance along the line of sight, $M$ the absolute magnitude, $\phi (M)$ the lu\-mi\-no\-si\-ty function and \begin{equation} \langle\sigma _{\rm t}\rangle=\int _0^{\infty }dr\ r^2\langle\rho \rangle(r)\int _{M_{\rm min}(r)} ^{M_{\rm max}(r)} dM\ \phi(M) .\label{sigmat2}\end{equation} The minimum and maximum values of $M$ for a distance $r$ depend on the minimum and maximum values of the apparent magnitude and the extinction along the line of sight. In this case, it is assumed that the absorption is not patchy, i.e. that it is independent of $\theta $ for small angles. This is not exactly true but it will be show in Section \ref{.extincorr} that the effects are negligible. The subscript `t' stands for `total', a projection of all distances and magnitudes, and $\langle\sigma _{\rm t}\rangle$ is the total two-dimensional density for all distances and magnitudes. In the literature, $\langle\sigma _{\rm t}\rangle$ is also called $A(m_{\rm min},m_{\rm max},l,b)$ and represents the star counts in the magnitude range (Bahcall 1986). This expression enables the TPACF to be found once the three-dimensional distribution of the stars is known and forms the basis of this article, in which we create a model distribution of the stars and compare the results obtained with those observationally in order to investigate the distribution of clustering in the structure of our Galaxy. In general, the TPACF cannot be inverted to give the TPCF due to the multiplicity of possible solutions and to the lack of precise knowledge of certain parameters. However, there are certain cases in which the equation can be inverted and TPCF obtained from TPACF (Fall \& Tremaine 1977). A trivial example where inversion is possible is that of a Poissonian three-dimensional distribution, which implies a Poissonian projected distribution and vice versa, i.e. $\xi =0$, $\omega =0$ on all scales. Another example is when $\langle\rho\rangle(r)$ is a constant independent of $r$. \subsection{ Definition of new variables} In order to simplify the comparison of the level of clustering for different regions of the sky, two new variables will be introduced. $\theta _{\rm max}$ is defined as the first zero of $\omega _{\rm t}(\theta)$. In this article (see for example Fig. \ref{Fig:a052}), $\omega _{\rm t}$ is positive up to a separation $\theta _{\rm max}$. For values greater than $\theta _{\rm max}$ this value is small and oscillates about zero, as there is no correlation among stars separated by large angular distances. Another definition, corresponding to the integration of $\omega _{\rm t}$ up to the limit $\theta =\theta _{\rm max}$ ($\theta >\theta _{\rm max}$ would give a null contribution to the integral), is \begin{equation} C_2\equiv \frac{\int _0 ^{\theta _{\rm max}}d\theta \ \theta \omega _{\rm t}(\theta ) }{\theta _{\rm max}^2} ,\label{C2}\end{equation} which means the excess (when $C _2$ is positive) or deficit (when $C _2$ is negative) of the relative number of objects with respect to a Poissonian distribution in a circle centred on an arbitrary star on the celestial sphere, within the observed solid angle and with angular radius $\theta _{\rm max}$. The relative correlation within the angular scale $\theta _{\rm max}$ is therefore measured. We call $C_2$ the `accumulation parameter' (N.B. there are also other definitions in the literature of the TPACF integral, e.g. Wiedemann \& Atmanspacher 1990). The variable $C_2$ has a clear meaning associated with projected clustering and is also a useful number to measure. Since it sums several values of $\omega $ for different angles, it condenses the information of interest into a single number that can be compared for different samples of stars and give the degree of clustering. This parameter is a mathematical expression of the degree of clustering seen in fields of stars. The idea that we wish to stress here is that all mathematical developments described in this paper are designed to put the intuitive idea of clustering to a reliable test. These calculations are necessary for a quantitative, as opposed to a merely qualitative, description of clustering. Applying the expression (\ref{omegat2}) of $\omega _{\rm t}$ in $C_2$, we get \[ C_2=\frac{1}{\langle\sigma _{\rm t}\rangle^2\theta _{\rm max}^2} \int _0^{\infty }dr\ r^2\langle\rho (r)\rangle^2 \int _{M_{\rm min}(r)}^{M_{\rm max}(r)} dM\ \phi(M) \]\begin{equation} \times \int _{M_{\rm min}(r)}^{M_{\rm max}(r)} dM'\ \phi(M') \int _0^{\theta _{\rm max}r} dy\ y \Xi (y;r,M,M') \label{C2a},\end{equation} where $\Xi $, an integrated TPCF, is \[ \Xi(y;r,M,M') \]\begin{equation}= \int_{r-\Delta r}^{r+\Delta r}dr' \xi\left (\sqrt{y^2+(r-r')^2};r,M,M'\right) .\label{Xi}\end{equation} \subsection{ Further approximations} In order to simplify the above calculations, it will be assumed that the distribution of stars does not depend on their luminosity, i.e. that $\xi (y;r,M,M')=\xi (y;r)$. This is not completely true as there is a small dependence on the distribution of stars in a cluster according their masses, and the luminosities are dependent on the masses. A complete calculation taking the luminosity function into consideration would be of great value. However, the relationship between the TPCF and the luminosity function is uncertain, although the effects of this approximation for the detection of clusters are expected to be small. With this approximation, and from (\ref{sigmat2}) and (\ref{omegat2}), \begin{equation} \langle\sigma _{\rm t}\rangle=\int _0^{\infty}dr\ \langle N^*\rangle(r) \end{equation} and \begin{equation} \omega _{\rm t}(\theta)=\frac{1}{\langle\sigma _{\rm t}\rangle^2} \int _0^{\infty}dr\ \langle N^*\rangle^2(r) \Xi (r\theta ;r) ,\label{omegat*}\end{equation} where $\langle N^*(r)\rangle$ is the number of stars observed per unit solid angle at a distance $r$: \begin{equation} \langle N ^*\rangle(r)=r^2\langle\rho \rangle(r) \int _{M_{\rm min}(r)}^{M_{\rm max}(r)} dM\ \phi(M) .\label{N*}\end{equation} The variable $\omega _{\rm t}$ can also be expressed as \begin{equation} \omega _{\rm t}(\theta)=\overline{\Xi (\overline{r}\theta )} \label{omegaavXi}\end{equation} where the averages $\overline{r}$ and $\overline{\Xi }$ are such that match (\ref{omegat*}). Also, from (\ref{C2a}), \begin{equation} C_2=\frac{1}{\langle\sigma _{\rm t}\rangle^2\theta _{\rm max}^2} \int _0^{\infty }dr\ \frac{\langle N^* (r)\rangle^2} {r^2} \int _0^{\theta _{\rm max}r} dy\ y \Xi (y;r) .\label{C2*}\end{equation} \noindent This last equivalence is a way of averaging the function $\xi $. Hence, high values of $C_2$ indicate that there must be high projected clustering in the direction of the beam. \subsection{Patchiness of extinction} \label{.extincorr} It is clear that extinction can distort the observed counts, the amount of the distortion being a matter of controversy. It is generally accepted that in the optical wavelengths this influence is very severe, particularly in regions near to or in the Galactic plane in the inner Galaxy, where the strong and patchily distributed obscuration makes it difficult to penetrate deep into the Galaxy. The amount of extinction decreases substantially with increasing wavelength. Maihara et al. (1978) quoted a value of 0.17 mag kpc$^{-1}$ as typical for extinction in the Galactic plane in the $K$ band, compared with 1.9 mag kpc$^{-1}$ for the $V$ band (Allen 1973). This has two important consequences. First, the $K$ band is more effective at penetrating the interstellar dust. Secondly, the observed stellar dis\-tri\-bu\-tion more closely resembles the true distribution. For the second argument to be true it is necessary that the obscuration in the $K$ band should not only be smaller in amount than in the $V$ band, but also that its patchiness should be less important. This rather uniform distribution of the interstellar extinction in $K$ can be inferred from the TMGS histograms in several cuts across the Galactic plane. Garz\'on et al. (1993, their Fig. 8) compared the observed stellar distribution in the TMGS and the GSC in the $V$ band. It is noticeable how uniform the $K$ histograms are, particularly when compared with those for the GSC. Except for small portions highly concentrated in the Galactic plane and more marked in the central regions, the shape of the high spatial resolution distribution curves of the TMGS does not exhibit the `noisy' pattern of the GSC plots, which is certainly due to the presence of strong and patchily distributed extinction. Hammersley et al. (1994) showed similar histograms for different areas which also have similar shapes. Moreover, a good fit to a classical ex\-po\-nen\-tial disc can be seen in Fig. 3 of that paper; this would not be the case if the extinction were important and non-uniform. This conclusion can also be reached from the contour maps of the bulge of the Galaxy of Dwek et al (1995), who showed the residuals of the DIRBE data after disc subtraction and extinction correction. Again, the general shape of the maps proves the basic uniformity of extinction distribution in the near infrared. We now estimate these effects. From (\ref{sigmat2}) with the change of variable $r=10^{(5+m_{\rm max}-M_{\rm max})/5}$ and \begin{equation} \Phi (M_{\rm max})=\int _{-\infty} ^{M_{\rm max}} dM\ \phi(M) \approx \int _{M_{\rm min}}^{M_{\rm max}} dM\ \phi(M) ,\end{equation} the local cumulative counts $\sigma _{\rm t}$ follow the expression \[ \sigma _{\rm t}=\langle\sigma _{\rm t}\rangle _{local} =200(\ln\ 10)10^{3m_{\rm max}/5}\omega \int_{-\infty }^\infty dM_{\rm max} \]\begin{equation}\times D\left( 10^{(5+m_{\rm max}-M_{\rm max})/5}\right) 10^{-3M_{\rm max}/5} \Phi (M_{\rm max}), \end{equation} ignoring the variation of extinction with the distance. If we take the density $D$ as constant, then \begin{equation} \sigma _{\rm t}= N(m_{\rm max})\propto 10^{3m_{\rm max}/5} .\end{equation} Taking $D$ as constant is sufficient for estimating the the order of magnitude of the patchiness due to extinction. In any case, the above proportionality is followed in the observed cumulative counts but with a constant value of between 1 and 2 instead of $3/5$ in the exponent. An excess of extinction, $\Delta a(\theta)$, due for example to a cloud at an angular distance $|\vec{\theta } -\vec{\theta _0}|$ with respect to a given point $\theta _0$, will cause a reduction in the apparent flux of a fraction, $f$, of stars (behind the cloud), thereby creating the same effect as a reduction in the maximum apparent magnitudes of these stars by $\Delta a(\theta)$ mag, or, if $\Delta a(\theta)$ is relatively small, a reduction in $m_{\rm max}$ by $f\Delta a(\theta)$ mag for all stars. Hence, \begin{equation} \sigma _{\rm t}(\theta)\sim \sigma _{\rm t} (\theta _0)10^{-3(a(\theta )-a(\theta _0))f/5} \label{sigmada} .\end{equation} If it is assumed that the observed flux fluctuations, $\Delta F$, are due mainly to extinction variations, with the small-fluctuation approximation, then both are related by \begin{equation} \Delta a=-2.5\log \left( 1-\frac{1}{f}\frac{\Delta F}{F} \right)\approx \frac{5\log _{10}e}{2f}\frac{\Delta F}{F} \end{equation} (the factor $f$ appears again here for the same reasons as above). So, from equation (\ref{sigmada}), using the small-fluctuation approximation, \begin{equation} \frac{\sigma _{\rm t}(\theta )} {\sigma _{\rm t} (\theta _0)} \sim \frac{3}{2} \frac{F(\theta)}{F(\theta _0)} \label{omegaext} .\end{equation} This means that the angular correlation of star density is about 3/2 times the angular correlation of the flux. Averaging the DIRBE $K$ flux (Boggess et al. 1992) fluctuations from the maps with $2520''$ resolution over $\mid b\mid \le 3^\circ$ for constant-$l$ strips over the range $-35^\circ < l< 35^\circ $ (where the effects of extinction are most relevant), we get root mean squares of \begin{equation} 0.03< \sigma (\Delta F/F)_{l={\rm const.}, \mid b \mid \le 3^\circ} < 0.23 ,\end{equation} with an average of \begin{equation} \overline{\sigma (\Delta F/F) _{l={\rm const.}, \mid b \mid \le 3^\circ}}= 0.10 .\end{equation} The oscillations of flux fluctuations are not very high in the plane, their maximum being $2.3$ times the average. From equation (\ref{omegaext}), and taking into account that the root mean square is $\sqrt{\omega (0)}$, \begin{equation} \overline{\omega (0) _{l={\rm const.}, \mid b \mid \le 3^\circ}} \sim 0.015 \end{equation} for regions of 2520$''$ in size. In the most unfavourable case, where the extinction is highest (multiplied by a factor of 2.3$^2$ because the maximum root mean square is $2.3$ times greater than the average), $\omega (0) \sim 0.08$. Higher-resolution flux maps are not available in the $K$-band for the whole sky so we cannot derive these numbers for smaller scales, but they are not expected to be much higher since average cloud size is of the order of degrees (rather higher than 42$'$) and the cloud distribution is fairly smooth. A fractal distribution would increase the contamination but this may apply only to very cold gas clouds (Pfenninger \& Combes 1994) which are not the main cause of extinction in the $K$ band. We conclude that extinction in $K$ cannot be responsible for correlations $\omega (0)$ greater than $\sim 0.08$. This is just an estimate, but the order of magnitude should not be very different. As will be shown, the results when applied to the TMGS are above this value (see, for example, Fig. \ref{Fig:a052}), and causes other than patchy extinction must explain this. | A technique is developed for searching for clustering in stellar surveys using correlation functions. The mathematical tools are useful for any field of stars and can be applied to any survey, especially those at carried out at infrared wavelengths, which permit a study of the distribution of stars throughout almost the entire Galaxy. The DENIS (Epchtein 1997) or 2MASS (Skrutskie et al. 1997) surveys will be ideal for this technique as the increased numbers of stars will reduce the errors. It is even possible, with a large number of stars in the survey, to apply the technique for different ranges of apparent magnitude. Studying the clustering of stars at different apparent magnitudes is equivalent to do studying in three dimensions ($l$, $b$ and the average distance $\overline{r}$ which is associated with the treated range of magnitudes). A simple model has been developed. This model could be improved by introducing a density dependence as a function of the distance from the centre of the cluster, perhaps a power-law dependence. In this paper the method has been applied to the TMGS. Is has been shown that a simple model in which old open clusters trace the whole Galaxy with a density of clusters proportional to the density of stars agrees quite well with the data. An exception to the general agreement are specific regions in the plane where the higher-than-expected clustering can be a\-ttri\-bu\-ted to star formation in the spiral arms. A second departure from the simple model is the reduced $C_2$ in the outer disc and in the bulge due to a lack of young clusters. In one of the regions with an excess, at $l=70^\circ $ in the plane, the approximate limits for the cluster density and the density of stars inside the cluster are derived. These are, respectively, $5\times 10^{-7}\ {\rm pc}^{-3} <{\rm \langle}n_{\rm cl}{\rm \rangle}< 2 \times 10^{-6}\ {\rm pc}^{-3}$ and $140\ {\rm pc}^{-3} < \rho _{\rm cl} < 700\ {\rm pc}^{-3}$. There is, however, a lower-than-expected correlation at $l=27^\circ $, $b=0^\circ $. There is believed to be a huge star formation region in this direction and the lack of correlation could be due to the star formation region being far larger than the sample area. As has been pointed out by Friel (1995), the oldest open clusters may be viewed from two perspectives with regard the formation of the Galaxy: a halo collapse or a continuous accretion and infall of material from the halo on to the Galactic disc. Either perspective is possible. The first should justify which were the original star formation regions that were the origin of the present old clusters in the outer disc and how they travelled there from their place of origin. The second perspective needs to test the infall of matter from the halo as well as the existence of star systems in the halo. Further improvements on these cluster searches and better numbers will give us a hint concerning these questions on the origin of the Galaxy. A better determination of $C_2$ in the bulge region will tell us about the age of bulge clusters if these exist. In this article we have observed a relative absence of correlation in the bulge that is somewhat less than the prediction of our simple model, but at best the prediction could say, as in the case of the anticentre, whether the correlation is greater or less than the improved model and enable us to reach further conclusions. \subsection* | 98 | 4 | astro-ph9804214_arXiv.txt |
9804 | astro-ph9804164_arXiv.txt | We analyze ultraviolet ($\sim 1500$~\AA) images of the old open clusters M67, NGC 188, and NGC 6791 obtained with Ultraviolet Imaging Telescope (UIT) during the second flight of the {\em Astro} observatory in March 1995. Twenty stars are detected on the UIT image of M67, including 11 blue stragglers, seven white dwarf candidates, and the yellow giant -- white dwarf binary S1040. The ultraviolet photometry of the blue stragglers F90 (S975) and F131 (S1082) suggests that these stars have hot subluminous companions. We present a semi-empirical integrated ultraviolet spectrum of M67, and show that the blue stragglers dominate the integrated spectrum at wavelengths shorter than 2600~\AA. The number of white dwarfs in M67 is roughly consistent with the number expected from white dwarf cooling models. Eight candidate sdB/sdO stars are detected in NGC 6791, and two are detected in NGC 188. The luminosity range $1.10 < \log$ L/\lsun\ $< 1.27$, derived from the ultraviolet photometry of the six sdB candidates, is consistent with theoretical models of metal-rich hot horizontal branch (HB) stars. The fraction of hot HB stars in both NGC 6791 and NGC 188 is about 30\%, implying that the integrated spectra of both clusters should show a UV turnup at least as strong as that observed in any elliptical galaxy. | As a stellar population ages, the main-sequence turnoff becomes cooler and fainter, until, after an age of about 2 Gyrs, the ultraviolet ($\sim$ 1500~\AA) flux is no longer dominated by stars near the main-sequence, but instead must be due to a more exotic population of hot stars, such as blue stragglers, hot subdwarf (sdB, sdO) stars, hot post-asymptotic giant branch (post-AGB) stars, white dwarfs, or interacting binaries. The blue stragglers and hot subdwarfs are of particular interest since they might have significant roles, respectively, in two current problems in extragalactic astronomy: the age-dating of old galaxies from their rest-frame near-UV ($\sim 2600$~\AA) spectra, and understanding the origin of the observed UV-upturn in elliptical galaxies. In the absence of blue stragglers, the rest-frame near-UV spectrum of an old galaxy should be dominated by the turnoff population, and thus fitting the spectrum to population synthesis models should be a reliable method for age-dating the galaxy, although metallicity effects also need to be properly disentangled (\cite{heap98}, \cite{bruz97}). However, as noted by \markcite{spin97}Spinrad et al. (1997), a significant underestimate of the galaxy age could result if blue stragglers contribute a large fraction of the galaxy near-UV flux, but are not included in the adopted population synthesis model. \markcite{greg90}Greggio \& Renzini (1990) outlined the various candidates for the origin of the observed UV-upturn in elliptical galaxies, which is characterized by a rising UV flux shortward of 1800~\AA\ (\cite{bica96b}, \cite{brown97}). Among these candidates, white dwarfs and hot post-AGB stars have the advantage of being relatively well-understood phases of stellar evolution, but lack sufficient fuel to explain the strongest UV-upturn galaxies. The more promising candidates are currently the extreme horizontal-branch (EHB) stars and their hot progeny; the AGB-manqu\'e stars, which miss the AGB phase entirely, and the post-early AGB stars, which leave the AGB before the thermally pulsing phase. Although the horizontal branch (HB) morphology of the globular clusters becomes generally redder with increasing metallicity, a significant EHB population is theoretically expected at metallicities above solar (and an age larger than about 5 Gyr), if reasonable assumptions are made concerning the metallicity dependence of the mass loss rate and helium abundance (\cite{dorm95}, \cite{yi97a}). Spectroscopically, EHB and AGB-manqu\'e stars are expected to appear, respectively, as sdB and sdO stars, and the measured effective temperatures and gravities of the field hot subdwarf stars is consistent with this scenario (\cite{saff94}). However, the field hot subdwarfs apparently have a large binary fraction (\cite{theis95}, \cite{green98b}), which instead suggests that binary interactions might be involved in their formation (\cite{bailyn95}). Despite their probable contributions to the ultraviolet spectra of galaxies, blue stragglers and hot subdwarfs are either neglected or poorly constrained in existing population synthesis models of old stellar populations. The currently favored explanations for the origin of blue stragglers involve binary interactions, either through mass-transfer processes, binary mergers, or direct stellar collisions (\cite{bailyn95}, \cite{leonard96}), but the large parameter space of these binary processes limits their inclusion in synthesis models to special cases (e.g.\ \cite{pols94}). Some attempts have been made to include EHB stars in synthesis models of metal-rich populations, but these are hampered by large uncertainties, for example, in the dependence of mass loss and helium abundance on metallicity (\cite{bress94}, \cite{yi97b}). The modeling of galaxy ultraviolet spectra should clearly benefit from empirical studies of the ultraviolet content of resolved, old, metal-rich stellar populations. Here we report on ultraviolet ($\sim$ 1600~\AA) images of the three old ($> 4$ Gyr) open clusters M67, NGC 188, and NGC 6791, which were obtained in March 1995 using the {\em Ultraviolet Imaging Telescope} (UIT, \cite{stech97}) during the second flight of the {\em Astro} observatory. The $40'$ diameter field of view of UIT is sufficiently large to allow a complete census of the hot stellar population in the observed clusters. Thus, the UIT images of these three clusters can be used for an empirical study of the ultraviolet content of old, metal-rich stellar populations. As it turns out, M67 contains blue stragglers, but no hot subdwarfs, whereas NGC 188 and NGC 6791 contain hot subdwarfs, but are too old, distant and reddened for the blue stragglers to be detectable with UIT. M67 is also sufficiently nearby that all the white dwarfs hotter than $\sim$ 21,000 K should have been detected on the UIT image. In the population synthesis models of \markcite{magris93}Magris \& Bruzual (1993), white dwarfs supply about 10\% of the 1750~\AA\ flux in elliptical galaxies with a weak UV-upturn, and in Section 3.2 we re-examine the white dwarf contribution to galaxy ultraviolet spectra, using more recent stellar atmospheres and evolutionary models. A brief summary of this work was given by \markcite{last97}Landsman \& Stecher (1997). A subsequent paper will discuss the UIT observations of the intermediate age ($\sim$ 2 Gyr) clusters NGC 752 and NGC 7789, which are sufficiently young to allow the turnoff population to be detected in the ultraviolet. | There are at least three ways in which Galactic open clusters are an inadequate model for understanding the stellar populations of old galaxies. First, the open clusters have experienced dynamical interactions which can alter the stellar population mix; for example, by leading to evaporation of low-mass stars while concentrating the blue stragglers toward the cluster center. Second, the higher mean stellar density of open clusters might result in enhanced stellar interactions and possibly increased blue straggler or hot subdwarf formation. Finally, due to the relatively small population of open clusters, the rapid evolutionary phases may not be adequately sampled. For example, hot post-AGB stars are believed to be a significant (although not the dominant) contributor to the observed UV-upturn in elliptical galaxies, but consistent with their short ($\sim 10^5$ yr) lifetime, there are no post-AGB stars present in M67, NGC 188, or NGC 6791. And while hot subdwarfs have now been detected in NGC 188 and NGC 6791, their total number ($<8$) remains uncomfortably low to draw any robust conclusions. Keeping this limitations in mind, we summarize below the main results derived from the UIT observations of M67, NGC 188, and NGC 6791, and their implications for the study of the integrated ultraviolet spectra of old stellar populations. \begin{enumerate} \item The UIT image of M67 is dominated by the eleven detected blue stragglers; in particular, the blue straggler F81 contributes 90\% of the integrated flux of M67 at 1520~\AA. The ultraviolet flux of the two blue stragglers, F90 and F131, is significantly higher than that predicted on the basis of optical photometry, and probably indicates the presence of hot, subluminous companions. A semi-empirical calculation of the integrated ultraviolet spectrum of M67 shows that, even when neglecting the possibly anomalous star F81, the blue stragglers dominate at wavelengths shorter than $\sim$ 2600~\AA. As pointed out by Spinrad et al.\ (1997), neglect of blue stragglers in population synthesis model fits of the rest-frame near-ultraviolet spectra of a galaxy will result in an underestimate of the galaxy age. \item Eight white dwarf candidates are identified in the UIT image of M67, including the core-helium white dwarf companion of the yellow giant S1040. Optical spectroscopy of two of these sources has been obtained by \markcite{flem97}Fleming et al.\ (1997): G152 is a hot (\teff $\sim$ 68,000 K) DA white dwarf, and MMJ 5973 is cooler (\teff\ $\sim$ 18,000 K) DB white dwarf. The number of white dwarf candidates is in reasonable agreement with that expected from theoretical white dwarf cooling models and a cluster age of 4 Gyr. The integrated ultraviolet flux at 1500~\AA\ along a white dwarf cooling model is $\log$ E$_{1500} = -4.23$ L$_{V}^{\odot}$ Gyr \AA$^{-1}$, and the contribution of white dwarfs to the integrated spectra of old galaxies is roughly 10\% of that expected from hot post-AGB stars. \item Eight probable cluster members are detected on the UIT image of NGC 6791, including the five hot subdwarfs studied spectroscopically by Liebert et al.\ (1994), and the two additional sdB/O candidates, B9 and B10, reported by \markcite{kal95}Kaluzny \& Rucinski (1995). Three probable cluster members are detected on the UIT image of NGC 188, including the sdB spectroscopic binary II-91. The star D702 in NGC 188 is probably a composite including a hot subdwarf, since it has a hot ultraviolet color (\mbi\ -- \mbv\ = --0.11), but a relatively cool optical color (\bv\ = 0.26). The derived luminosity range, $1.10 < \log$ L/\lsun\ $< 1.27$, of the five sdB stars in NGC 6791, and II-91 in NGC 188, is consistent with that expected for metal-rich, hot HB stars. The fraction of hot HB stars in both clusters is about 30\%, implying that the integrated spectra of both clusters should show a pronounced UV-upturn, as strong as that observed in any elliptical galaxy. \end{enumerate} | 98 | 4 | astro-ph9804164_arXiv.txt |
9804 | astro-ph9804328_arXiv.txt | Galaxy distance indicators are subject to different uncertainties and biases which may reflect in the peculiar velocity field. In this work, we present different statistical analysis applied to peculiar velocities of samples of galaxies and mock catalogs taken from numerical simulations considering these observational effects. Our statistical studies comprise the mean relative velocity and velocity correlation for pairs of galaxies as a function of separation, and a bulk flow analysis determined on spheres of $10 h^{-1}$ $Mpc$ radius. In order to test the statistical analysis we use COBE normalized CDM numerical simulations with different density parameters and cosmological constant and take into account the Tully-Fisher (TF) scatter and a possible TF zero-point offset, as well as variations in the results of the simulations arising from different observer positions. We compare the results of the mock catalogs with samples of spiral galaxies taken from the Mark III catalog. The results of our analysis indicate the importance of errors in deriving the density parameter $\Omega$ from statistics applied to the observational data and it is found that large values $\Omega \geq 0.8$ may be derived from the analysis if errors are not taken into account. However, the observed value of the TF scatter ($\simeq 0.35$ $mag$) make CDM models with $\Omega > 0.6$, inconsistent with the statistical tests involving relative velocities. A suitable agreement is found for models with $\Omega \leq 0.6$, although requiring a TF zero-point offset of the order of a tenth of a magnitude to provide consistency with the observed flow coherence. | Measurements of peculiar velocities provide direct probes of the mass distribution in the Universe putting constraints on models of large-scale structure formation (e.g. \cite{peebles}, \cite{vittorio}). Advances on new distance indicators (\cite{djor}; \cite{dresslera}) have allowed estimates of peculiar flows in the local Universe up to $\sim 50-100 h^{-1}$Mpc (see \cite{giova}, or \cite{strauss} for a review). The results from using elliptical galaxies (\cite{dresslerb}) suggest a large coherence length and amplitude of the local peculiar velocity field. The results from analysis of 1,355 spirals using the Tully-Fisher relation to determine distances (\cite{mathew}) would lead to an even larger coherence length and a similar amplitude. Other samples are also in agreement with the Great Attractor findings (\cite{willick}). In a more controversial finding \cite{postman} find from analysis of clusters of galaxies that the peculiar flows can be coherent over a much larger scale than the Great Attractor findings. All the results indicate significant deviations from the Hubble flow with a large coherence length. Since peculiar velocities probe the mass distribution and also depend on the density parameter $\Omega$ one can determine the latter by comparing the mass distribution implied by the velocity field with the observed distribution of galaxies. The density parameter is thus determined to within the uncertainty of the bias factor $b$ by determining the factor $\Omega^{0.6}/b$. Other analysis of the peculiar velocity information may be applied in order to obtain the parameter $\Omega$. \cite{berta} developed the POTENT method whereby the mass distribution is reconstructed by using the analog of the Bernoulli equation for irrotational flows. They used the method to analyze to great detail the peculiar velocity field out to about 60$h^{-1}$Mpc (\cite{bertb}). \cite{dekel} compared the previously determined velocity field with the observed distribution of galaxies and concluded that the best fit is with $\Omega^{0.6}/b \simeq 1$. A similar analysis can be applied to the components of the velocity tensor $U_{ij}$ (\cite{gorskia}; \cite{groth}). Recently, \cite{kashlinsky} has analyzed this tensor for the MarkIII peculiar velocity data to reconstruct the large-scale power spectrum obtaining consistency within the CDM model with $\sigma_{8} \Omega^{0.6} \simeq 0.8$. In a similar study \cite{zaroubi} obtain a different result $\sigma_{8} \Omega^{0.6} \simeq 0.35$ indicating the need of further analysis. In this paper we study the peculiar velocity field through a statistical analysis of COBE normalized CDM numerical simulations and observational data taken from the Mark III catalog. Our comparison of models and observations take into account observational uncertainties and possible biases, and variations of the results according to different observers in fully non-linear numerical simulations. \section {Data} We use the Mark III catalog (\cite{markIII1};\cite{markIII2};\cite{markIII3}) as a suitable data set to analyze the peculiar velocity flow. This Catalog lists Tully-Fisher and $Dn-\sigma$ distances and radial velocities for spiral, irregular, and elliptical galaxies. In our analysis we have used only spiral galaxies given their large number and smooth spatial distribution (see Table 1). \begin{deluxetable}{cccc} \tablewidth{35pc} \tablecaption{Observations: The Mark III spirals} \tablehead{ \colhead{Subsample}& \colhead{$N^o$ of Gx.}& \colhead{TF relation}& \colhead{$\sigma_{TF}$} } \startdata \phm{II} Aaronson et al. Field & 359 & $M_H=-5.95+10.29 \eta$ & $0.47$ \nl \phm{II} Mathewson et al. (1992) & 1355 & $M_I=-5.79+6.8 \eta$ & $0.43$ \nl \phm{II} Willick, Perseus Pisces (1991) & 383 & $M_r=-4.28+7.12 \eta$ & 0.38 \nl \phm{II} Willick, Cluster Galaxy (1991) & 156 & $M_r=-4.18+7.73 \eta$ & $0.38$ \nl \phm{II} Courteau-Faber (1993) & 326 & $M_r=-4.22+7.73 \eta$ & $0.38$\nl \phm{II} Han-Mould et al., Cl. Gx. (1992) & 433 & $M_I=-5.48+7.87 \eta$ & 0.4 \nl \enddata \end{deluxetable} The velocity parameter $\eta = Log \Delta V - 2.5$ is determined either from HI profiles or from optical $H_{\alpha}$ rotation curves. The TF relations and their corresponding scatters for the different samples of spiral galaxies are given by \cite{markIII3} and are shown in Table 1, where the absolute magnitud $M$ satisfies $M=m-5\log cz$. The galaxy apparent magnitudes $m$ of the Tully-Fisher distances are corrected for Galactic extinction, inclination and redshift (see Willick et al. 1997 for details). The selection bias in the calibration of the forward TF relation can be corrected once the selection function is known. But then, the TF inferred distances and the mean peculiar velocities suffer from Malmquist bias. Suitable procedures to consider these biases, induced both by inhomogeneities and selection function, have been discussed (see for instance /cite{freud95} and references therein) where the spatial distribution, selection effects and observational uncertainties are realistically modeled through Monte-Carlo simulations. We have used in our analysis forward TF distances, fully corrected for Inhomogeneous Malmquist Bias (\cite{markIII1}, \cite{markIII2}, \cite{markIII3}). Nevertheless we have found that the results of the statistics studied in this work do not change significantly if inverse TF distances are used as it will be discussed below. Radial velocities used to infer the peculiar velocity of the galaxies are referred to the Cosmic Microwave Background frame. It should be remarked that galaxy distance estimates are subject to errors due to the scatter in the TF relation (\cite{mo}; \cite{willickth}; \cite{mathew}) and uncertainties of the TF zero-point (\cite{shanks}; \cite{willickth}). Also, the possible presence of a small fraction of spurious velocities in the data induced by either galaxy peculiarities or observational errors in distance estimates (\cite{jacoby}) should be taken into account. | In this work we have attempted to asses the effects of galaxy distance uncertainties and biases on statistical tests of the peculiar velocity field using samples of spiral galaxies from the Mark III compilation and mock catalogs taken from numerical simulations. Our studies comprise relative velocity tests and pair velocity correlations as a function of separation, as well as a bulk flow analysis determined on spheres of $10 h^{-1}$ $Mpc$ radius. We construct mock catalogs using numerical simulations corresponding to COBE normalized CDM models with different values of density parameter and cosmological constant. The models take into account the Tully-Fisher (TF) scatter and possible TF zero-point shift, as well as variations in the results of the simulations arising from different observer positions. The analysis of the departures from gaussianity of the observational velocity distributions as measured by the kurtosis tests $K$ and $K1$ show that only a small fraction ($<2 \%)$ of possible spuriously high velocities (and therefore biased distance estimates) might be present in the observational data. We find that the observed scatter of the TF relation plays a very important role when deriving constraints to the cosmological parameters. If errors in galaxy distance estimates are neglected, the observed magnitude of the $\Delta V$ results show consistency with high density CDM models ($\Omega>0.8$). When the observed TF scatter $\Delta M \simeq 0.35$ is taken into account a significant disagreement with observations is found for the $\Delta V$ and $D1$ statistics in models with $\Omega > 0.6$. A possible offset in the TF zero-point used to determine the distances of the galaxies in the observational data artificially enhance velocity correlations ($\Pi$ and $M$ statistics). Due to this fact, if zero-point offsets of the order of a tenth of a magnitude are considered, the CDM models explored provide a more satisfactory fit to observations in these tests. The results presented in this work are not sensitive to sample variations. We have applied the same statistical tests to the subsample corresponding to Mathewson data and we find similar results than those of the total sample. It should be noted that the variations on the statistics arising from different observers are significantly enhanced when errors on the peculiar velocities are included. This fact make more difficult the distinction between the models. Error bars in figures 1 to 5 show the variations of the results arising from different observer positions in the models and serve as a test for the dependence of the observational results on our particular position in space. In our analysis of biased models we have not found relevant differences in the results of the statistical analysis of samples with different density thresholds. The selection of low (high) density particles in the simulations does not produce very relevant effects in the statistics although lower (higher) values of $\Pi$ and $\Delta V$ are observed due to the oversampling (undersampling) of high density particles where random motions dominate. Finally when comparing models and observations, we find that neither the K, K1, M, nor $\Pi$ statistical tests are sensitive to the density parameter of the models. We find crucial to consider properly the intrinsic scatter of the Tully-Fisher relation in studies of the peculiar velocity field. Although low values of the density parameter $\Omega < 0.6 $, are favored by our statistical analysis $D1$ and $\Delta V$ when this scatter is taken into account $\pm 0.15$ $mag$ zero-point offsets in the TF relation are required in order to provide the observed values of $\Pi$ and $M$ statistics in the models explored. A positive TF zero-point offset of this magnitude would imply lower values of the Hubble constant which may be worth to consider in a controversial topic (see for instance \cite{shanks} and \cite{th} and references therein). | 98 | 4 | astro-ph9804328_arXiv.txt |
9804 | astro-ph9804091_arXiv.txt | The conversion of neutron matter into strange matter in a neutron star occurs through the non-leptonic weak-interaction process. We study the energy loss of the neutron star by the emission of axions in that process. Owing to that process, the neutron star will liberate the energy which can in no way be negligible as an axion burst. \\ \\ PACS numbers : 14.80.Mz, 95.30.Cq, 97.60.Jd \\ Keywords : axions, astrophysics, neutron stars, strange star \\ \\ | The possibility of the existence of stable quark matter in the early universe or inside a neutron star or in relativistic heavy-ion collision experiments has been studied in many works \cite{sqm}\@. In this work, we consider the conversion of a neutron star into a strange star and its energy loss during the conversion process. In order to form strange matter in the interior of a neutron star, neutron matter should be converted into strange matter. Several conversion mechanisms have been discussed by Alcock et al.\ \cite{alcock}\@. When a neutron and a stable strangelet (strange quark matter droplet) meet, the neutron is readily absorbed, while a proton can coexist with a strangelet due to the Coulomb barrier between them \cite{alcock}\@. Therefore, a stable strangelet as a seed for the process will trigger a conversion in such a way that it grows by absorbing neutrons and, eventually, convert most of the neutron star into a strange star \cite{olinto,daietal}\@. Then the conversion would liberate about $10^{52}$ ergs in binding energy \cite{olinto}\@. A variety of mechanisms have been suggested for seeding the interior of a neutron star with stable strangelets so far \cite{alcock,lugones}\@. They can be divided into two main categories \cite{olinto}; (1) the primary mechanisms in which the seed is formed inside the neutron star, (2) the secondary mechanisms in which the seed comes from the interstellar medium. In a viewpoint of the primary mechanisms \cite{alcock}, the central high densities and pressures in a neutron star may be sufficient for a phase transition from neutron matter to two-flavor quark matter to occur. Subsequently two-flavor quark matter can then easily decay into the lower energy strange matter through weak interactions. If we assume that strangelet has formed inside a neutron star, subsequent conversion process of the rest of the star is described as follows \cite{olinto,olesen}\@. The volume over which strange matter equilibrates was shown to be much smaller than that of the total strange matter region \cite{lugones,horvath}, so that the problem can be treated hydrodynamically in one-dimensional geometry. As the conversion front sweeps into neutron matter, the small region behind the conversion front has an excess of down quarks relative to strange quarks due to the flux of neutrons ($udd$) at the conversion front. The excess down quarks will convert into strange quarks via the non-leptonic weak process \cite{madsen}, $d \; + \; u \; \rightarrow \; s \; + \; u$, as long as $\mu_d > \mu_s (n_d > n_s)$\@. By this process a $d$ quark can change itself to an $s$ quark until the Fermi energies of all the flavors become equal and the energy per baryon drops in comparison to the ordinary two-component nuclear matter. Other leptonic decay processes are suppressed considerably \cite{madsen}\@. The conversion will liberate $\sim 10^{52}$ ergs of energy (assuming that $\sim$ 10 $MeV$ is liberated per neutron converted)\@. This energy will be radiated as neutrinos, photons, $e^{+}e^{-}$ pairs, and axions, etc. In this stage neutrinos and axions can escape the star. Many authors \cite{daietal,aggs} have considered the escape of neutrinos. In this work we calculate the energy loss due to the emission of axions in the process of such a conversion of a neutron star into a strange star. At the chemical equilibrium between the quarks and the electrons, non- or semi-leptonic weak interaction is not important in neutrino and axion emission in quark matter because the weak interactions coupling $s$ and $u$ quarks are Cabibbo suppressed relative to the interactions coupling $d$ and $u$ quarks. For the nonequilibrium processes such as the conversion from two-flavor to three-flavor quark matter, however, as the strange quark semi-leptonic processes in neutrino emissions and non-leptonic processes in axion emissions are of significance. The axion is a pseudo-Goldstone boson which was introduced by Peccei and Quinn (PQ) to solve the strong CP-problem in a natural way \cite{PQ}\@. However, theoretical and experimental investigations give little guidance on the PQ symmetry-breaking scale, $f_a$, and therefore on the mass of the axion \cite{WW}\@. The axion decay constant $f_a$ is related to the axion mass \cite{axion} \[ m_a \simeq \left( \frac{0.62 \times 10^{7} GeV}{f_a} \right) \; eV . \] There are two generic types of invisible axions; the Dine-Fischler-Srednicki-Zhitnitskii (DFSZ) axion which couples to both quarks and leptons at tree level \cite{DFSZ} and the hadronic, or Kim-Shifman-Vainshtein-Zakharov (KSVZ), axion which has no tree-level coupling to leptons but does couple to them at loop levels \cite{KSVZ}\@. In invisible axion models, the axion mass $m_a$ is in principle arbitrary, however astrophysical and cosmological considerations \cite{axion} can provide an upper and lower bounds for $m_a$\@. The astrophysical bounds on $m_a$ are due to the fact that axion emission is an additional energy loss mechanism for stars. If such axions could be copiously produced during the conversion of a neutron star into a strange star, it might drastically alter the energy budget of stars. The axions might thus deplete the stellar energy and change the usual course of stellar evolution. The emission of axions would have hastened the cooling process. The quarks are expected to be highly relativistic and degenerate in such matter with their Fermi energies in the range of 300 to 500 $MeV$\@. A calculation of the axion energy flux from strange quark matter has been made by Anand et al.\ \cite{anand}\@. They have shown that the axion emission rate is several orders of magnitude smaller compare to the neutrino emission rate. We investigate here the energy loss of the neutron star during the conversion of non-strange quark matter into strange quark matter in the interior of the star. We assume that the axions escape freely out of the star as soon as they are produced, and that the temperature in the interior is constant during the conversion for the reasons that strong and electromagnetic interaction timescales are much smaller than those of weak interactions. Unless otherwise noted all equations assume $\hbar = c = k_B = 1$\@. | Our main result for the total energy flux rate for the emission of axions from conversion of non-strange quark matter into strange quark matter is given by Eq. (12). Here we can carefully quote our bound on the axion coupling $f_a$, using the simple inequality \cite{ishizuka}, \be {\cal E}_a \cdot V \cdot \delta t < E_m \ee The allowed energy loss, $E_m / \delta t$ is taken to be $10^{52}$ erg/sec for definiteness, and the volume of axion emitting region $V = \frac{4}{3} \pi R^3$ with $R = 10 \sim 20 km$\@. As for the conversion timescale, if $T=10 MeV$, it varies roughly between one and ten minutes \cite{olinto}\@. So, we take the timescale typically $\delta t = 100 sec$\@. Then, from Eq. (14), we can obtain the bound on $f_a$, \be \sim 10^5 \; GeV < f_a \,. \ee We note that the bound based on Eq. (14) is legitimate only if the emitted axions never interact on their way out of the star. If the axion coupling is large enough, axions once produced may interact many times, namely be absorbed and reemitted more than once in the hot core of the star, and may be trapped thermally. As applications, Mikheev and Vassilevskaya \cite{mik} recently investigated the radiative decay of the axion $a \rightarrow \gamma\gamma$ in an external electromagnetic field in the DFSZ model. They concluded that the decay probability is strongly catalyzed by the external field, namely, the field removes the main suppression caused by the smallness of the axion mass. Therefore, the radiative decay of the axion in strong magnetic fields of the neutron star could give interesting astrophysical phenomena. Furthermore, we can consider that the sudden conversion from a neutron star to a strange star may account for the gamma ray bursts at comological distances \cite{olinto}\@. Most of the energy is probably released in the form of neutrinos. If a part of the total energy goes into $\gamma$ rays decayed from the axions, it will be large enough to account for the gamma ray bursts at cosmological distances and to explain their isotropic distribution. The outcome of such a conversion event will be the emission of as much as $\sim 10^{58} MeV$ of energy in $\sim sec$ to $\sim min$ as radiation with a typical temperature of tens of $MeV$\@. This conversion can be observed as a gamma ray burst \cite{olinto}\@. In summary, as a possible mechanism for production of axions, we considered the conversion of non-strange quark matter into strange quark matter. We also estimated the energy loss of neutron stars through the emission of axions in addition to the cooling provided by the neutrino emission. During the conversion period, the important process is the non-leptonic weak-interaction. We assume that the axions are not trapped and escape freely out of the star as soon as they are produced, and that the temperature in the interior is constant during the conversion. We find that the energy carried away by axions during the conversion is not negligible. It is also found that the axion emission rate is three orders of magnitude larger compared to the chemical equilibrium case. In addition, we discussed the bound on the axion coupling. \vspace*{3cm} | 98 | 4 | astro-ph9804091_arXiv.txt |
9804 | astro-ph9804058_arXiv.txt | We have analyzed deep $B$ and $V$ photometry of the Carina dwarf spheroidal reaching below the old main-sequence turnoff to $V \sim 25$. Using simulated color-magnitude diagrams to model a range of star formation scenarios, we have extracted a detailed, global star formation history. Carina experienced three significant episodes of star formation at $\sim 15$ Gyr, 7 Gyr, and 3 Gyr. Contrary to the generic picture of galaxy evolution, however, the bulk of star formation, at least 50\%, occured during the episode 7 Gyr ago, which may have lasted as long as 2 Gyr. For unknown reasons, Carina formed only 10-20\% of its stars at an ancient epoch and then remained quiescent for more than 4 Gyr. The remainder ($\sim 30\%$) formed relatively recently, only 3 Gyr ago. Interest in the local population of dwarf galaxies has increased lately due to their potential importance in the understanding of faint galaxy counts. We surmise that objects like Carina, which exhibits the most extreme episodic behavior of any of the dwarf spheroidal companions to the Galaxy, are capable of contributing to the observed excess of blue galaxies at $B \sim 24$ only if the star formation occurred instantaneously. | The Carina dwarf galaxy, one of the nine known dwarf spheroidal companions to the Galaxy, has a surprisingly complex star formation history. Detection of carbon stars provided the first suggestion of a significant intermediate-age population (Cannon, Niss, \& Norgard-Nielsen 1981), followed by main-sequence photometry which revealed a young turn-off due to a population perhaps only 6--9 Gyr old (Mould \& Aaronson 1983; MA hereafter). MA fit a simulated luminosity function comprised of 7 Gyr-old stars to their data, and estimated the contribution of an old population to be relatively small in comparison. Saha, Monet, and Seitzer (1986) observed a large number of RR Lyraes in the central region of the galaxy and established a lower limit of 2--3\% for the fraction of old stars, depending on the yield of RR Lyraes. Mighell (1990) estimated the relative sizes of the populations from the double-peaked color distribution near the MSTO region to be 85\% for the intermediate-age burst and 15\% for the old episode. Using simulated luminosity functions, Mighell \& Butcher (1992) later fit intermediate-age burst models to this deep, main-sequence photometry and estimated an upper limit to the old population of 40\%. More recently, Smecker-Hane, et al. (1994; hereafter SHSHL) resolved two separate horizontal branches (HBs). Although the separation is not independent of metallicity effects, it is very suggestive of the multi-episode nature of Carina's history. We are left with an estimate between 2\% and 40\% for the fraction of old stars, and no clear result on the duration of the star formation episodes. The chemical evolution of the galaxy is likewise still in question. Spectroscopy of 15 giants in Carina (Da Costa 1994) resulted in an average [Fe/H] of --1.9, with one giant significantly more metal poor ([Fe/H] = --2.2). The excellent areal coverage of the SHSHL data reveal a very thin giant branch. Given the huge apparent age spread, a metallicity spread may be necessary to compensate and produce the observed narrow RGB. Improved methods of analysis are needed. Detailed analyses of stellar luminosity functions have some advantages over more traditional isochrone fitting. These methods, however, were designed for coeval systems. Because dwarfs exhibit a range of ages and metallicities, these approaches are problematic at best, and misleading at worst. A much better way to unravel complex star formation histories like those exhibited by dwarf galaxies such as Carina and Leo~I is to use all of the information embodied by color-magnitude diagrams. In these cases, where the galaxies appear to have experienced bursts at intermediate epochs, a conventional luminosity function would not distinguish between the old subgiant branch and the young MS stars, for example. Color information proves essential to resolve multiple components in these systems. Simulated color-magnitude diagrams have been used successfully in studies of bright stars in dwarf irregulars (Tosi et al. 1991; Tolstoy 1996; Gallart et al. 1996a, 1996b, Aparicio et al. 1997a, 1997b), of LMC field stars (Bertelli et al. 1992; Vallenari et al. 1996), and of LMC clusters (Vallenari et al. 1994). We have deep CCD photometry (limiting magnitude $V = 24.5$) with reasonably small photometric errors ($\sim 0.02$ at $V=23$) of three fields in the Carina dwarf galaxy, reaching well below the old main-sequence turnoff. Because the photometry reaches below the old MSTO, these data are well-suited to a detailed extraction of the star formation history by comparing model color-magnitude diagrams to that of Carina. Section 2 is a discussion of the observations and reductions of the Carina data. Section 3 outlines the analysis applied to these data, including the development of the pseudo-LF and the generation of similulated color-magnitude diagrams. The results, presented in section 4, are summarized and briefly discussed in the context of galaxy evolution in section 5. | Carina experienced an episode of star formation 7 Gyr ago which lasted no more than 2 Gyr and which was responsible for at least 50\% of its stars. The old population ($12-15$ Gyr) may amount to $10\%-20\%$ of Carina. The bulk of the remainder (($20\%-30\%$) is relatively young, between 2.5 and 3.5 Gyr old. While the details of one galaxy's SF history may seem inconsequential, taken in context and as a member of the Local Group dwarf population, the details are relevant to deeper, unresolved issues in cosmology and galaxy evolution. In the case of Carina, all studies confirm that at least one relatively long pause in star formation lasting $\sim 4$ Gyr occurred. The mechanism by which this kind of low mass system could experience such a pause and then another strong burst of star formation is not understood. Current ISM simulations can produce such a large gap in an isolated dwarf only by ejecting the gas out to $\sim 20$ kpc (Babul 1996). A dwarf such as Carina, residing in the potential well of the Galaxy, should lose that gas, preventing any further episodes of star formation. Recent Hubble Space Telescope observations of Carina have been studied by Mighell (1997). He claims to detect a significant number of stars in the region of the gap in star formation lasting from roughly 8 to 12 Gyr ago. This is interpreted as indicative of significant star formation which began in the central region of Carina, and propagated to the outer regions. However, the number of stars in the same section of our CMD, derived from fields several arcminutes from the center of Carina and overlapping the Mighell (1990) field, is consistent with the number in the HST CMD. Whether or not the number of stars in this area implies significant on-going star formation, there does not seem to be evidence of a difference between the central region and regions further out, such as that seen in the recently-discovered Antlia dwarf galaxy (Aparicio et al 1997c). Assigning ages to individual faint stars based on the isochrones which constrain them in color and magnitude must be handled with care. Firstly, at faint magnitudes, stars have a larger photometric error and thus a larger implied age range. Age determinations for specific stars are therefore subject to greater statistical uncertainty. Secondly, the lifetime of stars crossing the Hertzprung gap increases as stars become less massive. This must be accounted for when predicting the relative contributions of different ages. In light of these considerations, the results of the WFPC2 study may not differ from earlier ground based results. Evidence of intermediate-age stars in the halo leads to the question of whether dwarf spheroidals shredded by the Galaxy could be a significant source of the halo population (Preston et al. 1994; Mateo 1996). The Sagittarius dSph is currently being ripped apart by the Milky Way, proving that this type of interaction between a large galaxy and its tiny neighbors does occur (Ibata, Gilmore, \& Irwin 1994; Mateo et al. 1995). In addition, the dSph, especially Carina and Leo I, contain significant numbers of intermediate-age stars. Mateo (1996) shows that the fraction of relatively young and intermediate-age stars in the entire population of dwarf spheroidals is not inconsistent with the fraction in the Galactic halo. The detailed interpretation of deep galaxy counts and redshift surveys depends on the SFH of the dwarfs which comprise the excess of faint galaxies. For example, `flashing' dwarfs - i.e. bursting quickly and intensely - contribute to the deep counts to a dramatically different extent than do dwarfs that experience more drawn-out bursts that slowly turn on and off (Campos 1997). The prevalence of bursting behavior in the SFH of Local Group dwarfs may statistically constrain the degree to which galaxies similar to these local dwarfs could be responsible for the faint excess. We can place an upper limit on the star formation rate (SFR) of Carina 7 Gyr ago by assuming that the episode of star formation took place instantaneously. By choosing a reasonable model, we can easily calculate the total number of stars formed during this episode in the volume of the galaxy covered by our three fields, taking into account the incompleteness. We can use this number to normalize the IMF, and then integrate to get the total mass formed during this episode. The SFR is then given by: \begin{displaymath} SFR = \frac{M_{total}}{\Delta t}*\frac{1}{\gamma } M_{\odot } yr^{-1} \end{displaymath} where $\Delta t$ is the duration of the episode in years and $\gamma$ is the fraction of the galaxy's total surface brightness in the three fields. If the episode lasted 10 Myr, the implied SFR for Carina at that time would be roughly 0.8 $M_{\odot } yr^{-1}$. If the episode lasted 2 Gyr, the SFR is down by more than a factor of 100 to 0.004 $M_{\odot } yr^{-1}$. Either SFR is comparable to the instantaneous SFR in even such luminous objects as blue compact dwarfs (Fanelli et al. 1988). Could Carina-type galaxies contribute to the excess of faint galaxies at around $B \sim 24$? These galaxies have been shown to have redshifts between 0.3 and 0.7 (Glazebrook et al. 1995). Assuming $H_{\circ}$ = 50 km/s/Mpc, and $\Omega$ = 1, 7 Gyr corresponds to $z \sim 0.5$, which places the episode of star formation in the appropriate epoch. If the burst were essentially instantaneous, the implied total luminosity of Carina at 7 Gyr would be $\sim 2 \times 10^8 L_\odot$. Carina would be visible to $z \sim 0.3$ in a sample limited to $V \sim 25$, putting it at the near edge of objects that could contribute to these counts. Galaxies 3-10 times more luminous would fall in the range $z \sim 0.5$ to 1.0. Although the local dSph cover a range up to 50 times the luminosity of Carina, Carina is the most extreme example of this type of episodic behavior. Further, if the episodes of star formation are extended in time, then these dwarfs would only be visible locally. Thus, despite concerted efforts at unravelling the SFH of these galaxies, their role in faint galaxy counts problem remains unresolved. | 98 | 4 | astro-ph9804058_arXiv.txt |
9804 | astro-ph9804172.txt | We have analyzed the projected galaxy distributions in a subset of the ENACS cluster sample, viz. in those 77 clusters that have z $<$ 0.1 and R$_{\rm ACO} \ge$ 1 and for which ENACS and COSMOS data are available. For 20~\% of these, the distribution of galaxies in the COSMOS catalogue does not allow a reliable centre position to be determined. For the other 62 clusters, we first determined the centre and elongation of the galaxy distribution. Subsequently, we made Maximum-Likelihood fits to the distribution of COSMOS galaxies for 4 theoretical profiles, two with `cores' (generalized King- and Hubble-profiles) and two with `cusps' (generalized Navarro, Frenk and White, or NFW, and de~Vaucouleurs profiles). We obtain average core radii (or characteristic radii for the profiles without core) of 128, 189, 292 and 1582 kpc for fits with King, Hubble, NFW and de~Vaucouleurs profiles respectively, with dispersions around these average values of 88, 116, 191 and 771 kpc. The surface density of background galaxies is about 4 10$^{-5}$ gals arcsec$^{-2}$ (with a spread of about 2 10$^{-5}$), and there is very good agreement between the values found for the 4 profiles. There is also very good agreement on the outer logarithmic slope of the projected galaxy distribution, which is that for the non-generalized King- and Hubble-profile (i.e. $\beta_{King}$ = $\beta_{Hubble}$ = 1, with the corresponding values for the two other model-profiles). We use the Likelihood ratio to investigate whether the observations are significantly better described by profiles with cusps or by profiles with cores. Taking the King and NFW profiles as `model' of either class, we find that about 75 \% of the clusters are better fit by the King profile than by the NFW profile. However, for the individual clusters the preference for the King profile is rarely significant at a confidence level of more than 90 \%. When we limit ourselves to the central regions it appears that the signifance increases drastically, with 65 \% of the clusters showing a strong preference for a King over an NFW profile. At the same time, about 10 \% of the clusters are clearly better fitted by an NFW profile than by a King profile in their centres. We constructed composite clusters from the COSMOS and ENACS data, taking special care to avoid the creation of artificial cusps (due to ellipticity), and the destruction of real cusps (due to non-perfect centering). When adding the galaxy distributions to produce a composite cluster, we either applied no scaling of the projected distances, scaling with the core radii of the individual clusters or scaling with r$_{200}$, which is designed to take differences in mass into account. In all three cases we find that the King profile is clearly preferred (at more than 95 \% confidence) over the NFW profile (over the entire aperture of 5 core-radii). However, this `preference' is not shared by the brightest (M$_{b_j}$ $\la$ -18.4) galaxies. We conclude that the brighter galaxies are represented almost equally well by King and NFW profiles, but that the distribution of the fainter galaxies clearly shows a core rather than a cusp. Finally, we compared the outer slope of the galaxy distributions in our clusters with results for model calculations for various choices of fluctuation spectrum and cosmological parameters. We conclude that the observed profile slope indicates a low value for $\Omega_0$. This is consistent with the direct estimate of $\Omega_0$ based on the $\frac ML$-ratios of the individual clusters. | Until fairly recently, the projected galaxy density in rich galaxy clusters was generally described by King or Hubble profiles. In these profiles, the logarithmic slope of the mass distribution is essentially zero near the cluster centre. The core radius which is the characteristic scale of the distribution, was sometimes also regarded as the distance which more or less separates dynamically distinct regions in a cluster. From the kinematics of the galaxy population it appears that in clusters the relaxation time is significantly shorter than the Hubble time {\em only} in the very central region within at most a few core radii (see e.g. den Hartog and Katgert 1996). The concept of cores in clusters has been seriously challenged, on observational grounds (e.g. Beers $\&$ Tonry 1986) and as a result of numerical simulations. Navarro, Frenk and White (1995, 1996) found e.g. that the equilibrium density profiles of dark matter halos in universes with dominant hierarchical clustering all have the same shape, which is essentially independent of the mass of the halo, the spectrum of initial density fluctuations, or the values of the cosmological parameters. This `universal' density profile (NFW profile hereafter) does not have a core, but has a logarithmic slope of --1 near the centre which, at large radii, steepens to --3, and thus closely resembles the Hernquist (1990) profile except for the steeper slope of the latter at large radii of --4. Navarro, Frenk and White (1997) argue that the apparent variations in profile shape, as reported before, can be understood as being due to differences in the characteristic density (or mass) of the halo, which sets the linear scale at which the transition of the flat central slope to the steep outer slope occurs. They also argued that the existence of giant arcs in clusters requires that the mass distributions in clusters does not exhibit a flat core in the centre. In other words: if clusters have cores, the lensing results require that the core radii are very small, at least quite a bit smaller than the values usually quoted. It is not clear that galaxy clusters should have cores; after all, the dynamical structure of galaxy clusters is quite different from that of globular clusters, for which Michie \& Bodenheimer (1963) and King first proposed density profiles with cores, in particular the King profile (see e.g. King 1962). On the other hand, the X-ray data for clusters are quite consistent with the existence of a core in the density distribution. More specifically, it was argued recently by Hughes (1997) that the NFW profile would induce a temperature gradient. The existence of such a gradient in the Coma cluster can be excluded at the 99$\%$ confidence level. Similarly, the galaxy surface density in clusters is generally found to be consistent with a King profile. For galaxy clusters, little use has been made of the de Vaucouleurs profile to describe the galaxy density, even though the latter was found to arise quite naturally in N-body simulations of the collapse of isolated galaxy systems (e.g. van Albada 1982). In view of the claimed universality of the NFW profile found in the simulations, it seems useful to have a closer look at the projected distribution of the galaxies in clusters. After all, the NFW profile refers to the total gravitating mass, and it is not obvious that the galaxy distribution should have exactly the same shape as the distribution of total mass; although in numerical experiments no strong biasing between dark and luminous matter in clusters was seen (e.g. van Kampen 1995). In this respect, it is noteworthy that Carlberg et al. (1997) find that the combined galaxy density profile of 16 high-luminosity X-ray clusters at a redshift of $\approx$0.3 closely follows the NFW profile. More precisely, the logarithmic slope in the central region is consistent with the value of --1 of the NFW and Hernquist profiles, while the outer slope is consistent with both --3 (the NFW value) and --4 (the Hernquist value). The outer slope of the density profile was found by several authors (e.g. Crone et al. 1994, Jing et al. 1995, and Walter and Klypin 1996) to reflect the details of the formation scenario, and in particular the value of the density parameter of the universe. In addition, this slope is unlikely to be constant in time but is expected to get steeper with decreasing redshift. For that reason, it is important to study both the characteristics of the density profiles of rich clusters and their dependence on redshift. In this paper, we investigate the projected galaxy distributions for a sample of 62 rich and nearby (z $\la$ 0.1) clusters. These clusters are taken from the volume-limited ENACS (ESO Nearby Abell Cluster Survey) sample of R$_{\rm ACO} \ge 1$ clusters (see e.g. Katgert et al. 1996 (paper I), Mazure et al. 1996 (paper II), Biviano et al. 1997 (paper III), Adami et al. 1998 (paper IV), Katgert et al. 1998 (paper V) and de Theije $\&$ Katgert 1998 (paper VI)). In $\S$ 2, we first describe the sample of clusters that we used, and the data on which we based our analysis. In $\S$ 3, we discuss the results of Maximum-Likelihood fitting of profiles with and without a core, to the individual galaxy distributions taken from the COSMOS catalogue. In $\S$ 4 we discuss the galaxy density distribution for composite clusters (COSMOS and ENACS), in which the individual clusters are combined. In $\S$ 5 we discuss the constraints provided by the outer slope of the density distributions for the parameters of the formation scenario and in $\S$ 6 we present the conclusions. | We have studied the projected galaxy distributions in 77 clusters from the ESO Nearby Abell Cluster Survey. The present sample is an unbiased subset of the volume-limited ENACS sample, and thus forms a representative local (z $<$ 0.1) sample of rich (R$_{ACO}$ $>$ 1), optically selected clusters. We used both COSMOS and ENACS data to test the character of the projected galaxy distributions. In particular, we have investigated whether the galaxy distributions in rich clusters have cusps or cores in their central regions. We have made maximum Likelihood fits to the observed distribution of COSMOS galaxies to solve for the position and the elongation of the clusters. For 15 of the 77 clusters, no reliable centre could be determined and these clusters were not considered further. Using the positions and elongations, we subsequently solved for each of the 62 remaining clusters the three parameters that describe each of the four theoretical profiles that we tested, as well as the density of background galaxies. The four model profiles that we tested against the data are the King, Hubble, NFW (Navarro, Frenk and White) and the de~Vaucouleurs profiles. Although the solutions do not converge for all of the clusters nor for all four profiles, we obtain reliable results for between 75 and 95 \% of the clusters (depending on the model profile). We find mean values for r$_c$, the characteristic scale of the 2-D galaxy distribution, and dispersions around the means of 128 $\pm$ 88, 189 $\pm$ 116, 292 $\pm$ 191 and 1582 $\pm$ 771 kpc, for the King, Hubble, NFW and de Vaucouleurs profiles respectively. The outer logarithmic slopes of the distributions were generalized by the usual $\beta$-parameter, which we find to have the following average values: 1.02 $\pm$ 0.08, 1.03 $\pm$ 0.07, 0.61 $\pm$ 0.05 and 7.6 $\pm$ 0.5, for the King, Hubble, NFW and de Vaucouleurs profiles respectively, which are consistent. The average background density at the limit of the COSMOS catalogue is about 4 10$^{-5}$ galaxies arcsec$^{-2}$. In order to investigate whether the galaxy distributions in our clusters preponderantly have cores or cusps, we have determined the likelihood ratio for the King and NFW profiles. Using all galaxies down to the COSMOS magnitude limit of about m$_{b_j}$ $\approx$ 19.5, we find that in general the King profile is more likely to be a good representation of the data than the NFW profile. However, for the individual clusters this preference for the King profile is generally not statistically significant. If we restrict the analysis to the central regions, the significance of the preference for the King-profile fits increases, even though the number of galaxies decreases. We have increased the statistical weight for the likelihood analysis by combining the galaxy distributions in a subset of 29 of the 62 clusters, which show a regular galaxy distribution. We take special care to avoid the creation of an artificial cusp (by taking the ellipticities into account), and to avoid the destruction of a real cusp by summing distributions with different scale lengths. We have also checked that it is unlikely that the uncertainty in the centre positions has erased a cusp. For the test we summed without scaling projected distances, after scaling with r$_c$, as well as with r$_{200}$. In all three cases we find that the King profile provides a better fit to the data than the NFW profile, at confidence levels of more than 95 \%. Interestingly, this preference is not shared by the brighter galaxies. Finally, we have used the outer profile slope (i.e. the result that $\beta$ is very close to 1.0), in combination with several results from numerical models to conclude that the density parameter $\Omega _0$ is likely to be considerably smaller than unity. In addition, the available models indicate that the Universe probably has an open geometry (i.e. no closure through $\Lambda $ is indicated). This low implied value of $\Omega _0$ is fully consistent with a direct determination based on the $\frac ML$ ratios of our clusters. | 98 | 4 | 9804.172 |
9804 | astro-ph9804085_arXiv.txt | We perform a fractal analysis of the Southern Sky Redshift Survey 2, following the methods prescribed by Pietronero and collaborators, to check their claim that the galaxy distribution is fractal at all scales, and we discuss and explain the reasons of some controversial points, through tests on both galaxy samples and simulations. We confirm that the amplitude of the two--point correlation function does not depend on the sample depth, but increases with luminosity. We find that there is no contradiction between the results of standard and non--standard statistical analysis; moreover, such results are consistent with theoretical predictions derived from standard CDM models of galaxy formation, and with the hypothesis of homogeneity at large scale ($\sim 100$ \h). However, for our SSRS2 volume--limited subsamples we show that the first zero--point of the autocorrelation function $\xi(s)$ increases linearly with the sample depth, and that its value is comparable to the radius of the maximum sphere which can be completely included in the sampled volume; this implies that the true zero--crossing point of $\xi(r)$ has not been reached. We conclude that the apparent fractal behavior is due to a combination of a luminosity--dependent correlation amplitude and the recovering of power at larger scales in deeper samples. | One of the pillars of the standard cosmological models is the large--scale homogeneity of the Universe (e.g. Peebles 1993). The standard statistical methods to analyze the large--scale structure, as the two--point correlation function $\xi(r)$, are described by Peebles (1980); they rely on the definition of a mean galaxy density $\bar{n}$, which is meaningful only if the assumption of large--scale homogeneity is true. However, at small scales the autocorrelation function of the galaxy distribution is positive and can be fitted by a power--law, which is a property of a fractal set (see Peebles 1980). On the basis of the observational evidence that galaxies are clustered in ever increasing systems, from groups and clusters to superclusters, de Vaucouleurs (1970, 1971) presented ``the case for a hierarchical cosmology''\footnote{The idea of a hierarchical Universe has a long history, which dates back to the XVIII$^{th}$ century (see E. Harrison, 1987, {\em Darkness at night, A Riddle of the Universe}, Harvard University Press, Cambridge); in this century, Fournier d'Albe and Charlier were the first to propose a hierarchical (now we would say fractal) model of the Universe (see Mandelbrot 1982).}. Mandelbrot (1982) developed this concept, suggesting that the galaxy distribution in the Universe is fractal, with dimension $D = 1$. In addition, in the last 20 years redshift surveys at increasing depths have revealed ever larger structures and voids (see e.g. Davis et al. 1982; de Lapparent et al. 1986, da Costa et al. 1994, Vettolani et al. 1997, and references therein). Einasto et al. (1986) found evidence that $r_0$ increases with sample volume, but they estimated that it should reach a value $\sim 10$\h~ for a fair sample of the Universe. Pietronero (1987) stressed that, if homogeneity is not reached, the correlation length $r_0$ cannot be taken as a measure of the clustering amplitude, and suggested a slightly but significantly different definition of the autocorrelation function. Adopting this approach, Coleman et al. (1988) reanalyzed the CfA1 redshift survey (Huchra et al. 1983), concluding that the distribution of galaxies was fractal to at least 20 \h. On the other hand, Davis et al. (1988) found that $r_0$ increased as $r_0^{0.5}$, and not linearly with the depth as predicted for a simple fractal (see also Maurogordato et al. 1992). Others advocated the need for a multifractal approach (e.g. Balian \& Schaeffer 1989; Martinez \& Jones 1990; Martinez et al. 1990); this is however another issue and we will not discuss it: see for example the review by Borgani (1995) and references therein. The first redshift surveys sampled relatively small volumes, and the reality and nature of correlation amplitude variations with depth remained an open --and much debated-- question. Could these variations simply reflect fluctuations due to local structures? Were they a consequence of the fractal distribution of galaxies? Or were they an indirect consequence of the dependence of clustering on galaxy luminosity? It should be pointed out that even at a scale $\sim 1000$ \h~ we do not expect a {\em perfect} homogeneity, as COBE has found evidence of anisotropy in the CMB radiation at a level $\delta T / T \sim 10^{-5}$ (Smoot et al. 1992), a value necessary and sufficient, at least in some standard cosmological models, to explain the formation of the observed structures. The existence of very large structures in the Universe implies that the galaxy or cluster autocorrelation function must be positive to a scale corresponding to the size of these structures; but this is not necessarily inconsistent with the measured value of the galaxy autocorrelation length $r_0 \sim 5-6$ \h~ for $L \sim L_*$ galaxies (as claimed by Pietronero and collaborators), as larger structures have also a lower contrast, and the correlation of luminous matter may be significantly amplified relatively to the underlying dark matter correlation function (Kaiser 1984; Bardeen et al. 1986). Therefore, we expect that the galaxy and cluster distribution will not be perfectly homogeneous even at large scales. The claim for a fractal Universe is obviously much more stronger than that: it implies that there is no convergence to homogeneity and that it is not possible to define a mean density $\bar{n}$ of the Universe. In the last years Pietronero and his collaborators applied statistical indicators which do not imply a universal mean density on an ever increasing number of catalogs (see e.g. Di Nella et al. 1996, Sylos Labini, Montuori \& Pietronero 1996, Montuori et al. 1997), claiming evidence for a galaxy fractal distribution at all scales. Their results are impressive, as recently stressed by Coles (1998), but they appear to be at variance with the results of other groups who analyzed the same catalogs with standard indicators. It is clear that the situation is unsatisfactory, despite many articles and even public debates on the subject (see Pietronero et al. 1997; Davis 1997; Guzzo 1997); while the strongest support to large--scale homogeneity comes indirectly from the high level of isotropy of the CMB radiation and from two--dimensional catalogs (Peebles 1996), there is still confusion on the interpretation of the available three--dimensional data, mainly due to the difference in the statistical indicators, which do not allow a direct and quantitative comparison of the results. Therefore we analyzed the Southern Sky Redshift Survey 2 (SSRS2; da Costa et al. 1994) following the approach of Pietronero and collaborators, in order to independently check their claims and to answer to their criticisms (Sylos Labini et al. 1997, hereafter SMP) about the work of Benoist et al. (1996, hereafter Paper I). In section 2 we discuss the apparent dependence of $r_0$ on sample depth; in section 3 we describe the different statistics, and their relations; in section 4 we present and discuss our fractal analysis of the SSRS2 in comparison with the standard analysis; in section 5 we show that theoretical predictions derived from the standard CDM model can reproduce our results, and are therefore consistent with a homogeneous Universe; our conclusions are in section 6. | In this paper, we have carefully examined the claim that the Universe is a fractal, analyzing the SSRS2 and using the same statistical approach of Pietronero and collaborators. Here are our main results: \begin{itemize} \item the SSRS2 has a luminosity dependent correlation amplitude; \item results obtained with the standard and fractal approach give consistent results; \item concerning large--scale homogeneity, observational evidence is --for the moment-- consistent with theoretical predictions derived from standard CDM models of galaxy formation, and does not require a fractal Universe; \item on the other hand, the first zero--point of the correlation function we measure in SSRS2 volume--limited subsamples increases linearly with the sample depth, and has approximately the same value of the radius of the largest sphere which can be included in the sample. This implies that the zero--point of $\xi(s)$ is beyond 40 \h, a lower limit set by very luminous galaxies (see Benoist et al. 1996; Cappi et al. 1998) and consistent with results from clusters (e.g. Cappi \& Maurogordato 1992); on this point, we agree with SMP that measurements of $\xi(s)$ beyond a scale corresponding to the radius of the largest sphere included in the sample are not reliable. \end{itemize} We conclude that the dependence of the galaxy correlation function on the galaxy luminosity and the power still present at large scales ($\ge 40$ \h) can satisfactorily explain those effects interpreted by SMP as evidence of an inhomogeneous, fractal Universe. | 98 | 4 | astro-ph9804085_arXiv.txt |
9804 | astro-ph9804170_arXiv.txt | Observations of clusters and super clusters of galaxies have indicated that the Universe is more dominated by baryons than ever estimated in the homogeneous cosmological model for primordial nucleosynthesis. Recent detections of possibly low deuterium abundance in Lyman-$\alpha$ clouds along the line of sight to high red-shift quasars have raised another potential difficulty that \he4 is overproduced in any cosmological models which satisfy the low deuterium abundance constraint. We show that the inhomogeneous cosmological model with degenerate electron-neutrino can resolve these two difficulties. | One of the cosmological impacts of primordial nucleosynthesis is on the universal baryon mass density $\rho_b $. It is of significance to answer the question how much fraction of universal mass is made of ordinary matter baryons. Homogeneous Big-Bang model for primordial nucleosynthesis~\cite{wagoner67,copi95}, assuming the standard model for light neutrino families, predicts small $\Omega_b$, $0.03 \le \Omega_b h_{50}^2 \le 0.06$, where $\Omega_b = \rho_b / \rho_c$, $\rho_c$ is the critical density which marginally closes the Universe, and $h_{50}$ is the Hubble constant $H_0$ divided by $50\,km/s/Mpc$. However, recent observations of rich clusters and super clusters of galaxies have indicated much larger baryon fraction, 0.1 $\le \Omega_b h_{50}^{3/2} \le$ 0.3~\cite{white93,bahcall95}. Inhomogeneous Big-Bang model~\cite{appligate85}--\cite{mathews96}, which allows inhomogeneous baryon density distribution due to various physical processes in the early Universe, has been proposed in order to resolve this discrepancy. In this model difference in diffusion effects between neutrons and charged nuclei plays an important role in fluctuating density distribution to suppress overproduction of \he4, and resultant $\Omega_b$ is relaxed to $\Omega_b \sim 0.1$~\cite{orito97}. However, another potential difficulty has been imposed by recent observations of deuterium absorption line in Lyman-$\alpha$ clouds along the line of sight to high red-shift quasars~\cite{rugers96,tytler96}. Several detections~\cite{tytler96} among them provide too small deuterium abundance to accept concordant $\Omega_b$ which satisfies the abundance constraints on the other light elements \he3, \he4 and \li7. The observed deuterium abundance still scatters largely by one order of magnitude depending on different Lyman-$\alpha$ systems, and there are still many error sources unclear in the analysis of abundance determination. However, if these detections are real, the abundance found there is presumed to constrain most strongly the primordial abundance because these clouds are the primitive gas which still resides in the epoch of galaxy formation and has not been processed very much in its evolutionary history. It is the purpose of this paper to propose that the inhomogeneous cosmological model with degenerate electron-neutrino can resolve these two difficulties simultaneously within the framework of the standard model for neutrino. In the next section we first discuss neutrino properties in the early Universe which can affect strongly the primordial nucleosynthesis. We then present the results of primordial nucleosynthesis calculated in both homogeneous and inhomogeneous cosmological models in sect. 3, and the $\Omega_b$ problem and the problem of overproduction of \he4 are discussed in details. Finally, in sect. 4, we summarize this paper. | We studied the effects of lepton asymmetry of partially degenerate electron-neutrino on the primordial nucleosynthesis. Homogeneous Big-Bang model with neutrino degeneracy parameter $\xi_{\nu_e}$ = 0.05 can recover the concordance between \he4 and low deuterium abundance which was found in some Lyman-$\alpha$ clouds along the line of sight to high red-shift quasars, but the resultant $\Omega_b$ is less than detected in rich clusters 0.1 $\le \Omega_b h_{50}^{3/2} \le$ 0.3. It was found that the inhomogeneous Big-Bang model with the same degeneracy parameter can predict $\Omega_b$ as large as 0.22, which is in reasonable agreement with observation. This degeneracy parameter corresponds to a small chemical potential of electron-neutrino of order 10$^{-5}$ eV. It is desirable to detect the asymmetry of background neutrinos. | 98 | 4 | astro-ph9804170_arXiv.txt |
9804 | astro-ph9804036_arXiv.txt | We present a detailed analysis of a high resolution spectrum of the damped Ly$\alpha$ system at $z_{\rm abs}$~=~2.8112 toward PKS~0528-250. The absorption redshift is slightly larger than the emission redshift of the quasar. We estimate the column density of H$_2$ molecules $N$(H$_2$)~$\sim$~6$\times$10$^{16}$~cm$^{-2}$ and the fractional abundance of H$_2$, $f$~=~5.4$\times$10$^{-5}$. The excitation temperature derived for different transitions suggests that the kinetic temperature of the cloud is $\sim$200~K and the density $n$~$\sim$~1000~cm$^{-3}$. The cloud therefore has a dimension of $\sim$1~pc along the line of sight. Since it obscures the broad-line emission region, its transverse dimension should be larger than 10~pc.\par We obtain upper limits on the column densities of C~{\sc i} ($<$~10$^{12.7}$~cm$^{-2}$) and CO ($<$~10$^{13.2}$~cm$^{-2}$; $N$(CO)/$N$(H~{\sc i})~$<$~7$\times$10$^{-9}$). We suggest that the ratio $N$(H$_2$)/$N$(C~{\sc i}) is a useful indicator of the physical conditions in the absorber. Simple photo-ionization models assuming solar relative abundances show that radiation fields with spectra similar to typical AGNs or starbursts are unable to reproduce all the constraints and in particular the surprisingly small $N$(C~{\sc i})/$N$(H$_2$) and $N$(Mg~{\sc i})/$N$(H$_2$) ratios. In view of the models we explored, the most likely ionizing spectrum is a composite of a UV-"big bump" possibly produced by a local starburst and a power-law spectrum from the QSO that provides the X-rays. Dust is needed to explain the production of molecules in the cloud. The amount of dust is broadly consistent with the [Cr/Zn] abundance determination. | \label{intr} QSO absorption line systems probe the baryonic matter over most of the history of the Universe (0~$<$~$z$~$\la$~5). The so-called damped Ly$\alpha$ (hereafter DLA) systems are characterized by a very large H~{\sc i} column density ($N$(H~{\sc i})~$\ga$~2$\times$10$^{20}$ ~cm$^{-2}$), similar to the one usually seen through local spiral disks. The case for these systems to be produced by proto-galactic disks is supported by the fact that the cosmological density of gas associated with these systems is of the same order of magnitude as the cosmological density of stars at present epochs (Wolfe 1996). Moreover the presence of heavy elements ($Z \sim 0.1 ~ Z_\odot$) and the redshift evolution of metallicity suggest the ongoing star formation activities in these systems (Pettini et al. 1997), while strong metal line systems have been demonstrated to be associated with galaxies at low and intermediate $z$ (e.g. Bergeron \& Boiss\'e 1991). It has also been shown that the profiles of the lines arising in the neutral gas show evidence for rotation (Wolfe 1996, Prochaska \& Wolfe 1997). Whether these arguments are enough to demonstrate that DLA systems arise in large disks is a matter of debate however. Indeed simulations have shown that the progenitors of present day disks of galaxies could look like an aggregate of well separated dense clumps at high redshift. The kinematics could be explained by relative motions of the clumps with very little rotation (Haehnelt et al. 1997, Ledoux et al. 1998). Moreover, using {\sl HST} high spatial resolution images of the field of seven quasars whose spectra contain DLA lines at intermediate redshifts (0.4~$\la$~$z$~$\la$~1), Le~Brun et al. (1996) show that, in all cases, at least one galaxy candidate is present within 4~arcsec from the quasar. There is no dominant morphological type in their sample: three candidates are spiral galaxies, three are compact objects and two are amorphous low surface brightness galaxies. Therefore, although the nature of the DLA systems is unclear they trace the densest regions of the Universe where star formation occurs.\par It is thus surprising that despite intensive searches, the amount of H$_2$ molecules seems quite low in DLA systems in contrast to what is observed in our own galaxy. Two detections of H$_2$ molecules in high redshift DLA systems have been reported. Recently Ge \& Bechtold (1997) have found strong absorptions in the $z_{\rm abs}$~=~1.9731 DLA system toward Q~0013--004. They derive $N$(H$_2$)~=~6.9$\times$ 10$^{19}$~cm$^{-2}$, $b$~=~15~km~s$^{-1}$, $T_{\rm ex}$~$\sim$~70~K and $n$(H)~$\sim$~300~cm$^{-3}$ for a total hydrogen column density $N$(H)~=~6.4$\times$10$^{20}$~cm$^{-2}$. This system has by far the largest H$_2$ abundance $f$~=~2$N$(H$_2$)/[2$N$(H$_2$)~+~$N$(H~{\sc i})] $\sim$~0.22$\pm$0.05 observed in high z DLA systems. However the exact number should be confirmed using a higher resolution data. Other searches have led to much smaller values or upper limits ($f$~$<$~10$^{-6}$, Black et al. 1987, Chaffee et al. 1988, Levshakov et al. 1992). Levshakov \& Varshalovich (1985) suggested that H$_2$ molecules could be present in the $z_{\rm abs}$~=~2.8112 system toward PKS~0528--250. This claim has been confirmed by Foltz et al. (1988) using a 1~\AA~ resolution spectrum. The latter authors derive $N$(H$_2$)~=~10$^{18}$~cm$^{-2}$, $b$~=~5~km~s$^{-1}$, $T_{\rm ex}$~=~100~K and log~$N$(H~{\sc i})~=~21.1$\pm$0.3. By fitting the damped absorption together with the Ly$\alpha$ emission from the quasar, M\o ller \& Warren (1993) find log~$N$(H~{\sc i})~=~21.35. Three Ly$\alpha$ emission-line objects have been detected within 100$h^{-1}$~kpc from the quasar by M\o ller \& Warren (1993) and confirmed by Warren \& M\o ller (1996) to have redshifts within 200~km~s$^{-1}$ from the redshift of the DLA system ($z_{\rm abs}$~=~2.8112 as measured on the Ni~{\sc ii} lines by Meyer \& York 1987). The widths of the Ly$\alpha$ emission lines are very large ($>$~600~km~s$^{-1}$) and continuum emission could be present (Warren \& M\o ller 1996); this suggests that the gas is not predominantly ionized by the quasar and that star-formation may occur in the clouds, a conclusion reached as well by Ge et al. (1997). The proximity of the quasar makes the case difficult however and careful analysis is needed.\par In this paper we use much higher spectral resolution data to reanalyze the molecular lines in this system. We present the observations in Section~2, the results in Section~3 and discuss the low $N$(C~{\sc i})/$N$(H$_2$) ratio inferred in Section~4. | \label{s4} By fitting the different H$_2$ transitions that we detect in a high resolution spectrum of PKS~0528--250, we derive log~$N$(H$_2$)~$\sim$~16.78 and $T_{\rm ex}$~$\sim$~200~K that is most certainly the true kinetic temperature of the gas. For this temperature, the relative populations of rotational levels 0 to 4 indicate that the density is of the order of 1000~cm$^{-3}$. Therefore the dimension of the neutral cloud along the line of sight is less than 1~pc. It must be noticed that the damped absorber must cover the broad line region. Indeed there is no residual flux in the bottom of the Ly$\alpha$ absorption line which completely absorbs the Ly$\alpha$ emission from the quasar over more than 5000~km~s$^{-1}$ (M\o ller \& Warren 1993). Since the dimension of the BLR in QSOs can be approximated as $R$~$\sim$~0.3~$L_{46}^{0.5}$ where $L_{46}$ is the bolometric luminosity in units of 10$^{46}$~erg~s$^{-1}$ (Collin, private communication). The radius of the BLR in PKS~0528--250 is thus of the order of 10~pc. The transverse dimension of the damped cloud must be thus larger than 10~pc and the cloud must be quite flat. We can derive an upper limit on the transverse dimension of the cloud by interpreting the non-dection of redshifted 21 cm absorption (Carilli et al. 1996) as an effect of partial covering factor of the continuum radio emission by the cloud. Indeed the size of the radio source is 1~arcsec or 5$h^{-1}_{75}$~kpc. If the spin temperature is equal to the kinetic temperature we derived in sect. 3.2, this implies that the covering factor should be less than 0.3 and thus the radius of the cloud along the transverse direction is less than 1~kpc. We determine upper limits for the C~{\sc i}, Mg~{\sc i} and CO column densities (log~$N$(C~{\sc i})~$<$~12.7, log~$N$(Mg~{\sc i})~$<$~12.8 and log~$N$(CO)~$<$~13.2). Based on simple photoionization models we conclude that (i) no simple model can reproduce at the same time the low $N$(C~{\sc i})/$N$(H~{\sc i}), $N$(Mg~{\sc i})/$N$(H~{\sc i}) ratios and the presence of molecules at the level observed; (ii) steep spectra are required to reproduce the low $N$(C~{\sc i})/$N$(H~{\sc i}) and $N$(Mg~{\sc i})/$N$(H~{\sc i}) ratios but they predict temperatures smaller than 100~K in conflict with the excitation temperature derived from the H$_2$ transitions; (iii) the only way to keep the temperature as high as 200~K is to allow for some X-ray flux heating the gas; (iv) in the framework of these models, dust is needed to produce the observed amount of molecules; (v) the gas phase abundances of carbon and especially magnesium should be smaller than 0.1 of solar. In view of the models we explored, the most likely ionizing spectrum is a composite of a UV-"big bump" possibly produced by a local starburst and a power-law spectrum from the QSO that provides the X-rays. | 98 | 4 | astro-ph9804036_arXiv.txt |
9804 | astro-ph9804292_arXiv.txt | The effects that large scale fluctuations had on small scale isothermal modes at the epoch of recombination are analysed. We find that: (a)~Albeit the fact that primordial fluctuations were at this epoch still well in the linear regime, a significant nonlinear radiation hydrodynamic interaction could have taken place. (b)~Short wavelength isothermal fluctuations are unstable. Their growth rate is exponential in the amplitude of the large scale fluctuations and is therefore very sensitive to the initial conditions. (c)~The observed CMBR fluctuations are of order the limit above which the effect should be significant. Thus, according to their exact value, the effect may be negligible or lead to structure formation out of isothermal fluctuations within the period of recombination. (d)~If the cosmological parameters are within the prescribed regime, the effect should be detectable through induced deviations in the Planck spectrum. (e)~The sensitivity of the effect to the initial conditions provides a tool to set limits on various cosmological parameters with emphasis on the type and amplitude of the primordial fluctuation spectrum. (f)~Under proper conditions, the effect may be responsible for the formation of sub-globular cluster sized objects at particularly high red shifts. (g)~Under certain circumstances, it can also affect horizon sized large scale structure. | The problem of structure formation in the Universe has probably been one of the foremost and most studied question in cosmology. Perhaps the greatest achievement of cosmology was the prediction and following discovery of the Cosmic Microwave Background Radiation -- the CMBR (Penzias \& Wilson, 1965) and its anisotropy (Smoot et al. 1992). Not only did it provide support for the Big Bang theory, it proved that structure is a result of small fluctuations growing into large inhomogeneities and not vice versa. The fact that the observed fluctuations in the CMBR are small ($\Delta T/T \sim 10^{-5}$), naively implies that one can treat the fluctuation within the linear approximation when modeling the evolution of structure before the time of last photon scattering, i.e., one can assume that fluctuation modes of different wavelengths are completely decoupled from each other before radiation-matter decoupling. The first to develop a linear theory for the perturbed Fridmann - Robertson - Walker metric was Lifshitz (1946), and it can be found in various text books such as Weinberg (1972), Peebles (1980) and Kolb \& Turner (1990). The theory has since then been applied to describe cosmologies with various components and various initial parameters. Baryonic matter, radiation and other massless particles, cold (massive) or hot (light) dark matter are some of the ingredients that enter the primordial soup. While other parameters such as the Hubble constant $H_0$, the vacuum energy through the cosmological constant $\Lambda$, and the initial spectrum affect the evolution. A review of various current cosmological models, their evolution and implication can for example be found in White et al. (1994) with emphasis on the microwave background radiation and in Primack (1997). Less current reviews that cover the basic principles of structure formation are found in the aforementioned textbooks. One of the major parameters that affect the qualitative behaviour of our universe and is highly relevant to the present paper, is the type and form of the primordial spectrum of fluctuations. The fluctuations are generally classified into curvature (or adiabatic) and isocurvature (or isothermal) fluctuations. The first arise naturally in inflationary scenarios (e.g., Liddle and Lyth, 1993, and ref. therein). They are fluctuations of both the baryonic fluid and the radiation and they propagate at the adiabatic speed of sound that is equal to $c/\sqrt{3}$ if the radiation energy density dominates. These waves are found to decay on scales smaller than the Silk scale, namely, scales comparable to galaxy sized objects (Silk, 1967). Therefore, a top-down structure formation is a natural consequence of adiabatic perturbations. Small scale objects can form after recombination (and initiate a bottom-up scenario) only if one adds the undamped perturbation of a cold component as is the case in the standard CDM model (Peebles 1982, Blumenthal et al. 1984 and Davis et al. 1985) and its variants. Isocurvature (or isothermal) fluctuations on the other hand, are a natural consequence of topological defects formed in phase transitions (e.g, strings, monopoles and textures) or if more than one field contributes significantly to the energy density during inflation. They correspond to fluctuations that alter the entropy density but not the energy density. Unlike the first type of fluctuations, these do not suffer from Silk damping. Consequently, objects with a mass as small the post-recombination Jeans scale, namely the size of globular clusters, can form after decoupling. It is interesting and important to note that even if the primordial spectrum was purely adiabatic, isothermal perturbations of a given wavelength are formed as the second order effect of Purcell clustering from adiabatic waves of similar wavelengths (Press \& Vishniac, 1979). Through this effect, noninteracting particles (baryons) that are viscously coupled to a stochastically oscillating background gas (radiation fluid) undergo secular clustering. Originally, it was thought that the effect can produce a bottom-up scenario even from a pure baryon model with an adiabatic spectrum. However, the typical post-recombination Jeans scale amplitude of the isothermal waves will only be $\rho_{iso}/\rho\sim (10^{-2} - 1) \times \left( \delta \rho_{ad} / \rho\right)^2$, with $\delta \rho_{ad} /\rho$ the typical adiabatic fluctuations at horizon crossing. Namely, if the primordial spectrum is flat and adiabatic, the typical isothermal fluctuation at recombination is roughly $\delta \rho_{iso} /\rho \sim 10^{-10}-10^{-8}$. For a flat isothermal primordial spectrum, one should expect typical amplitudes of $\rho_{iso}/\rho \sim 10^{-5} - 10^{-4}$. The smallness of nonlinear effects such as the Purcell clustering (or shock waves which are of an even higher order) led to the consensus that the evolution of the fluctuations can be treated linearly and modes of different wavelengths are decoupled from each other. This delays the nonlinear treatment to the time when radiation does not play any dynamic role anymore. At face value, it certainly appears to be the case as $\delta\rho /\rho$, $v/c$, and $\delta T/T$ are all much smaller than unity. One should nevertheless be extremely careful when assuming linearity, especially in view of the fact that not all of the dimensionless parameters are actually smaller than unity. One of the dimensionless numbers that appears in the solution of the nonlinear fluctuations' equations of motion and that is not small at all is the root of the radiation to gas pressure ratio. Just before recombination one finds that $\sqrt{P_{rad}/P_{gas}}\approx{10}^{5}$! Consequently, we expect that the interaction between the radiation and matter at this period will have profound impact on the evolution of the fluctuations. In this paper we examine the linear hypothesis and its validity by adding the force large scale perturbations exert on short wavelength waves. We begin in \S2 by overviewing the problem of solving the nonlinear equations of motion and estimating the effect with a very simple analysis. We proceed in \S3 to write the Newtonian equations describing the evolution of short wavelength isothermal modes. In \S4 we analyse the simplified solution, while in \S5 we proceed to estimate the effect in a few cosmological scenarios. In \S6 \& \S7, we study the possible ramifications to structure formation and study the possibility of measuring and using this effect in the study of cosmological parameters. In \S8 we show that the effect can influence large scale structure as well. | We have analysed the nonlinear growth of small scale isothermal fluctuation through the induced effect that large-scale perturbations have. The equations are solved through the separation of the large scale inhomogeneities into smaller regions where they can be considered uniform, and the shear $\Delta v_0$ exerted by the large scales is constant in space. This can be performed because the waves that interest us are waves with very small wavelengths and a very small propagation speed, such that over the integration time, the waves practically remain in the same region of space. Small scale isothermal waves are found to be unstable if the large scale shear velocity between the ``photon fluid'' and the baryon fluid is larger than roughly twice the isothermal speed of sound at the time of recombination. A period when the logarithmic derivative of the opacity with respect to the density does not vanish but is $-1/2$ instead. When the shear is smaller than the critical value, the solution found is qualitatively the same as the solution for the completely decoupled equations. The shear in this case may induce some quantitative corrections but the behaviour of the system does not change. The behaviour of the system is however radically different if the shear is larger than a critical value. For wavelength shorter than the roughly the Jeans scale, the growth rate increases with $k$ and it saturates when $\tau_{e\gamma}=\tau_o$, i.e., at a scale which is a few times smaller than the Jeans scale. The maximal growth rate attained is: \begin{equation} r_{\infty} = {\alpha \over 2} {\Delta v_0 \over v_s} {1\over \tau_e{\gamma}}. \end{equation} If $r_{\infty}$ is much larger than all other rates in the system, it can be easily integrated over the duration of decoupling $\Delta t$. The growth factor obtained is then: \begin{equation} {\delta_1\over\delta_0} \approx \exp (G_{\infty}) = \exp \left(\left\langle {\alpha \over 2} {\Delta v_0 \over v_s} {1\over \tau_e{\gamma}}\right\rangle_t \Delta t\right), \end{equation} with the brackets denoting a temporal average. The isothermal waves are unstable because their sound speed is very small. Consequently, the acoustic energy in a wave with a given amplitude is very small and the amount of work needed to increase its amplitude isn't large, in fact, it is so small that nonlinear effects start to take place already at $\delta\rho/\rho\sim 10^{-5}~{\rm to}~10^{-4}$! Even more surprising is the fact the the typical shears found in typical cosmological scenarios is {\em exactly} at the verge of having nonlinear effects take place! If the amplitude of the large scale fluctuations would have been an order of magnitude larger, then any isothermal perturbations, however small, could have been amplified right into the nonlinear regime. On the other hand, if the typical amplitude would have been a few times smaller, the effect would have been completely meaningless. For typical cosmological scenarios normalized to COBE, the growth parameter is roughly $G_{rms} \sim 1 - 5$, e.g., for an $\Omega=1,~\Omega_b=0.05,~h=.66$ CDM model one finds $G_{rms} \approx 5$, if the spectral index is tilted to $n=0.9$, one finds $G_{rms}\approx 4$, if a hot component is added at a 20\% level to the untilted model, it falls to $G_{rms} \sim 1$. The fate of the regions that do reach nonlinearity is still an open question. Do these regions collapse and form black holes? Do they form an early generation of massive stars? Do they release enough energy to the environment and affect it, or perhaps, even form structure on much larger scales? Although some speculations of what might occur do exist in the context of early small scale structure formation, it is not entirely clear. Moreover, nonlinear hydrodynamic simulations is probably unavoidable if we are to really solve the problem. The only relatively certain consequence is that the dissipation taking place after recombination heats the matter and it is likely to ionise it, raise its temperature to the ionisation temperature and leave an imprint on the CMBR in the form of a deviation from a Planckian spectrum. In scenarios where $G_{rms}\gtrsim 5$, isothermal perturbations with an amplitude of as low as $\delta\rho_{iso}/\rho\sim10^{-6}$ will be in fact theoretically detectable in the future. An amplitude of $\delta\rho_{iso}/\rho\sim10^{-4}$ will leave 1\% of the background sky with a detectable deviation. This fraction can change considerably if the density perturbations are not Gaussian. The detection of deviations from a Planckian spectrum or the placing of an upper limit for them can result with interesting implications. Any deviations found will first of all directly prove the existence of primordial isothermal fluctuations. Second, such a detection will place stringent limits on cosmological parameters, as only those scenarios that produce a large enough $G_{rms}$ are capable of producing nonlinear structure, and because the amount of nonlinearity is extremely sensitive to $G_{rms}$, it is also sensitive to the exact cosmological parameters. Third, if the cosmological parameters are known with an large accuracy (for example, through the fitting of the CMBR spectrum) then the primordial isothermal perturbation spectrum at very large $k$'s can be estimated. Even if no detection of a deviation from a Planckian spectrum is found, one can place limits on cosmological parameters. Moreover, if in the future it will be found that the cosmological parameters actually correspond to a high $G_{rms}$ model, one will be able through the lack of detection of a $y$-parameter to place extremely powerful limits on the amplitude of the isothermal component of the primordial spectrum. Through the Press \& Vishniac effect, limits can also be placed on the adiabatic spectrum. Another interesting implication is the possibility of transferring enough energy from the large to the small scales and considerably change the amplitude of the large-scales. This possibility relies on whether a significant volume fraction can be amplified out of the linear regime and on the maximum dissipation rate of nonlinear isothermal waves. Under favourable circumstances, the large scale amplitude can be significantly reduced such that $G_{rms}$ calculated from the observed large scale fluctuations would only be of order unity. The fact that the observed fluctuations in the CMBR correspond to values of this order raises a very interesting question. Why do the values of $G_{rms}$ corresponding to the observed fluctuations happen to fall in a small region around unity? Is it because the universe was created with a fluctuation spectrum corresponding to this region, or, is it because $G_{rms}$ that evolved from the initial cosmological parameters was actually larger but it was naturally reduced to the oberved value? One can summarise the several plausible scenarios according to both the value of $G_{rms}$ predicted by the cosmological model and the amplitude $\delta = \delta \rho_{iso}/\rho \approx 10^{-10} - 10^{-5} $ of the sub Jeans scale isothermal fluctuations before the advent of recombination, these possibilities are: \begin{enumerate} \item If $G_{rms}$ is of order unity or less, the typical growth of isothermal waves is at most a few $e$-folds. The effect will be insignificant as the predicted micro degree size fluctuations are neither fluctuations of the temperature nor on a scale measurable in the near future. Moreover, there are no implications at all to structure formation. \item For a value of $G_{rms} \gtrsim - (\ln \delta)/3 \approx 4 - 8$, the effect will be measurable as patches in the CMBR with a distorted Planckian spectrum. As long as $G_{rms} \lesssim - \ln \delta \approx 10-20$, only a small fraction of the universe would have reached nonlinearity and formed small scale structure by the end of recombination. Under certain circumstances however, it can lead to large scale structure formation through explosive amplification. \item For a value of $G_{rms}$ such that $ G_{rms} \gtrsim - \ln \delta \approx 10-20$, a large fraction of the universe reaches nonlinearity by the end of recombination and small scale structure is subsequently formed in most of the universe soon after recombination. The large scale spectrum is damped by transferring a significant amount of energy to the small scales, thus, the $G_{rms}$ measured from the damped spectrum is smaller than the $G_{rms}$ calculated when neglecting this process. In some cases, $G_{rms}$ will be reduced to a values of order unity and it will mimic values on the boundary between the first and second scenarios. Note that it will not affect fluctuations in components that decouple from the matter-radiation fluid before recombination (e.g. dark matter fluctuation). \end{enumerate} The analysis presented here is the first step in the investigation of the radiation-matter interaction instability. More accurate integration is needed to improve the actual transfer function for isothermal waves. The analysis here was restricted to the evaluation of $G_{rms}$ and it does not include the actual rates which are $k$ dependent. This approximation overestimates the rate of growth of finite sized $k$'s. In the present paper we have simulated the freezing of the instability due the freezing of recombination by stopping the integration abruptly at a given $z$. However, the equilibrium conditions and therefore the switching off of the effect depend on the isothermal wavelength as well. By assuming this assumption, we have actually underestimated the contribution from waves of order the Jeans size. Many of the possible implications depend on the quantitative behaviour of of the nonlinearities once they are reached. A numerical hydrodynamic study of large amplitude waves will certainly help us understand of the fate of the nonlinear objects and the possible implications they have on large scale structure as well. | 98 | 4 | astro-ph9804292_arXiv.txt |
9804 | astro-ph9804017_arXiv.txt | I summarize some recent models and ideas for the formation of axisymmetrical structures of planetary nebulae and the three rings of SN 1987A, as follows. (a) I review the general role of binary companions, including brown dwarfs and planets. (b) I propose a mechanism for axisymmetrical mass loss on the AGB that may account for the axially symmetric structures of elliptical planetary nebulae and that operates for slowly rotating AGB stars, $10^{-4} \Omega _{\rm Kep} \lae \Omega \lae 10^{-2} \Omega_{\rm Kep}$, where $\Omega_{\rm Kep}$ is the equatorial Keplerian angular velocity. (c) I propose a model for the formation of the two outer rings of SN 1987A, which is based on the numerical simulation of Soker (1989), and discuss a mechanism for their displacement from the exploding star. | Scanning through recent images of SN 1987A (e.g. Burrows {\it et al.} 1995) and through catalogs of planetary nebulae (PNs; e.g., Acker {\it et al.} 1992; Schwarz, Corradi, \& Melnick 1992; Manchado {\it et al.} 1996) we find that the circumstellar media of many stars at their final nuclear burning phase have axisymmetrical, rather than spherical, structures. Axisymmetrical PNs which have two lobes with a morphological ``waist'' between them are termed ``bipolar PNs'' (also ``butterfly'' or ``bilobal''), while PNs which have a more elliptical than bilobal structure are termed elliptical PNs (Schwarz, Corradi, \& Stanghellini 1993). The axisymmetrical structures of most PNs led to a debate on whether elliptical PNs can be formed through single-stellar evolution, or whether a binary companion is necessary (e.g., Fabian \& Hansen 1979; Livio 1982, 1998; Livio, Salzman, \& Shaviv 1979; Webbink 1979; Morris 1981; Zuckerman \& Gatley 1988; Pascoli 1992; Iben \& Livio 1993; Soker 1997, 1998a; Balick {\it et al.} 1994; Pottasch 1995; % Pollacco \& Bell 1997; Corradi {\it et al.} 1996; Kastner {\it et al.} 1996). In the last decade this debate was extended to the formation of the nonspherical explosion and three rings of SN 1987A. In many PNs, as well as in the three rings of SN 1987A, there are displacements of the nebulae from the central stars, which hint at the interaction of the progenitors with wide binary companions, with close binaries having eccentric orbits, or with the ISM ($\S 5$). In a recent paper (Soker 1997) I suggest that four main evolutionary routes determine the degree of asymmetry of the axially symmetric structures of PNs. I then classify 458 PNs according to the process which caused their progenitors to blow axisymmetrical winds. The classification is based primarily on the morphologies of the different PNs, assuming that binary companions, stellar or substellar, are necessary for axisymmetrical mass loss on the AGB. The four evolutionary classes, according to the binary-model hypothesis, are: \newline (a) Progenitors of planetary nebula which did not interact with any companion, and therefore they rotate extremely slowly when reaching the AGB. These amount to $\sim 10 \%$ of all planetary nebulae. \newline (b) Progenitors which interact with stellar companions which avoided a common envelope, $11 ^{+ 2}_{-3} \%$ of all nebulae. These form bipolar PNs, as is the case in symbiotic nebulae (Morris 1990; Schwarz \& Corradi 1992; Soker 1998a). % \newline (c) Progenitors which interact with stellar companions via a common envelope phase, $23^{+11}_{-5} \%$ of all nebulae. These form extremely asymmetrical structures, i.e., tori, elongated elliptical PNs, and in some cases bipolar PNs. \newline (d) Progenitors which interact with {\it substellar} (i.e., planets and brown dwarfs) companions via a common envelope phase, $56^{+5}_{-8} \%$ of all nebulae. These form elliptical PNs with relatively small deviation from sphericity. \newline These numbers are compatible with other studies (e.g., Yungelson, Tutukov, \& Livio 1993; Han, Podsiadlowski, \& Eggleton 1995). In $\S 2$ I discuss the problem of angular momentum of AGB stars, which suggests that to account for the $\sim 60 \%$ elliptical PNs either there are many planetary systems ($\S 3$; Soker 1996; 1997) or there is a mechanism to induce axisymmetrical mass loss from very slowly rotating AGB stars. Such a model for singly evolved very slowly rotating AGB stars is the mechanism of mode-switch to nonradial oscillations, proposed by Soker \& Harpaz (1992). In $\S 4$ I propose yet another model (Soker 1998c) which may operate in singly evolved AGB stars. This model is based on both magnetic activity and radiation pressure on dust. In $\S 5$ I propose a model for the two outer rings of SN 1987A. | 98 | 4 | astro-ph9804017_arXiv.txt |
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9804 | astro-ph9804221_arXiv.txt | We investigate a one-zone chemophotometric evolution model of disk-disk galaxy mergers in order to clarify whether or not galaxy mergers with the widely spread merging epoch can reproduce reasonably well the observed small scatter of the color-magnitude ($C-M$) relation in cluster ellipticals at low and intermediate redshift ($z<1$). We consider that merger progenitor disks begin to consume interstellar gas at moderate rate from $z \sim 5$ and then merge to form an elliptical with the secondary starburst at $z=z_{\rm merge}$. We find that even if the epoch of galaxy merging is rather extended ($0.3<z_{\rm merge}<3.0$), the dispersion in the rest-frame $U-V$ color among galaxy mergers is well within the observed one ($\sim$ 0.05 mag at $z$ = 0). We also find that the $z_{\rm merge}$ is required to be within a certain range to keep the observed $C-M$ relation tight at a given $z$. For example, the required range of $z_{\rm merge}$ in galaxy mergers between Sa disks is $1.3<z_{\rm merge}<3.0$ for cluster ellipticals at $z$ = 0.895, $0.9<z_{\rm merge}<3.0$ for $z$ = 0.55, and $0.3<z_{\rm merge}<3.0$ for $z$ = 0. The main reason for the derived small scatter is that younger stellar populations, which are formed during the secondary starburst of galaxy mergers, are formed preferentially from more metal-enriched interstellar gas. This result reinforces the Worthey's suggestion (Worthey et al. 1996) that the age-metallicity conspiracy, which means that younger stellar populations are preferentially more metal-enriched, can operate to keep the tight $C-M$ relation. These numerical results imply that the observed small scatter in the $C-M$ relation at low and intermediate redshift ($z < 1$) {\it does not necessarily} require the coevality of elliptical galaxies in clusters or their formation at high $z$, which has been conventionally believed in the classical passive evolution picture. | Redshift evolution of fundamental physical relations in elliptical galaxies is generally considered to give strong constraints on the formation and evolution of elliptical galaxies. For example, evolution of the color-magnitude ($C-M$) relation with redshift ($z$) suggests that elliptical galaxies are old, coeval, and homogeneous systems passively evolving after the single initial burst of star formation associated with dissipative galaxy formation at $z > 2.0$ (Arag\'on-Salamanca et al. 1993; Ellis et al. 1997; Stanford, Eisenhardt, \& Dickinson 1998). This classical picture of coeval elliptical galaxy formation also appears to be supported by small redshift evolution of both the mass-to-light-ratio (van Dokkum \& Franx 1996) and the $\rm Mg_{2} - \sigma$ relation (Ziegler \& Bender 1997). The considerably tight $C-M$ relation (Bower, Lucey, \& Ellis 1992) and the Fundamental Plane (e.g., Djorgovski \& Davis 1987) at the present epoch, and redshift evolution of the slope and the zero-point of the $C-M$ relation (Kodama \& Arimoto 1997; Gladders et al. 1998; Kodama et al. 1998) furthermore seem to support the coevality of elliptical galaxy formation. An increasing number of recent observational results, however, shed a strong doubt on this long-standing view of elliptical galaxy formation, and suggest that there is great variety of star formation history among elliptical galaxies, such as the epoch of major star formation, the duration and efficiency of star formation (Worthey, Faber, \& Gonzalez 1992; Faber et al. 1995; Worthey, Trager, \& Faber 1996). In particular, Faber et al. (1995) suggested that the `apparent age spread', which is inferred from the combination of line index analysis of elliptical galaxies, amounts to $\sim$ 10 Gyr. Schweizer \& Seitzer (1992) found that in merger remnants with morphologically fine structures, the last merging epoch, which corresponds to elliptical galaxies formation, ranges from 4.6 Gyr to 8.0 Gyr ago. These observed spread in `apparent mean age' seem to be inconsistent with the aforementioned coevality of elliptical galaxy formation expected mainly from the redshift evolution of the $C-M$ relation. The purpose of this paper is to give a plausible answer to the above apparent inconsistency in the epoch of elliptical galaxy formation. We adopt the merger scenario of elliptical galaxy formation (e.g., Toomre \& Toomre 1972) and thereby investigate to what degree the difference in the epoch of major galaxy merging (i.e., the epoch of elliptical galaxy formation) can be allowed to preserve the observed small scatter of the $C-M$ relation of cluster ellipticals ($\sim 0.05$ mag) at $z$ = 0 (Bower et al. 1992), 0.55 (Ellis et al. 1997), and 0.895 (Stanford et al. 1998). We find that owing to the age-metallicity conspiracy proposed by Worthey et al. (1996), the observed small scatter in the $C-M$ relation can be reproduced reasonably well even in star-forming galaxy mergers with the widely spread merging epoch. This result accordingly reinforces the recent results of Kauffmann \& Charlot (1998), in which the tight $C-M$ relation can be successfully reproduced by merger scenario of elliptical galaxy formation based on a hierarchical clustering scenario. This result furthermore implies that the previously suggested interpretation of the tightness of the $C-M$ relation at low and intermediate redshift ($z < 1$) is {\it not unique}, thus that the formation epoch of elliptical galaxies can be more widely spread than the classical passive evolution picture predicts. Thus, the above apparent inconsistency in the interpretation of the $C-M$ relation can be due primarily to the fact that previous studies claiming the coevality of elliptical galaxy formation did not explore so extensively possible variety in star formation history of elliptical galaxies. | The present study predicts that even if the epoch of major galaxy merging (i.e., the epoch of elliptical galaxy formation) is rather spread, both the tightness and the slope of the $C-M$ relation can be kept owing to the age-metallicity conspiracy originally proposed by Worthey et al (1996). This result accordingly provides a heuristic explanation for the result of Kauffmann \& Charlot (1998) in which the tight $C-M$ relation has been already reproduced in the merger scenario of elliptical galaxy formation based on the hierarchical clustering model. The conclusions derived in the present study however {\it seem} to be inconsistent with those derived in previous ones on the redshift evolution of the slope, zero-point and tightness of the $C-M$ relation of elliptical galaxies (Bower et al. 1992; Arag\'on-Salamanca et al. 1993; Kodama \& Arimoto 1997). In particular, the present numerical results {\it seem} to disagree with those of Kodama \& Arimoto (1997) and Kodama et al. (1998) (see also Gladders et al. 1998), which claim that the considerably less significant evolution of the slope of the $C-M$ relation rejects the age spread larger than 1 Gyr among elliptical galaxies. The apparent disagreement between the present study and the previous ones (e.g., Kodama \& Arimoto 1997; Kodama et al. 1998) is due essentially to the fact that the previous studies inevitably have over-interpreted the redshift evolution of the $C-M$ relation owing to the ad hoc assumption adopted in the previous studies. Although the previous studies are considerably sensible and valuable, it is important to point out the ad hoc assumptions adopted in the previous studies and thereby clarify the reason why the present conclusions are not consistent with those derived by the previous studies of Kodama \& Arimoto (1997) and Kodama et al. (1998). The following three are the ad hoc assumptions which inevitably lead the previous studies to draw the strong and general conclusion that formation of elliptical galaxies (especially in the cores of clusters) are {\it as a whole} coeval and occurred at high redshift. First is that elliptical galaxies are formed by {\it only one} initial starburst. Owing to this assumption, time evolution of global colors of elliptical galaxies depends exclusively on the epoch of initial burst of star formation (i.e., the epoch of elliptical galaxy formation in the previous study). As a result of this, the age difference between elliptical galaxies (i.e., the difference of the epoch of elliptical galaxy formation in the previous studies) can be more clearly reflected on the redshift evolution of the slope of the $C-M$ relation in the previous studies. Accordingly the observed less significant evolution of the slope of the $C-M$ relation is more likely to be interpreted as an evidence that supports the coevality of elliptical galaxy formation. It is certainly reasonable to claim that the observed evolution of the $C-M$ relation reject the `pure age' sequence model which demands that less luminous ellipticals have younger age. However, it seems not to be so reasonable to draw strong and general conclusion that elliptical galaxies are formed at $z>2$ {\it only} from the redshift evolution of the $C-M$ relation. Considering the first ad hoc assumption in the previous studies, what is more accurate and plausible interpretation on the observed evolution of the $C-M$ relation is just that the formation of {\it stellar populations} in {\it some} elliptical galaxies in the cores of {\it some} clusters (not the formation of galaxies with structural and morphological properties similar to those of ellipticals) can be coeval and occurred at higher redshift ($z > 2$). The second is that an elliptical galaxy in a cluster of galaxies at higher redshift is a precursor of an elliptical galaxy in a cluster at lower redshift. The third is that a cluster of galaxies observed at higher redshift is a precursor of a cluster of galaxies at lower one. These two ad hoc assumptions actually enable us to discuss the origin of elliptical galaxies in a more general way and thus lead us to draw more general conclusions on the formation epoch of elliptical galaxies. However, since there are no observational evidences which can provide the firm physical basis for the above assumptions at least now, it is questionable to give any general conclusions on the coevality of elliptical galaxy formation. Thus, these three assumptions adopted in the previous studies inevitably lead them to provide the strong and general conclusion that formation of elliptical galaxies are coeval and occurred at higher redshift. The present study, on the other hand, does not adopt the above three ad hoc assumptions, and rather relaxes these assumptions. Furthermore the present study instead allows both continuous and moderate star formation (not strong initial starburst) and the secondary starburst associated with galaxy merging, and assumes that the epoch of morphological transformation (into ellipticals) does not necessarily coincide with the epoch of galaxy formation (i.e., the epoch when the star formation begins). The evolution of the $C-M$ relation in the present study consequently does not depend so strongly on the difference in the formation epoch between elliptical galaxies (i.e., the epoch of major galaxy merging with the secondary starburst). As a result of this, the present merger model predicts that even if the formation epoch of elliptical galaxies (i.e., the epoch of galaxy merging) are rather spread, both the slope and tightness of the $C-M$ relation can be kept. Thus, the essential reason for the aforementioned apparent disagreement on the coevality of elliptical galaxy formation is that the present study does not adopt the above three ad hoc assumptions whereas the previous studies do. The interpretation on the redshift evolution of the $C-M$ relation in each model can depend strongly on the assumptions adopted by each model. It is safe for us to say that it is not clear, at least now, which of the two different conclusions on the coevality of elliptical galaxy formation is more plausible and reasonable. However, considering the above three ad hoc assumptions adopted in the previous studies, what is more reasonable interpretation on the redshift evolution of the $C-M$ slope is that only {\it stellar populations} (not elliptical morphology) in {\it some} ellipticals located in the cores of {\it some} clusters of galaxies are formed at higher redshift. We should not draw any {\it general} conclusions from the redshift evolution of the slope of the $C-M$ relation. Environmental difference of stellar populations (in particular, the existence of intermediate-age population) in early-type galaxies has been already indicated by a number of observational studies (e.g., Bower et al. 1990; Rose et al. 1994; Mobasher \& James 1996). On the other hand, the tightness and the slope of the $C-M$ relation of early-type galaxies are observationally revealed not to depend so strongly on galaxy environments. These two apparently inconsistent observational results on spectrophotometric properties of elliptical galaxies have called into the following question: ``Why does not the $C-M$ relation of early-type galaxies depend strongly on galaxy environments (e.g., between rich clusters and poor ones), though stellar populations and star formation histories in early-type galaxies probably depend on galaxy environments?'' To give a plausible answer for this question seems to be important because the above apparently inconsistent observational results give us valuable information both on the environmental difference in the details of physical processes of elliptical galaxy formation and on a certain mechanism for the tight $C-M$ relation. However, no extensive theoretical studies have yet addressed the above important question. The present study has shown that the age-metallicity conspiracy, which is achieved by younger and more metal-enriched stellar populations created in the secondary starburst of galaxy mergers, allows both the apparent age spread of elliptical galaxies and the tightness of the $C-M$ relation. This result seems to provide a clue to the above question. Since the {\it real} question concerning the tight $C-M$ relation is not to determine the typical epoch of elliptical galaxy formation but to give a convincing explanation for the reason why possible diversity in star formation histories of elliptical galaxies can allow the tight $C-M$ relation, more extensive theoretical studies including more variety of star formation history of elliptical galaxies and its likely dependence on galaxy environments are certainly worth for our deeper understanding of the origin of the tight $C-M$ relation. The present numerical results are consistent with recent observational results which suggest that coeval elliptical galaxy formation with initial starburst at higher redshift ($z > 2.0$) is not promising. Kauffmann, Charlot, \& White (1996) revealed that only about one-third of bright E/S0 galaxies in the sample of Canada-France Redshift Survey were already in the passive evolution phase at $z \sim 1.0$. Franceschini et al. (1997) found a remarkable absence of early-type galaxies at $z > 1.3$ in the $K$-band selected sample of early-type galaxies in the Hubble Deep Field (HDF), which suggests either that early-type galaxies are formed by galaxy merging with less prominent star formation or that a dust-polluted interstellar gas obscures forming elliptical galaxies till $z = 1.3$. Zepf (1997) demonstrated that strong deficit of galaxies with extremely red colors in the HDF means that the formation epoch of typical elliptical galaxies is $z < 5.0$. Sample galaxies in these studies are selected from {\it field ellipticals}, which possibly have star formation histories different from those of {\it cluster ellipticals}. Accordingly, it might not be plausible to derive strong conclusions on the formation epoch of ellipticals. However these observational results together with the present results seem to support the merger scenario which can naturally predict that the epoch of elliptical galaxy formation is rather extended ranging from high redshift to moderate one. Thus we have succeeded in pointing out that even if the epoch of elliptical galaxy formation (i.e., the epoch of major disk-disk galaxy merging, in this study) is rather widely spread, the tightness of the $C-M$ relation at low and intermediate redshift can be kept reasonably well. This result suggests that coevality of elliptical galaxy formation, which has been conventionally believed in the classical passive evolution picture, is {\it not unique} interpretation for the small scatter of the $C-M$ relation. This furthermore implies that {\it only} the tightness of the $C-M$ relation at a given redshift {\it does not necessarily} give strong constraints on the formation epoch of elliptical galaxies. Worthey et al. (1996) have already pointed out that the age-metallicity conspiracy can keep both the tightness of the Fundamental Plane and that of the $C-M$ relation in elliptical galaxies. The present numerical study, which is different from the Worthey's single stellar population analysis, has confirmed that the proposed age-metallicity conspiracy can actually operate to keep convincingly the tightness of the $C-M$ relation of ellipticals formed by disk-disk galaxy mergers. The present chemophotometric evolution model is, however, not so elaborated and realistic in that this model neither includes continuous gas accretion/merging expected from a specific cosmology (e.g., Baugh, Cole, \& Frenk 1996; Kauffmann \& Charlot 1998) nor considers important dynamical effects of galaxy merging on chemical and photometric evolution of galaxies (Bekki \& Shioya 1998). Accordingly it is our future study to confirm that the results derived in the present preliminary study can hold even for more sophisticated and realistic merger models. Furthermore, we should check whether or not observed redshift evolution of other fundamental relations such as the $\rm Mg_{2} - \sigma$ relation (Ziegler \& Bender 1997), the Fundamental Plane (van Dokkum \& Franx 1996), and the abundance ratio of [Mg/Fe] can be also reproduced self-consistently by our future merger model. | 98 | 4 | astro-ph9804221_arXiv.txt |
9804 | hep-ph9804378_arXiv.txt | In this paper we present the first analytic model for vorton formation. We start by deriving the microscopic string equations of motion in Witten's superconducting model, and show that in the relevant chiral limit these coincide with the ones obtained from the supersonic elastic models of Carter and Peter. We then numerically study a number of solutions of these equations of motion and thereby suggest criteria for deciding whether a given superconducting loop configuration can form a vorton. Finally, using a recently developed model for the evolution of currents in superconducting strings we conjecture, by comparison with these criteria, that string networks formed at the GUT phase transition should produce no vortons. On the other hand, a network formed at the electroweak scale can produce vortons accounting for up to $6\%$ of the critical density. Some consequences of our results are discussed. | \label{v-in} As first pointed out by Witten \cite{witten}, cosmic strings can in some circumstances (typically when the electromagnetic gauge invariance is broken inside the string) behave as `superconducting wires' carrying large currents and charges---up to the order of the string mass scale in appropriate units. The charge carriers can be either bosons or fermions (see \cite{vs} for a review). The former type occurs when it becomes energetically favourable for a charged Higgs field to have a non-zero vacuum expectation value in the string core; the latter happens when fermions couple to the string fields creating fermion zero modes. It is well known that arbitrarily large currents are not allowed---there is a critical value beyond which the current saturates. In other words, for large enough winding number per unit length, the superconducting condensate is quenched down, suppressing the current flow. Also, the current can decay by magnetic flux-line tunnelling; this can be used to impose constraints on allowed particle physics models. If superconducting strings carry currents, they must also carry charges of similar magnitude. This includes not only charges trapped at formation by the Kibble mechanism but also the ones due to string inter-commuting between regions of the string network with different currents. Just like with currents, charge densities cannot have arbitrarily large magnitude---there is a limit beyond which there will no longer be an energy barrier preventing the charge carriers from leaving the string. A rather important point is that the presence of charges on the string tends to counteract the current quenching effect discussed above. In fact, numerical simulations of contracting string loops at fixed charge and winding number have shown \cite{davsh} that a `chiral' state with equal charge and current densities is approached as the loop contracts. In this limiting chiral case, quenching is in fact eliminated completely. This has several important consequences. Strings that have trapped charges as a consequence of a phase transition can become superconducting even if the formation of a condensate was otherwise energetically unfavoured. More importantly, a string with both a charge and a current density will have a non-zero angular momentum. In the cosmological context, these strings would of course interact with the cosmic plasma, originating a number of interesting consequences. The most remarkable of these, however, has to do with the evolution of string loops. If a superconducting string loop has an angular momentum, it is semi-classically conserved, and it tries to resist the loop's tension. This will at least increase the loop's lifetime. If the current is too large, charge carriers will leave the string accompanied by a burst of electromagnetic radiation, but otherwise it is possible that dynamically stable loops form. These are called vortons \cite{vor}---they are stationary rings that do not radiate classically, and at large distances they look like point particles with quantised charge and angular momentum. Their cosmological significance comes from the fact that they provide very strong constraints on allowed particle physics models, since they behave like non-relativistic particles. According to current belief \cite{vor,vor2}, if they are formed at high enough energy scales they are as dangerous as magnetic monopoles, producing an over-density of matter in disagreement with observations. On the other hand, low-mass vortons could be a very interesting dark matter candidate. Understanding the mechanisms behind formation and evolution is therefore an essential cosmological task. The overwhelming majority of the work done on cosmic strings so far was concerned with the structureless Goto-Nambu strings (but see \cite{carternew} and references therein for some exceptions). In the case of work on vortons, this means that somewhat {\em ad-hoc} estimates had to be made for some properties of the cosmic string network---notably for microscopic quantities such as current and charge densities. This is despite the fact it has been recognised a long time ago that, even though they might be computationally very useful \cite{ms,ms2,ms2a}, Goto-Nambu models cannot realistically be expected to account for a number of cosmologically relevant phenomena, due to the very limited number of degrees of freedom available. Two such phenomena are the build-up of small-scale structure and charge and current densities. In this paper we fill this important gap by discussing the problem of vorton formation in the context of the superconducting string models of Witten \cite{witten} and of Carter and Peter \cite{wcp} (sections \ref{v-wt} and \ref{v-cp}). Strangely enough, the issue of the conditions for vorton formation has been so far neglected with respect to those of their stability and cosmological consequences. We will start by introducing these models and determining the microscopic string equations of motion in each case. It will be shown that in the relevant chiral limit these equations coincide---this also provides the first conclusive evidence of the validity of the supersonic elastic models of Carter and Peter \cite{wcp}. We then proceed to study the evolution of a number of loop solutions of these equations numerically (sections \ref{v-fl} and \ref{v-ex}), and from the results of this analysis parameters will be introduced which characterise the loop's ability to evolve into a vorton state (section \ref{v-vt}). Finally, we discuss a very simple phenomenological model for the evolution of the superconducting currents on the long cosmic string network \cite{mss}, based on the dynamics of a `superconducting correlation length' (sections \ref{v-cr}--\ref{v-ff}). Using this model we can therefore estimate the currents carried by string loops formed at all relevant times, and thus (in principle) decide if these can become vortons (section \ref{v-rs}) and calculate the corresponding density (section \ref{v-den}). Based on our results, we don't expect any GUT vortons to form at all. This is essentially because the friction-dominated epoch is very short for GUT-scale strings \cite{ms}, so their currents and charges are never large enough to prevent them from becoming relativistic---and therefore liable to losses. Even if they did form, they wouldn't be in conflict with the standard cosmological scenario if they decayed soon after the end of the friction-domination epoch. Hence we conclude that, in contrast with previously existing estimates \cite{vor,vor2}, one cannot at the moment rule out GUT superconducting string models. We should point out at the outset that there are essentially three improvements in the present work which justify the different end result for GUT-scale strings. Firstly, by analysing simple (but physically relevant) loop solutions of the microscopic string equations of motion for the Witten model, we can get a much improved idea of how superconducting loops evolve and of how (and under which conditions) they reach a vorton state. Secondly, by using a simple model for the evolution of the currents on the long strings \cite{mss} we can accurately determine the typical currents on each string loop at the epoch of its formation. Finally, the use of the analytic formalism previously introduced by the present authors \cite{ms,ms2a} allows us to use a quantitative description throughout the paper, and in particular to determine the loop sizes at formation. As will become clear below, when taken together these allow a detailed analysis of the process of vorton formation to be carried out, either in the Witten model (as is done in this paper) or any other that one considers relevant. In contrast, note that Davis \& Shellard \cite{vor} restrict themselves to the particular case of the initial Brownian Vachaspati-Vilenkin loops with Kibble currents, and do not consider the subsequent evolution of the network. On the other hand, Brandenberger {\em et al.} \cite{vor2} make rather optimistic order-of-magnitude estimates about the process of relaxation into a vorton state. As it turns out, for high energy GUT scales, all these loops become relativistic before reaching a vorton state. Finally, neither of these treatments has the benefit of a quantitative model for the evolution of the long-string network \cite{ms} which allows one to accurately describe the process of loop production. On the other hand, as we lower the string-forming energy scale we expect more and more efficient vorton production, and the 'old' scenario still holds. Therefore intermediate-scale superconducting strings are still ruled out, since they would lead to a universe becoming matter-dominated earlier than observationally allowed. Finally, at low enough energy scales, vortons will be a dark matter candidate. For example, for a string network formed around $T\sim10^2\,GeV$ (typical of the electroweak phase transition) they can provide up to $6\%$ of the critical density. A more detailed discussion of these issues is left to a forthcoming publication \cite{inprep}. Throughout this paper we will use fundamental units in which $\hbar=c=k_B=Gm^2_{Pl}=1$. | \label{v-cc} In this paper we have presented the first rigourous study of the cosmological evolution of superconducting strings in the limit of chiral currents. We have shown that in this limit the elastic string model of Carter \& Peter\cite{wcp} coincides with the model derived from first principles by Witten\cite{witten}. By analysing physically relevant loop solutions of the microscopic equations of motion for these strings, we have verified that the effect of frictional damping is crucial for vorton formation. We then defined suitable parameters characterising the evolution of these loops, and in particular whether or not they become vortons. In particular, we have established the usefulness of the `stability parameter' ${\overline n}$. In general, it is more difficult to form vortons when the string-forming phase transition is of first order. This is because such networks produce, during their evolution in the stretching regime, loops with a size close to that of the horizon; these will therefore be significantly affected by expansion, which tends to decrease the fraction of the loops's energy in the current---whereas friction tends to increase it. After introducing a simple `toy model' for the evolution of currents on the strings \cite{mss}, we have considered the cases of first and second-order GUT-scale string-forming and superconducting phase transitions (which is the most favourable GUT case of vorton formation since frictional forces can act longer). We have presented evidence suggesting that GUT-scale string networks might well produce no vortons, and that even if they do, this will not necessarily rule out such models. This is in contradiction with previous, less detailed studies \cite{vor,vor2}, and hence calls for a re-examination of a number of cosmological scenarios involving superconducting strings. Notably, these strings could be at the origin of the observed galactic magnetic fields \cite{msmag}. Finally, we have explicitly calculated the vorton density in two `extreme' cases to illustrate the method that one should follow once the microphysical properties of these networks are known in more detail. For electroweak-scale string networks, we have found that vortons can produce up to about $6\%$ of the critical density of the universe. On the other hand, it is conceivable that superconducting string networks formed at an energy scale $T\sim10^4-10^6\,GeV$ (depending on details of the model) can solve the dark matter problem. The detailed analysis presented in this paper for GUT stings can obviously be extended to other energy scales---this will be the subject of a forthcoming publication \cite{inprep}. Obviously, as we lower the energy scale, the frictional force becomes more and more important and acts for a longer time. Hence the vorton-forming region of parameter space increases, and by the electroweak scale almost all loops chopped off the long-string network will become vortons. We therefore conclude that in addition to the low-$G\mu$ regime (which as we saw includes the electroweak scale) where vortons can be a source of dark matter and to an intermediate-$G\mu$ range in which vortons would be too massive to be compatible with standard cosmology (thereby excluding these models), there is also a high-$G\mu$ regime (of which the GUT scale is part) in which vortons don't form at all and therefore no cosmological constraints based on them can be set. It is then curious (to say the least) that vorton constraints can be used to rule out cosmic string models in a wide range of energy scales $G\mu$, but not those formed around the GUT or the electroweak scales, where cosmic strings can be cosmologically useful. | 98 | 4 | hep-ph9804378_arXiv.txt |
9804 | astro-ph9804235_arXiv.txt | We investigate the possibility of accounting for the currently inferred primordial abundances of D, $^{3}$He, $^{4}$He, and $^{7}$Li by big bang nucleosynthesis in the presence of baryon density inhomogeneities plus the effects of late--decaying massive particles (X), and we explore the allowed range of baryonic fraction of the closure density $\Omega_{b}$ in such context. We find that, depending on the parameters of this composite model (characteristic size and density contrast of the inhomogeneities; mass--density, lifetime, and effective baryon number in the decay of the X particles), values as high as $\Omega_{b}h_{50}^{2}\simeq 0.25-0.35$ could be compatible with the primordial abundances of the light nuclides. We include diffusion of neutrons and protons at all stages, and we consider the contribution of the X particles to the energy density, the entropy production by their decay, the possibility that the X--products could photodissociate the light nuclei produced during the previous stages of nucleosynthesis, and also the possibility that the decay products of the X--particles would include a substantial fraction of hadrons. Specific predictions for the primordial abundance of Be are made. | Standard homogeneous big bang nucleosynthesis could have produced the observationally inferred primordial abundances of D, $^{3}$He, $^{4}$He, and $^{7}$Li, provided that the baryon fraction of the cosmic closure density $\Omega_{b}$ would lie in the range: $$0.04\lapprox \Omega_{b}h_{50}^{2}\lapprox 0.08\eqno(1)$$ \noindent where $h_{50}$ is the Hubble constant in units of 50 km s$^{-1}$ Mpc$^{-1}$ (Walker et al. 1991; Copi, Schramm, \& Turner 1995). For the long--time most favored cosmological model, a flat Universe with $\Omega_{M} = 1$ and $\Omega_{\Lambda} = 0$ (those being, respectively, the fractional contributions of matter and vacuum energy densities to the closure density), the upper bound to $\Omega_{b}$ would mean that most matter in the Universe should be in nonbaryonic form. Given the far--reaching implications of the dominance of nonbaryonic dark matter, possible alternatives to homogeneous big bang nucleosynthesis have been explored, especially during the last 15 years or so. The suggestion that the quark--hadron phase transition might be first--order and generate baryon inhomogeneities (Witten 1984) led to the calculation of the possible effects on primordial nucleosynthesis (Applegate \& Hogan 1985; Applegate, Hogan, \& Scherrer 1987; Malaney \& Fowler 1988). The goal was to see whether inhomogeneous big bang nucleosynthesis with $\Omega_{b} = 1$ might account for the primordial light--element abundances. Besides a first--order quark--hadron phase transition, other mechanisms might also generate baryon inhomogeneities. Much of the work in this line is reviewed by Malaney \& Mathews (1993). However, the recent studies, treating accurately the coupling between baryon diffusion and nucleosynthesis, show that the upper limit on $\Omega_{b}$ set by the light--element abundances does not significantly differ from that obtained for homogeneous big bang nucleosynthesis (Mathews, Schramm, \& Meyer 1993; Thomas et al. 1994). This last conclusion, though, has very recently been challenged by Orito et al. (1997), who explore the dependence of primordial nucleosynthesis on the geometry of baryon inhomogeneities and find that cylindrical geometry might allow to satisfy the observational constraints with baryon fractions as high as $\Omega_{b}h_{50}^{2}\lapprox 0.2$. A different approach has been to explore the possible modifications of the yields from homogeneous big bang nucleosynthesis by the effects of the decay of unstable massive particles ($M\gapprox few\ GeV$), produced at earlier stages in the evolution of the Universe and with half--lives longer than the standard nucleosynthesis epoch ($\tau_{x}\gapprox 10^{4}\ s$) (Audouze, Lindley, \& Silk 1985; Dom\'{\i}nguez--Tenreiro 1987; Yepes \& Dom\'{\i}nguez--Tenreiro 1988; Dimopoulos et al. 1988). Gravitinos produced during reheating at the end of inflation are a possible example of such particles. In Dimopoulos et al. (1988), the emphasis is put on the resulting hadron cascade. The main problem encountered in this model is the predicted overproduction of $^{6}$Li: $^{6}Li/^{7}Li\gg 1$, whereas observations show $^{6}Li/^{7}Li\lapprox 0.1$. Although they have only been considered separately, baryon inhomogeneities and the presence of unstable massive particles decaying when the Universe has already cooled down below $T_{9}\simeq 0.4$ are by no means mutually exclusive. Here we explore their combined effects on the primordial abundances of the light elements. The parameter space now has, of course, a dimension which is the sum of those for the two separate cases: characteristic size and density contrast of the inhomogeneities, mass--density, lifetime, and mode of decay of the massive particles. We find that there are regions in such extended parameter space where values of $\Omega_{b}$ as high as $\Omega_{b}h_{50}^{2}\simeq 0.35$ would still be compatible with the primordial abundances of the light nuclides inferred from observations. Such values of $\Omega_{b}$ are of the same order as the low values for $\Omega_{M}$ currently derived from a variety of sources, including high--redshift supernova searches (Perlmutter et al. 1998; Garnavich et al. 1998). Our results thus suggest again the possibility that all the matter in the Universe could be baryonic. On the other hand, recent determinations of the D abundance in high--redshift QSO absorbers, when confronted with the currently inferred primordial $^{4}$He abundance, might be in conflict with the predictions of standard, homogeneous big bang nucleosynthesis for $N_{\nu} = 3$ (Steigman 1998): the ``low'' high--redshift D abundances (which appear more reliable) would indicate too high a value of $\Omega_{b}$ to be compatible with that corresponding to the $^{4}$He abundance. Since it is hard to tell whether this conflict points to new physics or just to systematic errors in the derivation of abundances, Steigman, Hata, \& Felten (1998) have discarded the constraint on $\Omega_{b}$ from standard big bang nucleosynthesis and turned to other observational constraints to determine the key cosmological parameters. The results from our composite model, by showing how minor deviations from the standard hypotheses can produce agreement with the primordial abundances inferred from observations, support that attitude. Besides, as we will see, the combined effects of inhomogeneities plus late--decaying particles might solve the conflict between D and $^{4}$He abundances. | We have shown, by means of a simple model, that the combined effects on big bang nucleosynthesis of baryon inhomogeneities plus the decay of unstable, relatively long--lived massive particles, giving rise to both electromagnetic and hadron cascades, might be to allow agreement with the primordial light--element abundances inferred from observations for values of $\Omega_{b}$ much higher that those allowed by standard, homogeneous nucleosynthesis. The upper limit might be as high as $\Omega_{b}h_{50}^{2}\simeq 0.35$. The values obtained here are of the same order as the low $\Omega_{M}$ values now being derived from a variety of sources and, therefore, they pose in new terms the question of whether all matter in the Universe could be baryonic. A testable prediction of the model is the production of a $^{9}$Be abundance that is of the order of current observational upper limits. On the other hand, in the parameter region of our model where there is agreement between predicted and observationally inferred primordial light--element abundances, given values of $\Omega_{b}$ (or, equivalently, $\eta_{10}$) always predict ``low'' D abundances (in the sense of the high--redshift abundances referred to in the Introduction), thus potentially eliminating the conflict with the $^{4}$He abundance pointed out by Steigman (1998). The model presented here deals with inhomogeneities in a very simplified way. A futher step will be to examine the effects of the geometry of the density fluctuations on the outcome. Orito et al. (1997) have already shown that cylindrical shell geometry alone (without the extra effects of late--decaying particles) might allow $\Omega_{b}\lapprox 0.2$ (but for density contrasts $R\sim 10^{6}$, much higher than those considered here). Another extension of the model will be to consider particles with shorter half--lives, decaying at the time when thermonuclear reactions are still taking place. | 98 | 4 | astro-ph9804235_arXiv.txt |
9804 | astro-ph9804145_arXiv.txt | The neutrino emissivities resulting from direct URCA processes in neutron stars are calculated in a relativistic Dirac-Hartree approach in presence of a magnetic field. In a quark or a hyperon matter environment, the emissivity due to nucleon direct URCA processes is suppressed relative to that from pure nuclear matter. In all the cases studied, the magnetic field enhances emissivity compared to the field-free cases. | 98 | 4 | astro-ph9804145_arXiv.txt |
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9804 | astro-ph9804286_arXiv.txt | We present a detailed morphological analysis of the galaxy populations in the first two clusters to be completed in an extensive observational study of nine high-redshift clusters of galaxies (Oke, Postman \& Lubin 1988). These two clusters, CL0023+0423 and CL1604+4304, are at redshifts of $z = 0.84$ and $z = 0.90$, respectively. The morphological studies are based on high-angular resolution imagery taken with WFPC2 aboard the {\it Hubble Space Telescope}. These data are combined with deep, ground-based $BVRI$ photometry and spectra taken with the Keck 10-meter telescopes. The morphological classifications presented in this paper consist of two parts. Firstly, we provide a quantitative description of the structural properties of $\sim 600$ galaxies per cluster field using the Medium Deep Survey automated data reduction and object classification software (Griffiths et al.\ 1994; Ratnatunga, Ostrander \& Griffiths 1997). This analysis includes the galaxy position, photometry, and best-fit bulge+disk model. Secondly, for the brightest subsample of $\sim 200$ galaxies per cluster field, we provide a more detailed morphological description through a visual classification based on the revised Hubble classification scheme (e.g.\ Sandage 1961; Sandage \& Bedke 1994). Based on these classifications, we have examined the general relation between galaxy morphology and other photometric and spectral properties. We find that, as expected, the elliptical and S0 galaxies are redder, on average, than the spirals and irregulars. In addition, there is a strong correlation between morphology and spectral type. Of the galaxies that are visually classified as ellipticals, the majority show K star absorption spectra which are typical of nearby, red early-type galaxies; however, a few are actually blue compact galaxies with spectra characterized by fairly strong, narrow emission lines. Normal late-type galaxies typically have spectra with blue colors and [\ion{O}{2}] emission, while the presence of strong star-formation features, such as extremely high equivalent width [\ion{O}{2}], ${\rm H\beta}$, and/or [\ion{O}{3}] emission, is always accompanied by peculiar morphologies which suggest recent mergers or interactions. We have used the statistical distributions of cluster galaxy morphologies to probe the overall morphological composition of these two systems. This analysis reveals that the two clusters contain very different galaxy populations. CL0023+0423 has a galaxy population which is more similar to groups of galaxies and the field. This system is almost completely dominated by spiral galaxies. CL1604+4304, however, has a morphological composition which is more typical of a normal, present-day cluster; early-type galaxies comprise $\sim 76\%$ of all galaxies brighter than $M_{V} = -19.0 + 5~{\rm log}~h$ in the central $\sim 0.5~h^{-1}~{\rm Mpc}$. The ratio of S0 galaxies to ellipticals in this cluster is $1.7 \pm 0.9$, consistent with local cluster populations. The morphological results support the conclusions of the dynamical analysis presented in the second paper of this series (Postman, Lubin \& Oke 1998). CL0023+0423 consists of two galaxy groups which are separated by $\sim 2900~{\rm km~s^{-1}}$ in radial velocity. CL1604+4304, on the other hand, has a velocity distribution indicating that it is already well-formed and relaxed. The morphological composition, velocity dispersion, and implied mass of the CL1604+4304 system are consistent with an Abell richness class 2 or 3 cluster. | The study of the galaxy populations of rich clusters provides important constraints on the formation mechanisms of both clusters and galaxies. Present--day clusters show a distinct correlation between the structure of the cluster and the galaxy population. Irregular, open clusters, such as Virgo, are spiral--rich. These systems show no single, central condensation, though the galaxy surface density is at least five times as great as the surrounding field ($n_{\rm gal} > 30~h^{3}~{\rm galaxies~{Mpc}^{-3}}$). These clusters are often highly asymmetric and have significant degrees of substructure. Dense, centrally concentrated clusters, such as Coma, contain predominantly early--type galaxies in their cores (Abell 1958; Oemler 1974; Dressler 1980a,b; Postman \& Geller 1984). These clusters have a single, prominent concentration among the bright member galaxies and typically display a high--degree of spherical symmetry, though this does not preclude evidence of some substructure. Central densities can reach as high as $10^{4}~h^{3}~{\rm galaxies~{Mpc}^{-3}}$ (e.g\ Bahcall 1975; Dressler 1978). In these regions, spiral galaxies comprise less than 10\% of the cluster population, while elliptical (E) and S0 galaxies make up 90\% or more of the population. The ratio of S0s to ellipticals is typically S0/E $\sim 2$ (Dressler 1980a). The galaxy content of clusters is part of the general morphology--density relation of galaxies; as the local density increases, the fraction of elliptical and S0 galaxies increases, while the fraction of spiral galaxies decreases (Hubble 1936; Dressler 1980a,b; Postman \& Geller 1984). Previous studies of clusters of galaxies at $z < 1$ have revealed significant evolution in the morphology and the color of the cluster members. One of the most notable of these changes is the progressive blueing of cluster's galaxy population with redshift, a trend first observed by Butcher \& Oemler (1984). They found that the fraction of blue galaxies in a cluster is an increasing function of redshift, indicating that clusters at redshifts of $z \sim 0.5$ are significantly bluer than their low--redshift counterparts. At redshifts of $z \sim 0.4$, the fraction of blue galaxies is $\sim 20\%$. Recent HST image data reveal that many of these blue galaxies are either ``normal'' spirals or have peculiar morphologies, resulting in non--elliptical fractions which are 3 to 5 times higher than the average current epoch cluster (Dressler et al.\ 1994; Couch et al.\ 1994; Oemler, Dressler \& Butcher 1997; Dressler et al.\ 1997). Detailed photometric observations of other intermediate redshift ($z \simless 0.4$) clusters have confirmed the original results of Butcher \& Oemler (e.g.\ Millington \& Peach 1990; Luppino et al.\ 1991; Rakos \& Schombert 1995). Even though these clusters show an increased fraction of blue galaxies, they still contain a population of E/S0s distinguished by extremely red colors and a tight color--magnitude (CM) relation (a ``red envelope''). Both the mean color and the CM relation are consistent with that of present--day ellipticals (e.g.\ Sandage 1972; Butcher \& Oemler 1984; Arag\'on-Salamanca et al.\ 1991; Luppino et al.\ 1991; Molinari et al.\ 1994; Dressler et al.\ 1994; Smail, Ellis \& Fitchett 1994; Stanford, Eisenhardt \& Dickinson 1995). At redshifts of $z \simgreat 0.4$, the red envelope has moved bluewards with redshift (Arag\'on-Salamanca et al.\ 1993; Smail et al.\ 1994; Rakos \& Schombert 1995; Oke, Gunn \& Hoessel 1996; Lubin 1996; Ellis et al.\ 1997; Stanford, Eisenhardt \& Dickinson 1997). At $z \sim 0.9$, there are few cluster members with colors nearly as red as present--day ellipticals. The color distribution of this high-redshift elliptical population is relatively narrow, and the trend is uniform from cluster to cluster; this suggests a homogeneous population which formed within a narrow time span (e.g.\ Bower, Lucey \& Ellis 1992a,b). Dickinson (1995) finds similar results in a cluster of galaxies which is associated with the $z = 1.206$ radio galaxy 3C 324. The galaxies in this cluster exhibit a narrow, red locus in the CM magnitude diagram. This branch is $\sim 0.6$ mag bluer than the expected ``no--evolution'' value, though the intrinsic rms color scatter is only 0.2 mag. The observed color trend for the red envelope of galaxies in this data is consistent with passive evolution of an old stellar population formed by a single burst of star formation at redshifts of $z \simgreat 2$. The reasonably small color scatter would imply closely synchronized intra--cluster star formation (Bower et al.\ 1992a,b; Arag\'on-Salamanca et al.\ 1993; Dickinson 1995; Ellis et al.\ 1997; Stanford et al.\ 1997). The high-resolution imaging of HST has been essential in understanding the evolutionary processes occurring at intermediate redshifts (see e.g.\ Abraham et al.\ 1996a). Morphological classifications can be made on scales of $\sim 1~{\rm kpc}$, providing a direct comparison with ground-based classifications of nearby galaxies. A comprehensive survey of 10 clusters of galaxies at $z = 0.37 - 0.56$ has revealed a significant change relative to local clusters in the composition and behavior of the galaxy populations (Smail et al.\ 1997, hereafter S97; Dressler et al.\ 1997, hereafter D97). The authors have visually classified over 6000 galaxies based on the Revised Hubble Scheme used to classify nearby galaxies (e.g.\ Sandage 1961; Sandage \& Bedke 1994). These classifications are used to quantify the morphological composition of each cluster. Their results indicate that the morphology--density relation is qualitatively similar to that in the local universe in those intermediate redshift clusters which are centrally-concentrated and compact; however, the relation is non-existent in the loose, open clusters. Even so, all of the clusters exhibit a roughly similar make-up of galaxy morphologies. The fraction of ellipticals is the same or larger than that in local clusters; the S0 fraction, however, is $\sim 2-3$ times lower, with a corresponding increase in the cluster spiral population. These findings imply that the elliptical population is already in place by $z \sim 0.5$, but a large fraction of the S0 galaxies are formed between redshifts of $z \sim 0.5$ and $z = 0$ (D97). However, it should be noted that these classifications are typically derived from images which are not of comparable quality to the local data. Because of such uncertainty, the observed evolution in the S0 population is still in contention (e.g.\ Stanford et al.\ 1997; Andreon, Davoust \& Helm 1997; Andreon 1998). Because there appears to be significant evolution occurring between redshifts of $z \sim 0.5$ and the present epoch, it is critical to extend these detailed observations to even higher redshifts if we are to understand the formation of galaxy morphology, as well as the mechanisms and timescales of this evolution. Therefore, we have undertaken an extensive observational program to study nine candidate clusters of galaxies at redshifts of $z > 0.6$. The cluster sample was chosen from the Gunn, Hoessel \& Oke (1986) survey and the Palomar Distant Cluster Survey (PDCS; Postman et al.\ 1996). For each cluster, we are in the process of obtaining deep $BVRI$ photometry from Keck and deep $K$ photometry from the KPNO 4-meter, low-resolution spectra from Keck, and high angular resolution imagery from HST. The observations and data processing procedures of this survey are the subject of the first paper in this series (Oke, Postman \& Lubin 1998; hereafter Paper I). The first two clusters to be completed in this observational program are CL0023+0423 and CL1604+4304 at redshifts of $z = 0.84$ and $z = 0.90$, respectively (see Paper I). In this paper, we have used HST images to undertake a detailed morphological analysis of the galaxy populations in the central regions of these two clusters. The reduction and analysis of the Keck $BVRI$ photometry and spectra of the galaxies in these cluster fields are discussed in the second paper of this series (Postman, Lubin \& Oke 1998; hereafter Paper II). However, the galaxy parameters presented in Paper II are used in this paper, specifically for a comparison with the morphological properties. In Sect.\ 2, we provide a brief description of the data. In Sects.\ 3 and 4, we describe the automated and visual galaxy classification procedures used in this paper and present a comparison between the two techniques. In Sect.\ 5, we examine the morphologies of the galaxies in the two cluster fields. This includes the relationship between morphology and other galaxy properties, as well as the overall distribution of morphologies in the cluster. A discussion and summary of our conclusions are presented in Sects.\ 6 and 7. In the following analyses, we have assumed $q_{0} = 0.1$ (e.g\ Carlberg et al.\ 1996) and $H_{0} = 100~h~{\rm km~s^{-1}~Mpc^{-1}}$. | As part of an observational program to study distant clusters of galaxies, we have examined the morphological properties of the galaxies in two cluster fields, CL0023+0423 at $z = 0.84$ and CL1604+4304 at $z = 0.90$, using high-resolution HST images. The morphology of the individual galaxies have been studied by two methods; 1) a quantitative description of the structural properties of $\sim 600$ galaxies per cluster field is provided by the Medium Deep Survey automated data reduction and ``bulge+disk'' object classification software; 2) the brightest subsample of $\sim 200$ galaxies per cluster field are assigned a more detailed morphological description through a visual classification based on the revised Hubble scheme. A comparison between the two techniques shows that there is a reasonable correlation between the parameters of the automated and visual classifications (see also Lahav et al.\ 1995). To investigate the morphological composition of the two galaxy clusters, we have used the visual classifications of the brightest subsample of galaxies in each field. Our main conclusions are summarized below. \newcounter{discnt} \begin{list} {\arabic{discnt}.} {\usecounter{discnt}} \item The color-magnitude diagrams and the color histograms of all (field + cluster) galaxies in the two cluster field show a clear progression in color between early- and late-type galaxies. As expected, the elliptical and S0 galaxies are redder, on average, than the spirals and irregulars. This trend is also apparent in the color ages which represent the time since the last period of major star formation. 83\% of the galaxies classified as late-type (spiral or irregular/peculiar) have color ages less than 2 Gyr. In contrast, 55\% of the galaxies classified as early-type have color ages of greater than 2 Gyr, and 73\% of all galaxies with ages greater than 3 Gyr are classified as early-type. In addition, cluster galaxies are typically older than field galaxies at similar redshifts. This is due in large part to the fact that there are more early-type galaxies in these systems. \item We see a distinct correlation between the galaxy morphology and the corresponding spectral features. We have specifically examined this relation for those galaxies which are confirmed cluster members. The majority of galaxies that are visually classified as ellipticals show spectra which are typical of nearby, red elliptical galaxies. However, some of the galaxies visually classified as ellipticals turn out to be blue compact galaxies with spectra characterized by fairly strong, narrow emission lines. Normal late-type galaxies typically have spectra with blue colors and [\ion{O}{2}] emission, while the presence of strong star-formation features, such as extremely high equivalent width [\ion{O}{2}], ${\rm H\beta}$, and/or [\ion{O}{3}] emission, is always accompanied by peculiar morphologies which suggest recent mergers or interactions. \item Despite being at very similar redshifts, the two cluster systems contain very different galaxy populations as indicated by their background-subtracted morphological distributions. We have examined all galaxies brighter than $M_{V} = -19.0 + 5~{\rm log}~h$ in the central $\sim 0.5~h^{-1}~{\rm Mpc}$ of the cluster. CL0023+0423 has a galaxy population which is more typical of groups and the field population. The numbers from the statistical distribution are consistent with almost all of the galaxies being normal spirals. The spectral results support these numbers, independently indicating spiral fractions of 66\% or more. CL1604+4304, in contrast, has a morphological composition which is characteristic of a normal, present-day rich cluster. Early-type galaxies comprise 76\% of all galaxies in this region. In this population, the ratio of S0 galaxies to ellipticals is $1.7 \pm 0.9$, consistent with galaxy populations found in local clusters (Dressler 1980a). \item The morphological results support the conclusions of the dynamical analysis presented in Paper II. CL0023+0423 is apparently two independent systems separated in velocity by $\sim 2900~{\rm km~s^{-1}}$. The velocity dispersions and implied masses indicate that these systems are similar to local galaxy groups or poor clusters. Though this may be a chance projection, the dynamical and morphological evidence may indicate that we are seeing the merger of two spiral-dominated galaxy groups (see Lubin, Postman \& Oke 1998a). The velocity histogram of CL1604+4304, on the other hand, is consistent with a Gaussian distribution, implying that this system formed at an earlier epoch and is already relaxed. The velocity dispersion and implied mass of this system are consistent with an Abell richness class 2 or 3 cluster. \end{list} \vskip 0.5cm We thank the anonymous referee for his thorough review of this paper. Alan Dressler, Chris Fassnacht, and Ian Smail are thanked for useful discussions, comments, and material aids to this paper. It is also a great pleasure to thank Allan Sandage for his generous gift of time and invaluable expertise to this project. The W.M. Keck Observatory is operated as a scientific partnership between the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. It was made possible by generous financial support of the W. M. Keck Foundation. LML graciously acknowledges support from a Carnegie Fellowship. This research was supported in part by {\it HST} GO analysis funds provided through STScI grant GO-06000.01-94A and {\it HST} Archival grant 7536. \clearpage | 98 | 4 | astro-ph9804286_arXiv.txt |
9804 | astro-ph9804208_arXiv.txt | s{Galactic synchrotron emission is a potentially confusing foreground, both in total power and in polarization, to the Cosmic Microwave Background Radiation. It also contains much physical information in its own right. This review examines the amplitude, angular power spectrum and frequency spectrum of the synchrotron emission as derived from the presently available de-striped maps. There are as yet no maps at arcminute resolution at frequencies above 2.4 GHz. This incomplete information is supplemented with data from supernovae, which are thought to be the progenitors of the loops and spurs found in the Galactic emission. The possible variations of the frequency spectral index from pixel to pixel are highlighted. The relative contributions of free-free and synchrotron radiation are compared, and it is concluded that the free-free contribution may be smaller than had been predicted by COBE. New high resolution polarization surveys of the Galactic plane suggest detail on all scales so far observed. At high latitudes the large percentage polarisation means that the foreground contamination of the polarised CMB signal will be more serious than for the unpolarized radiation.} | Galactic emission at radio wavelengths is important to understand in its own right. Moreover it is crucial to be able to quantify and remove this component as a foreground to the cosmic microwave background (CMB). Both synchrotron and free-free emission contribute to this foreground, with the synchrotron emission dominating at low frequencies ($\leq$1 GHz). The synchrotron emissivity is a function of both the relativistic (cosmic ray) density and the local magnetic field strength. The luminosity at frequency $\nu$ is given by \be I(\nu) = const \; L N_0 B^{(p+1)/2} \nu^{-(p-1)/2} \ee \noindent where $N_0$ is the density of relativistic electrons, $L$ is the emission depth, $B$ is the magnetic field and the relativistic electron energy spectrum\cite{ve} is given by $dN/dE = N_0 E^{-p}$. The radio spectral index is $\alpha= (p-1)/2$ in energy terms or $2+\alpha$ when expressed as a brightness temperature $T_B \propto \nu^{-(2+\alpha)} = \nu^{-\beta}$. Within the interstellar magnetic field of 2 to 5 microgauss, emission at GHz frequencies is characteristically from relativistic electrons with an energy of 1 to 10 GeV. Both $B$ and $N_0$, as well as $p$, will vary from point to point in the Galactic disk and nearby halo. The cosmic ray electrons are thought to originate mainly in supernovae then diffuse outwards in the expanding remnant. Structure will be formed in the remnant as it collides with the non-uniform ambient medium. The magnetic field will be likewise amplified in compression regions and vary in strength and direction. The net effect is to produce elongated synchrotron emission structures on a wide range of scales. The spectral index of the emission will vary with position for two reasons. Firstly the electron spectral index varies from one supernova to another and secondly the spectrum steepens ($\Delta p = +1$) with time due to radiation energy loss thus giving an age-dependent spectral index. This paper will describe the synchrotron features in and near the Galactic plane which are believed to give rise to the structures seen at higher galactic latitudes. All-sky and large area surveys are assessed to give information about the amplitude and spectrum of the high latitude emission which is a potential confusing foreground to the CMB. Comments are given about the role of synchrotron polarization and of free-free emission. | 98 | 4 | astro-ph9804208_arXiv.txt |
|
9804 | astro-ph9804102_arXiv.txt | We report on {\it ROSAT} HRI observations of the $z=0.61$ radio galaxy 3C\,220.1. The X-ray emission from this object consists of an extended component, which we attribute to luminous cluster emission, and a compact central source. The compact component is too bright to be modelled as a cooling flow under some plausible assumptions for the hot gas temperature and distribution; we suggest instead that it is directly related to the core of the radio source. The X-ray flux of the compact component is consistent with the prediction of Worrall \etal\ (1994) that all powerful radio galaxies should have a central jet-related X-ray emission component that is proportional in strength to the radio core flux density. Other observations of distant 3CR radio sources are consistent with this model. | In unified models of powerful radio sources (e.g.\ Barthel 1989) core-dominated quasars, lobe-dominated quasars and radio galaxies are the same objects, with the apparent differences attributed to the effects of relativistic beaming and anisotropic obscuration on a population of sources oriented randomly with respect to the line of sight. In quasars we see the AGN directly, while radio galaxies have their jet axis at large angles to the line of sight so that a torus of gas and dust obscures optical continuum and broad-line emission from the AGN. This torus should also obscure soft X-ray emission originating close to the central engines, leading to suggestions (e.g.\ Crawford \& Fabian 1996) that X-ray emission from powerful radio galaxies should in general be dominated by thermal emission from hot cluster gas. However, an additional X-ray component may arise from the radio-emitting plasma directly, through synchrotron or synchrotron self-Compton radiation, or indirectly, through mechanisms such as the inverse-Compton scattering of external photons. If the different components can be separated, X-ray observations can provide physical insights into the active nucleus, jets and large-scale environment. There is substantial evidence that unabsorbed radio-related non-thermal X-rays are seen. Fabbiano \etal\ (1984) first stressed that a correlation between the total soft X-ray and radio-core luminosity in radio galaxies implied a nuclear, jet-related origin for at least some of the X-ray emission. More recent work has strengthened this conclusion, as high-resolution X-ray observations have allowed point-like and extended components to be separated (Worrall \& Birkinshaw 1994; Edge \& R\"ottgering 1995; Worrall 1997). Although component separation is generally better for closer, less powerful radio galaxies, compact soft X-ray emission is also seen in powerful narrow-line FRII radio galaxies that are known to lie in sparse environments, where cluster emission is not a source of confusion (e.g.\ Hardcastle, Birkinshaw \& Worrall 1998a). The full complexity is illustrated by X-ray observations of the nearby powerful cluster radio galaxy Cygnus A. {\it EXOSAT} and {\it Ginga} have found evidence for highly obscured core emission, ($N_H \sim 4 \times 10^{23}$ cm$^{-2}$; Arnaud \etal\ 1987, Ueno \etal\ 1994) which should not have been seen with {\it ROSAT}, given its low-energy X-ray passband. However, an unresolved core component was detected with the {\it ROSAT} HRI (Harris, Perley \& Carilli 1994), implying that Cygnus A's core has both an absorbed and an unabsorbed X-ray component; while the absorbed component may be associated with emission from the AGN, the unabsorbed component may be radio-related (Worrall 1997). Birkinshaw \& Worrall (1993) argued in the case of NGC 6251 that a plausible source for such compact X-ray emission is synchrotron self-Compton (SSC) radiation from the base of the radio jet, originating on scales larger than that of the torus and so avoiding absorbtion. Such emission would be suppressed, but not eliminated, by relativistic beaming effects in radio galaxies and [as discussed in Worrall \etal\ (1994) and references therein] almost certainly dominates the X-ray emission in core-dominated quasars. It is therefore important in interpreting the X-ray emission from radio galaxies to have spectral or spatial information capable of distinguishing a non-thermal or compact component from thermal or extended emission. High-redshift powerful radio galaxies are important as counterparts to quasars in unified models, and Worrall \etal\ (1994) performed such a spatial separation for the $z = 1$ radio galaxy 3C\,280. They showed that the unresolved component fell on an extension of the correlation between X-ray and radio core flux observed for core-dominated quasars, consistent with the model discussed above. In this paper we report X-ray observations of the radio galaxy 3C\,220.1 with the {\it ROSAT} HRI. We use the high resolution of the HRI to constrain the contributions from compact and extended emission. 3C\,220.1 is an FRII (Fanaroff \& Riley 1974) narrow emission-line radio galaxy at $z=0.61$ (Spinrad \etal\ 1985). With $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0$, used throughout the paper, its 178-MHz luminosity is $3.6 \times 10^{27}$ W Hz$^{-1}$ sr$^{-1}$. Radio imaging (Burns \etal\ 1984; Jenkins, Pooley \& Riley 1977; Harvanek \& Hardcastle 1998) shows it to be a typical classical double object with largest angular size 35 arcsec (see Fig.\ \ref{contour}, inset); at this redshift one arcsecond corresponds to 8.94~kpc, so the projected linear size of 3C\,220.1 is about 300 kpc. Burns \etal\ (1984) report an unusually prominent one-sided jet in the eastern lobe, and the radio core is also comparatively prominent, which may be an indication that the source is significantly affected by relativistic beaming, although no broad emission lines are reported by Spinrad \etal\ (1985). The prominent radio core, compared with the weak cores of the objects observed by Worrall \etal\ (1994), was the motivation for the present observations, since it allows us to probe the possible core X-ray and radio association. Optical observations show no evidence for a rich cluster near 3C\,220.1, although the presence of a gravitational lens arc with $z=1.49$ in HST observations implies a deep potential well (M.\ Dickinson, private communication, 1997). | We observed 3C\,220.1 in order to test the hypothesis of Worrall \etal\ (1994) that all powerful radio sources exhibit a compact X-ray component related to their radio core. We find that there is strong evidence for a compact X-ray component in this object, that it is too bright to be attributed to a cooling flow under some simple assumptions, and that it follows the expected positive correlation between radio and X-ray core flux or luminosity. 3C\,220.1 lies slightly above the line of slope unity plotted through the core-dominated quasars (Fig.\ \ref{cores}). Together with its unusually prominent radio core and one-sided radio jet, this may be an indication that we are viewing the source at an angle to the line of sight which is close to the radio galaxy-quasar boundary, so that its central X-ray emission includes a non-jet-related component originating close to the AGN. Extended X-ray emission is also unequivocally detected around 3C\,220.1, with core radius and luminosity comparable to that of nearby rich clusters. This is qualitatively consistent with the HST detection of a gravitational lensing arc near the source. | 98 | 4 | astro-ph9804102_arXiv.txt |
9804 | astro-ph9804334_arXiv.txt | Deep infrared and optical images are presented of three proposed remnants of Thorne-$\dot{\rm Z}$ytkow Objects (T$\dot{\rm Z}$O). In particular, the IR data go several infrared magnitudes deeper than previous observations and in at least one case reveal the existence of weak objects within the error circles. It is argued, however, that none of the objects is likely to be the binary companion to the X-ray source in that region. These data present severe limits on any possible star or residual envelope at the distances of the respective X-ray objects. | Recent work by several authors, especially van Paradijs et al. (1995) and Mereghetti \& Stella (1995), have identified a group of pulsating X-ray sources that appear to have no detectable optical counterpart. In addition, these objects show no evidence of any binarity and all have pulse periods in the 5-10s range. Their X-ray luminosities are all sufficiently high that accretion must be occurring on to the neutron star from some source other than the interstellar medium. One possible scenario is that they are the first objects identifiable as the remains of Thorne-$\dot{\rm Z}$ytkow Objects (T$\dot{\rm Z}$Os). The progenitor T$\dot{\rm Z}$Os would have consisted of a neutron star embedded in the core of another star (Thorne \&$\dot{\rm Z}$ytkow 1977). It is thought that these T$\dot{\rm Z}$Os could be formed as a result of the evolution of high mass systems through a common envelope phase (Ghosh et al. 1997). Once formed, thermonuclear burning via the {\it r-p} process occurs in the core region of the T$\dot{\rm Z}$O. This supports the highly convective outer envelope against gravitational forces. However, once there is insufficient fuel remaining for the {\it r-p} process to occur efficiently, the outer envelope collapses on the Kelvin-Helmholtz timescale (Podsiadlowski, Cannon and Rees, 1995). Some of the envelope material then forms into a disk around the neutron star with a radius that is dependent on the amount of angular momentum in the system. It is from this disk of material that the neutron star may accrete enough matter to power the X-ray pulsations seen. In order to explore this possibility deep infrared imaging observations have been carried out from UKIRT and the NASA IRTF to search for any IR emission from such a disk around three of the neutron stars. The observations reach as deep as J$\sim$20 and in some cases reveal the presence of faint objects in the X-ray error circles. However, the characteristics of the IR colours, and other considerations, strongly suggest that these are not the TZO envelopes around the neutron stars. However, the new stringent limits on any IR companion in these systems leave these objects as enigmatic as ever. The three systems investigated here are: \subsection{1E2259+586} The X-ray pulsar 1E2259+586 lies in the centre of the supernova remnat G109.1-1.0. The most recent estimates on its age make it approximately 3000 years old (Parmar {\it et al.} 1997), though values up to 10000 years have been quoted. G109.1-1.0 is at an estimated distance of about 4 kpc. Fahlman{\it et al.} (1982) demonstrated that the X-ray column densities to both the pulsar and the SNR were similar and hence that these two objects were probably related. However, the latest results from Parmar et al.(1997) indicate that this may not be the case. They find the X-ray column density to the pulsar to be greater than that to the supernova remnant. They suggest that this could be due to the pulsar being at a greater distance than the supernova remnant or to the presence of absorbing material local to the pulsar. However, assuming the same distance to 1E2259+586 as the SNR implies a pulsar X-ray luminosity of $\approx10^{28}$ Watts. Results from {\it EXOSAT}, {\it GINGA} and {\it BeppoSAX} (Hanson et al. 1988, Koyama et al. 1989 and Parmar et al. 1997 respectively) have shown that the 7s period of 1E2259+586 is slowly increasing, but the spindown is too slow to power the observed luminosity. There is no evidence of binary motion, the upper limit on $a_{x} \sin i$ from pulse timing measurements being as low as 30 light-ms (Mereghetti, Israel and Stella, 1998). Early searches for the optical counterpart (Fahlman {\it et al.} 1982; Margon \& Anderson 1983) found several possible candidates, one of which (star D) was tentatively identified with 1E2259+586 by Middleditch, Pennypacker \& Burns (1983) on the basis of IR pulsations. However, more sensitive fast photometric observations by Davies {\it et al.} (1989) failed to find any pulsations, and there would now appear to be little positive evidence associating star D with 1E2259+586. A fainter, multicolour photometric study of all the possible counterparts by Davies \& Coe (1991), together with revised astrometry, updated the list of possible counterparts. Coe \& Jones (1992) presented the first optical spectroscopy of all the candidates brighter than V=23 and re-analysed the Einstein X-ray images of this source taken in 1981 in order to check the position and size of the X-ray error circle. They found that the relocation and reduction in the size of the error circle stemming from the re-processing of these data considerably changes the perspective on possible counterparts to the X-ray pulsar. In addition, a careful study of the images reveals previously unreported features that are possibly associated with jet activity from the pulsar, a possibility originally suggested by Gregory and Fahlman (1980). However, Hurford \& Fesen 1995 report subsequent {\it Rosat} HRI observations which, while confirming the existence of such features, they suggest are just statistical fluctuations in the SNR emission. \subsection{4U0142+62} This persistent source has appeared in most X-ray catalogues and was originally confused (due to poor angular resolution) with the Rosat source RX J0146.9+6121 (=LSI +61 235) - the two sources are only 24 arcminutes apart. Measurements by White et al. (1987) with {\it EXOSAT} reported a 25m modulation from the region (later shown to be coming from LSI +61 235 (Mereghetti, Stella \& De Nile 1993)) but using the unmodulated lower energy data they were able to identify the position of 4U0142+62. The detection of X-ray pulsations at 8.7s was reported by Israel et al. (1994) who concluded that they were most likely coming from the optically unidentified system 4U0142+62 rather than the new {\it Rosat} source. Further measurements by White et al. (1996) using {\it ASCA} confirmed this conclusion and also demonstrated evidence for an excess in the X-ray halo around the source. This could be due to material local to the pulsar along the line of site, possibly in the form of a molecular cloud or perhaps the remains of a TZO envelope. No evidence of any binary motion has been reported from the X-ray signal. The accurate position reported by White et al. (1987) allowed them to locate the X-ray error circles from both {\it EXOSAT} and {\it Einstein} on the sky. As in the case of 1E2259+586 there is no evidence of any obvious optical candidates within the defined regions. They obtained an R band image and set a limit of R$\geq$22.5 for any counterpart. The ASCA observations (White et al. 1996) determined the column to the source to be 8 x 10${^{21}}$ cm${^{-2}}$ which in turn allows them to deduce an optical extinction of A$_{v}$$\sim$4.7. The distance to the source is not well defined but probably lies in the range 0.5--2.0 kpc. As in the case of 1E2259+586 a molecular cloud lies near to, or in front of 4U0142+62 (Mereghetti \& Stella 1995) and may be affecting the column measurement. \subsection{RX J1838.4--0301} This source was originally reported by Schwentker (1994) as 5.5s pulsar possibly associated with a SNR at a distance of 4kpc. Interestingly, in this case the X-ray positional error circle was reported to include a $\sim$14 mag optical candidate. Subsequently Mereghetti et al. (1997) reanalysed the same {\it Rosat} data and reported are different interpretation of the data. They obtained an optical spectrum of the star in the error circle and concluded that this object was simply a K5 main sequence star and that the X-ray emission originated from its corona . | The first deep infrared and optical images have been presented of three proposed remnants of Thorne-$\dot{\rm Z}$ytkow Objects (T$\dot{\rm Z}$O). In particular, the IR data go several infrared magnitudes deeper than previous observations and in at least one case, 1E2259+586, reveal the existence of weak objects within the X-ray error circle. It is argued, however, that in that case none of the objects seen is likely to be the binary companion to the X-ray source in that region. These data present new limits on any possible IR counterpart to the respective X-ray objects. | 98 | 4 | astro-ph9804334_arXiv.txt |
9804 | astro-ph9804272_arXiv.txt | Recently Alard\markcite{alard1996} proposed to detect the shift of a star's image centroid, $\delta x$, as a method to identify the lensed source among blended stars. Goldberg \& Wo\'zniak\markcite{goldberg1997} actually applied this method to the OGLE-1 database and found that 7 out of 15 events showed significant centroid shifts of $\delta x \gtrsim 0.2$ arcsec. The amount of centroid shift has been estimated theoretically by Goldberg.\markcite{goldberg1997} However, he treated the problem in general and did not apply it to a particular survey or field, and thus based his estimates on simple toy model luminosity functions (i.e., power laws). In this paper, we construct the expected distribution of $\delta x$ for Galactic bulge events by using the precise stellar LF observed by Holtzman et al.\markcite{holtzman1998} using HST. Their LF is complete up to $M_I\sim 9.0$ ($M_V\sim 12$), corresponding to faint M-type stars. In our analysis we find that regular blending cannot produce a large fraction of events with measurable centroid shifts. By contrast, a significant fraction of events would have measurable centroid shifts if they are affected by amplification-bias blending. Therefore, Goldberg \& Wo\'zniak's measurements of large centroid shifts for a large fraction of microlensing events confirms the prediction of Han and Alard that a large fraction of Galactic bulge events are affected by amplification-bias blending. | Experiments to detect Massive Astronomical Compact Halo Objects (MACHOs) by monitoring the light variations of stars undergoing gravitational microlensing events are being carried out by several groups (MACHO, Alcock et al.\ 1997a; EROS, Ansari et al.\ 1996; OGLE, Udalski et al.\ 1997; DUO, Alard \& Guibert \ 1997).\markcite{alcock1997a, ansari1996, udalski1997, alard1997b} Since the lensing probability for a single source star is very low, these searches are being conducted toward very dense star fields such as the Large Magellanic Cloud and Galactic bulge. While searches towards these crowded fields result in an increased event rate, it also implies that many of the observed light curves include light from unresolved stars that are not being lensed: the blending problem. Depending on the source for the blended light, blending can be classified into several types. The first type, ``regular blending'', occurs when a bright source star registered on the template plate is lensed and its flux is blended with the light from numerous faint unresolved stars below the detection limit imposed by crowding. Regular blending affects the results of lensing experiments in various ways. First, it makes the measured event timescale shorter than the true one. Since the lens mass scales as $M \propto t_{\rm E}^{2}$, the lens mass determined from the measured timescales will be underestimated and the lens population will be misinterpreted. In addition, since the optical depth is directly proportional to the summation of event timescales, i.e., $\tau \propto \sum_i t_{{\rm E},i}$, the Galactic MACHO fraction determined from the optical depth without a proper blending correction is subject to great uncertainty (Di Stefano \& Esin 1995).\markcite{distefano1995} The second type of blending occurs if the source for the blended light is the lens itself: ``lens blending'' (Kamionkowski 1995; Buchalter, Kamionkowski, \& Rich 1996; Buchalter \& Kamionkowski 1997; Alard 1997).\markcite{kamionkowski1995, bucjalter1996, buchalter1997, alrad1997} Besides the effects of regular blending, lens blending has several additional effects on the result of lensing experiments. First, because detecting events due to lenses close to the observer is comparatively more difficult than detecting events produced by lenses near the source, lens blending makes the optical depth depend on the geometry of the lens system. As a result, the matter distribution derived from the optical depth distribution deviates from its true one. Secondly, lens blending causes the lensing optical depth to depend on the lens mass function since more massive stars, which contribute more to the total optical depth, tend to be brighter, resulting in a larger blending effect (Nemiroff 1997; Han 1998).\markcite{nemiroff1997, han1998} Finally, ``amplification bias'' blending occurs when one of several faint stars in the seeing disk below the detection limit is lensed, and the flux of the lensed star is associated with the flux from other stars in the integrated seeing disk (Nemiroff 1997).\markcite{nemiroff997} In current experiments, photometry is carried out by comparing template images, in which only very bright stars are resolved and registered, with a series of images taken of the same field. The result is that in amplification-biased events the brightest star appears to be the source because the lensed star is too faint to be resolved. To be detected, the amplification-biased event must be highly amplified to overcome the high threshold flux from the brighter star. Therefore, the mean detection probability to detect events for each source star will be very low. However, if these faint stars comprise a significant fraction of the total number of stars, a considerable fraction of events might be amplification-biased (Han 1997b; Alard 1997a).\markcite{han1997b, alard1997a} Moreover, due to the large amount of blended light from a bright star, the effects of blending for these events would be much more significant than those caused by other blending types. There have been several methods proposed to correct for the blending problem. The first method is to introduce an additional lensing parameter representing the residual flux from unresolved faint stars into the light curve fitting process. However, this method suffers from large uncertainties in the derived lensing parameters as a result of parameter degeneracies (Wo\'zniak \& Paczy\'nski 1997).\markcite{wozniak1997} Early-warning systems (for MACHO: Pratt et al.\ 1996; for OGLE: Udalski et al.\ 1994; for PLANET: Albrow et al.\ 1995)\markcite{pratt1996, udalski1994, albrow1995} allow one to construct lensing light curves with high time resolution and small photometric errors, enabling one to detect small color shifts caused by blending (Buchalter et al.\ 1996)\markcite{buchalter996}. However, due to the narrow distribution of colors for Galactic bulge stars, the expected color shifts are small. Han (1997b)\markcite{han1997} proposed to use the Hubble Space Telescope (HST) to provide blending corrections. By using the high resolving power of HST combined with color information from ground-based observations, one can identify the lensed source star in the blended seeing disk, thus the uncertainty in the derived timescale can be significantly reduced. However, this method requires costly HST times. One can also correct the blending effect statistically if the luminosity function(LF) of stars well below the detection limit can be constructed (Alcock et al.\ 1997b)\markcite{alcock1996}, but in this case we lose information about individual events. Recently Alard (1996)\markcite{alard1996} proposed to detect the shift of a star's image centroid, $\delta x$, as a method to identify the lensed source among blended stars. Goldberg \& Wo\'zniak (1997)\markcite{goldberg1997} actually applied this method to the OGLE-1 database and found that 7 out of 15 events showed significant centroid shifts of $\delta x \gtrsim 0.2$ arcsec. The amount of centroid shift has been estimated theoretically by Goldberg (1997).\markcite{goldberg1997} However, he treated the problem in general and did not apply it to a particular survey or field, and thus based his estimates on simple toy model luminosity functions (i.e., power laws). In this paper, we construct the expected distribution of $\delta x$ for Galactic bulge events by using the precise stellar LF observed by Holtzman et al.\ (1998)\markcite{holtzman1998} using HST. Their LF is complete up to $M_I\sim 9.0$ ($M_V\sim 12$), corresponding to faint M-type stars. In our analysis we find that regular blending cannot produce a large fraction of events with measurable centroid shifts. By contrast, a significant fraction of events would have measurable centroid shifts if they are affected by amplification-bias blending. Therefore, Goldberg \& Wo\'zniak's measurements of large centroid shifts for a large fraction of microlensing events confirms the prediction of Han (1997a) and Alard (1997) that a large fraction of Galactic bulge events are affected by amplification-bias blending. | The fact that events affected by amplification-bias blending produce large $\delta x$ while the centroid shifts caused by regular blending are small can be understood analytically in the following way. To produce centroid shifts large enough to be measured, events should satisfy two conditions. First, the event should be highly amplified. For a very low amplification event ($A_{\rm abs}\sim 1$), the expected centroid shift will be small since $A_{\rm abs}-1 \sim 0$ and thus $\eta \sim 0$ in equation (1). On the other hand, for a very high-amplification event ($A_{\rm abs}\sim \infty$), one finds $\delta x\sim \left\vert\langle {\bf x}\rangle -{\bf x}_{j}\right\vert$ since the factor $\eta$ approaches unity. However, not all events that satisfy the first condition produce large centroid shifts. The second condition is that the lensed star should be one of the faint stars in the blended seeing disk. If the lensed star is the brightest one in the effective seeing disk and its flux dominates the flux over those from other faint blended stars, i.e., $\kappa\sim 1$, the position of the CL before gravitational amplification will be very close to that of the lensed star, resulting in a small amount of shift ($\left\vert\langle {\bf x}\rangle-{\bf x}_j\right\vert \sim 0$) since $\sum_{i\neq j} F_{0,i}{\bf x}_i + F_{0,j}{\bf x}_j\sim F_{0,j}{\bf x}_j$ and $\sum_i F_{0,i} \sim F_{0,j}$. On the other hand, if the lensed source is very faint, i.e., $\kappa\sim 0$, the light from the source star has negligible effect on the position of CL, resulting in high possibility of a large centroid shift. In summary, the expected centroid shifts for various extreme cases of amplification and source star brightness are: $$ \cases{ \eta\sim 0\ {\rm and}\ \delta x\sim 0 & for a low amplification event \cr & ($A_{\rm abs}\sim 1$) \cr \delta x\sim \left\vert\langle{\bf x}\rangle-{\bf x}_j\right\vert\sim 0 & for a high-amplification event\cr & with luminous source \cr & ($A_{\rm abs}\sim \infty$, $\kappa\sim 1$) \cr \delta x\sim \left\vert\langle{\bf x}\rangle-{\bf x}_j\right\vert & for a high-amplification event\cr & with faint source \cr & ($A_{\rm abs}\sim \infty$, $\kappa\sim 0$). \cr } \eqno(4) $$ For amplification-biased events, source stars are in general very faint ($\kappa\sim 0$), mostly far below the detection limit. Despite their low luminosities, the fact that they are detected implies that the source stars are highly amplified ($A_{\rm abs}\sim \infty$). Therefore, the conditions for large centroid shifts agree well with those for amplification-biased events. On the other hand, regular blended events do not meet these conditions. First, due to relatively small amount of blended light, regular blended events do not need to be highly amplified for detection ($A_{\rm abs}\sim 1$). Although they can be highly amplified ($A_{\rm abs}\sim \infty$), the dominance of their fluxes ($\kappa\sim 1$) over those from other faint blended stars will result in small $\delta x$. As demonstrated by the large centroid shifts for a significant fraction of events, the effect of amplification-bias blending on the results of lensing experiments is important. However, the methods to correct for the blending effect mentioned in \S\ 1 have various limitations in application. One very simple but very practical method to minimize the effect of amplification-bias blending is to monitor only very bright stars. With increasing reference image brightness, the required amplification for detection becomes higher, resulting in a lower probability of amplification-biased events. To show how the blending effect decreases with increasing reference image brightness, we simulate Galactic bulge events which are expected to be detected for various threshold reference image brightnesses. For each event, we compute the light fraction of the source star $\kappa$ and the timescale decrease factor $\eta$. The distributions of $\kappa$ and $\eta$ are shown in the upper panel of Figure 6. For a given fraction of source star flux, the observed timescale is reduced by $\eta=t_{\rm eff}/t_{\rm E}= \left[ 2(1-A_{\rm min}^{-2})^{-1/2}-2 \right]^{1/2}; \ A_{\rm min} = 0.34/\kappa + 1$. Here events are assumed to be detected as long as they can amplify the reference image flux by more than a factor of 1.34. However, highly blended events will have short $t_{\rm eff}$, resulting in low detectability. Therefore, we correct the distributions by the detection efficiency. We assume that the detection efficiency is linearly proportional to the timescale decrease factor, i.e., $\epsilon \propto \eta$. The efficiency-corrected distributions $f(\kappa)$ and $f(\eta)$ are presented in the middle panels. In the lower panels we present the distributions of the fraction of events with $\kappa \ge \kappa_{\rm lim}$, $1-\int_{0}^{\kappa_{\rm lim}} f(\kappa)d\kappa$, and $\eta \ge \eta_{\rm lim}$, $1-\int_{0}^{\eta} f(\eta')d\eta'$. From these distributions one finds that under the current threshold reference image brightness of $M_I\sim 3$, the fraction of events with little blending effects ($\kappa \gtrsim 0.9$ or $\eta \gtrsim 0.9$) is $\lesssim 10\%$. However, as the brightness of the threshold reference image increases, this fraction gradually increases until it becomes $\sim 80\%$ when only stars brighter than $M_I\sim 0$ are monitored, which corresponds to the brightness of Galactic bulge clump giant stars (Paczy\'nski \& Stanek 1998).\markcite{paczynski1998} The MACHO group (1997a)\markcite{alcock1997a} already applied this method and their optical depth determination is based on clump giant stars. However, by monitoring significantly fewer stars at a decreased event rate, the statistical precision of the lensing experiments will be lowered. A more general solution for the blending correction is provided by the rapidly progressing image subtraction technique (Alard \& Lupton 1997b; Tomaney 1998)\markcite{alard1997c, tomaney1998} which is also being applied to detect microlensing events towards M31 by the Colombia-Vatt group (Crotts \& Tomaney 1996; Tomaney \& Crotts 1996).\markcite{crotts1996, tomaney1996} | 98 | 4 | astro-ph9804272_arXiv.txt |
9804 | astro-ph9804050_arXiv.txt | We investigate if the gamma ray halo, for which recent evidence has been found in EGRET data, can be explained by neutralino annihilations in a clumpy halo. We find that the measured excess gamma ray flux can be explained through a moderate amount of clumping in the halo. Moreover, the required amount of clumping implies also a measureable excess of antiprotons at low energies, for which there is support from recent measurements by the BESS collaboration. The predicted antiproton fluxes resulting from neutralino annihilations in a clumpy halo are high enough to give an excess over cosmic-ray produced antiprotons also at moderately high energies (above a few GeV). This prediction, as well as that of one or two sharp gamma lines coming from annihilations into $\gamma\gamma$ or $Z\gamma$ can be tested in upcoming space-borne experiments like AMS and GLAST. | 98 | 4 | astro-ph9804050_arXiv.txt |
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9804 | astro-ph9804099_arXiv.txt | We investigate the nature of stellar populations of major galaxy mergers between late-type spirals considerably abundant in interstellar medium by performing numerical simulations designed to solve both the dynamical and chemical evolution in a self-consistent manner. We particularly consider that the star formation history of galaxy mergers is a crucial determinant for the nature of stellar populations of merger remnants, and therefore investigate how the difference in star formation history between galaxy mergers affects the chemical evolution of galaxy mergers. We found that the rapidity of star formation, which is defined as the ratio of the dynamical time-scale to the time-scale of gas consumption by star formation, is the most important determinant for a number of fundamental characteristics of stellar populations of merger remnants. Main results obtained in this study are the following five. (1) A galaxy merger with more rapid star formation becomes elliptical with larger mean metallicity. This is primarily because in the merger with more rapid star formation, a smaller amount of metal-enriched gas is tidally stripped away during merging and consequently a larger amount of the gas can be converted to stellar component. This result demonstrates that the origin of the color-magnitude relation of elliptical galaxies can be closely associated with the details of merging dynamics which depends on the rapidity of star formation in galaxy mergers. (2) Negative metallicity gradient fitted reasonably well by power-low can be reproduced by dissipative galaxy mergers with star formation. The magnitude of metallicity gradient is larger for an elliptical galaxy formed by galaxy merging with less rapid star formation. (3) Absolute magnitude of metallicity gradient correlates with that of age gradient in galaxy mergers in the sence that a merger remnant with steeper negative metallicity gradient is more likely to show steeper age gradient. (4) The outer part of stellar populations is both older and less metal-enriched than nuclei in an elliptical galaxy formed by galaxy merging with less rapid star formation. Moreover, the metallicity of the outer part of gaseous component for some models with less rapid star formation is appreciably smaller than that of stellar one. This result implies that the origin of metal-poor hot gaseous $X$-ray halo in real elliptical galaxies can be essentially ascribed to the dynamics of dissipative galaxy merging. (5) Irrespectively of the rapidity of star formation, the epoch of galaxy merging affects both the mean stellar metallicity and mean stellar age of merger remnants: Later galaxy mergers are more likely to become ellipticals with both younger and more metal-enriched stellar populations. This result reflects the fact that in the later mergers, a larger amount of more metal-enriched interstellar gas is preferentially converted into younger stars in the later star formation triggered by galaxy merging. These five results clearly demonstrate that even the chemical evolution of elliptical galaxies can be strongly affected by the details of dynamical evolution of galaxy merging, which is furthermore determined by the rapidity of star formation of galaxy mergers. In particular, tidal stripping of interstellar gas and total amount of gaseous dissipation during galaxy merging are demonstrated to play a vital role in determining a number of chemical properties of merger remnants. Based upon these results, we adopt a specific assumption of the luminosity dependence of the rapidity of star formation and thereby discuss how successfully the present merger model can reproduce a number of fundamental chemical, photometric, and spectroscopic characteristics of elliptical galaxies. | Elliptical galaxies have been generally considered to be old, coeval and homogeneous systems passively evolving after the single initial burst of star formation associated with dissipative galaxy formation. This classical picture of elliptical galaxy formation appears to have been supported by the considerably tight color-magnitude relation of elliptical galaxies ( Bower, Lucey, \& Ellis 1992; Ellis et al. 1997) and by relatively smaller redshift evolution of photometric properties of elliptical galaxies (Arag$\rm \acute{o}$n-Salamanca et al. 1993; Franx \& van Dokkum 1996). A growing number of recent observational results, however, shed a strong doubt on this long-standing view of elliptical galaxy formation, and suggest that there is great variety of star formation history between elliptical galaxies, such as the epoch of major star formation, the duration and efficiency of star formation (Worthey, Faber, \& Gonzalez 1992; Matteuchi 1994; Faber et al. 1995; Bender 1996; Worthey, Trager, \& Faber 1996). This tendency that elliptical galaxies show diversity in star formation history and nevertheless can actually keep the tightness of the color-magnitude relation is considered to be quite mysterious and thus to provide any theoretical models with a valuable insight on the elliptical galaxy formation. Such kind of mysterious nature observed in elliptical galaxies is demonstrated to hold equally for the dynamical and kinematical properties of elliptical galaxies. For example, considerably small thickness of the fundamental plane of elliptical galaxies implies a rather smaller range of admitted dynamical state of the galaxies (Djorgovski \& Davis 1987; Dressler et al. 1987 ; Djorgovski, Pahre, \& de Carvalho 1996) whereas the morphological dichotomy between boxy-disky elliptical galaxies (Kormendy \& Bender 1996) and the projected density profile systematically departing from de Vaucouleurs $R^{1/4}$ law (Caon, Capaccioli, \& D'Onofrio 1993) show a great variety of major orbit families consisting the galaxies. These fundamental characteristics that elliptical galaxies show both diversity and uniformity in their chemical, photometric and dynamical properties have imposed some stringent but valuable constraints on any theoretical models of elliptical galaxy formation. What is the most vital in challenging the origin of elliptical galaxy formation in this kind of situation is to investigate whether or not both the chemical and photometric properties and dynamical and kinematical ones can be reproduced successfully by a specific model of galaxy formation in a reasonably self-consistent manner. The previous theoretical models addressing this important issue on elliptical galaxy formation are divided basically into two categories: The dissipative galactic collapse model (e.g., Larson 1976; Carlberg 1984) and the galaxy merger model (e.g., Toomre \& Toomre 1972). As is suggested by Kormendy \& Sanders (1992), these two dominant and apparently competing scenarios for elliptical galaxy formation are now converging, thus it would be crucial to construct one more realistic and sophisticated model of elliptical galaxy formation. Although there are a large number of important studies exploring the origin of elliptical galaxy formation along the dissipative collapse scenario, especially in the context of the nature of stellar populations (e.g., Arimoto \& Yoshii 1987), we here restrict ourselves to the merger scenario of elliptical galaxy formation. Recent extensive studies of merger models of elliptical galaxy formation, mostly based upon numerical simulations, $appear$ to have succeeded in resolving most of the outstanding problems related to dynamical and kinematical properties of elliptical galaxies, such as the phase space density (Ostriker 1980; Carlberg 1986) and kinematical misalignment (Franx, Illingworth, \& de Zeeuw 1991; Barnes 1992), by invoking the inclusion of bulge component, gaseous dissipation, and multiplicity of galaxy merging (Hernquist, Spergel, \& Heyl 1993; Weil \& Hernquist 1996; Barnes \& Hernquist 1996). Although it would be safe to say that the galaxy merging between two late-type spirals is one of the most promising candidates explaining more clearly the origin of elliptical galaxies at least in the context of the $dynamical$ $and$ $kinematical$ $properties$, however, there still remain a number of unresolved and apparently serious problems concerning the merger model (e.g., van den Bergh 1995). One of the most crucial problems among these is on whether the fundamental characteristics of stellar populations of elliptical galaxies can be reproduced reasonably well by galaxy merging between two late-types spirals. Surprisingly, there are only a few works addressing this critical issue for the merger model, probably because it is considered to be rather difficult to solve the chemical evolution of galaxy mergers in which a number of competing physical processes are expected to affect strongly the chemical evolution of galaxy mergers. White (1980) and Mihos \& Hernquist (1994) found that the stellar populations of progenitor disks are not mixed so well even by the violent relaxation during galaxy merging and consequently the metallicity gradient of progenitor disks is not so drastically washed out. The metallicity gradient of merger remnant is furthermore found not to be fitted by power law observed in elliptical galaxies (Mihos \& Hernquist 1994). Schweizer \& Seitzer (1992) discussed whether or not the bluer integrated $UBV$ color of elliptical galaxies with morphologically fine structure can be explained by secondary starburst induced by major disk-disk galaxy mergers. Kauffmann \& Charlot (1997) construct a semi-analytic model of elliptical galaxy formation, which is based upon the hierarchical clustering in CDM universe and includes rather simple chemical enrichment process, and thereby demonstrate that the origin of the color-magnitude relation of elliptical galaxies can be reproduced successfully even in the CDM model of galaxy formation (See also Baugh, Cole \& Frenk 1996.). Thus, since there are only a few works addressing chemical and photometric properties for the merger model, it is essential for the merger model to investigate more throughly the fundamental chemical and photometric properties of merger remnants, including the origin of color-magnitude relation (Faber 1973; Visvanathan \& Sandage 1977), age and metallicity gradient (Peletier et al. 1990; Davies et al 1991), $\rm Mg_{2} - \sigma$ relation (Burstein et al. 1988), age-metal-conspiracy in stellar populations (Faber et al. 1995; Worthey et al. 1996), luminosity dependence of the line ratio [Mg/Fe] (Worthey et al. 1992), metal-poor gaseous $X$-ray halo (Matsumoto et al. 1997), and the substantially metal-enriched galactic nuclei at higher redshift (Hamann \& Ferland 1993). What should be recognized foremost in investigating the nature of stellar populations in merger remnants is that a glowing number of observational results have been accumulated which suggest the relatively earlier formation of elliptical galaxies. Tightness of the color-magnitude relation in the cluster of galaxies (Bower et al. 1992, Ellis et al 1996), relatively smaller photometric evolution of cluster ellipticals (Arag$\rm \acute{o}$n-Salamanca et al. 1993), and the redshift evolution of the fundamental plane (Franx \& van Dokkum 1996) all suggest the $typical$ formation epoch of elliptical galaxies is earlier than 2 in redshift. Furthermore, as is suggested by Kormendy \& Sanders (1992), the fact that no galaxy in the $K$-band survey of Cowie et al. (1994) shows the global color resembling that of the Arp 220, which is considered to be ongoing mergers and forming ellipticals, implies that the formation epoch of elliptical galaxies should be earlier than 1.0 in redshift. Silva \& Bothun (1997) revealed that the fraction of mass of stellar populations with intermediate age to total mass in elliptical galaxies with morphologically fine structure is less than 15 percent. These results imply that if elliptical galaxies are formed by galaxy merging, the epoch of galaxy merging should be relatively earlier and furthermore that the precursor disks of galaxy mergers may be extremely abundant in interstellar medium compared with the present spirals. Recent high quality imaging using $Hubble$ $Space$ $Telescope$ ($HST$) has revealed that a larger number of galaxies at faint magnitude are interacting/merging galaxies (e.g., van den Bergh et al. 1996), indicating furthermore that the potential candidate for elliptical galaxies formed by galaxy merging are ubiquitous in higher redshift universe. Hence it is quite reasonable and essential to study the nature of stellar populations of higher redshift galaxy mergers between disk galaxies with the gas mass fraction larger than 0.2, which is a typical value of the present late-type spirals, and thereby to confirm whether or not elliptical galaxies can be formed $actually$ by galaxy merging. The purpose of this paper is to explore the nature of the stellar populations of a gas-rich disk merger which is considered to be occurred the most frequently in the high redshift universe. We particularly investigate how successfully galaxy mergers between gas-rich spirals can reproduce a number of fundamental chemical, photometric, and spectroscopic properties of elliptical galaxies. The layout of this paper is as follows. In \S 2, we summarize numerical models used in the present study and describe in detail methods for analyzing the stellar populations produced by dissipative galaxy mergers with star formation. In \S 3, we demonstrate how a number of fundamental characteristics of stellar populations in merger remnants are affected by the star formation history of dissipative galaxy merging. In \S 4, we discuss how successfully the present merger model can reproduce a number of observational results concerning the chemical, photometric, and spectroscopic properties of elliptical galaxies. In this section, we also point out the advantages and disadvantages of galaxy mergers in explaining both the chemical, photometric, and spectroscopic properties and dynamical and kinematical ones in real elliptical galaxies. The conclusions of the present study are given in \S 5. | Main results obtained in the present study are summarized as follows. (1) Galaxy mergers with more rapid star formation become ellipticals with larger mean stellar metallicity, primarily because in the mergers with more rapid gas consumption, a smaller amount of metal-enriched gas is tidally stripped away during merging and consequently a larger amount of the gas can be converted into stellar component. This result is demonstrated not to depend so strongly on the other parameters such as the orbit configuration of galaxy merging and multiplicity of the mergers. These results suggest that the origin of the color-magnitude relation of elliptical galaxies can be closely associated with the details of merging dynamics which depends on the rapidity of star formation (thus on the galactic luminosity) in galaxy mergers. (2) Negative metallicity gradient fitted reasonably well by power-low can be reproduced by dissipative galaxy mergers with star formation, which is in good agreement with the recent observational results of elliptical galaxies. The absolute magnitude of metallicity gradient in each merger remnant depends on the orbit configuration of each galaxy merging, suggesting that the observed dispersion in the absolute magnitude of metallicity gradient for a given luminosity range of elliptical galaxies reflects the diversity in the orbit configuration of galaxy merging. (3) Absolute magnitude of metallicity gradient correlates with that of age gradient in a merger in the sence that a merger remnant with steeper negative metallicity gradient is more likely to show steeper age gradient. This result reflects the fact that the degree of violent relaxation and gaseous dissipation during merging strongly affect both the age gradient and metallicity one. (4) The outer part of stellar populations is both older and less metal-enriched than nuclei in an elliptical galaxy formed by galaxy merging with less rapid star formation. Moreover galaxy mergers with less rapid star formation are more likely to become ellipticals with metal-poor gaseous halo. This result suggests that the formation of metal-poor $X$-ray halo actually observed in elliptical galaxies can be essentially ascribed to the dissipative galaxy merging between late-type spirals, and furthermore provides a clue to a solution for the iron abundance discrepancy problem in elliptical galaxies. (5) The epoch of galaxy merging affects both the mean stellar metallicity and the mean stellar age in merger remnants: Later galaxy mergers become ellipticals with both younger and more metal-enriched stellar populations. This result suggests that the origin of Worthey's 3/2 rule (Worthey et al. 1996), which is invoked in maintaining the tightness of the color-magnitude relation of elliptical galaxies, can be understood in terms of the difference in the epoch of galaxy formation and transformation, that is, the epoch of galaxy merging, between elliptical galaxies. (6) Luminosity dependence of chemical, photometric, and spectroscopic properties in merger remnants, which is derived by adopting a specific assumption on the luminosity dependence of the rapidity of star formation of galaxy mergers, does not match so reasonable well with that observed in real elliptical galaxies. This result implies that other fundamental physical processes expected to be dependent on the galactic luminosity should be incorporated into the present merger model for more successful comparison with observational trends of luminosity-dependent chemical, photometric, and spectroscopic properties of elliptical galaxies. (7) As is described in the above (1) - (6), the details of gas dynamics of galaxy merging, in particular, the tidal stripping of metal-enriched interstellar gas and the degree of gaseous dissipation during merging, both of which depend on the star formation history of galaxy mergers, are demonstrated to determine even the chemical and photometric properties of merger remnants. These results can not be obtained until both the chemical and dynamical evolution during galaxy merging are solved numerically in a reasonably self-consistent way. | 98 | 4 | astro-ph9804099_arXiv.txt |
9804 | astro-ph9804320_arXiv.txt | Physical properties of the atomic gas in spiral galaxies are briefly considered. Although both Warm (WNM, 10$^4$~K) and Cool (CNM, $\sim$~100 K) atomic phases coexist in many environments, the dominant mass contribution within a galaxy's star-forming disk (R$_{25}$) is that of the CNM. Mass fractions of 60 to 90\% are found within R$_{25}$. The CNM is concentrated within moderately opaque filaments with a face-on surface covering factor of about 15\%. The kinetic temperature of the CNM increases systematically with galactocentric radius, from some 50 to 200~K, as expected for a radially declining thermal pressure in the galaxy mid-plane. Galaxies of different Hubble type form a nested distribution in T$_K$(R), apparently due to the mean differences in pressure which result from the different stellar and gas surface densities. The opaque CNM disappears abruptly near R$_{25}$, where the low thermal pressure can no longer support the condensed atomic phase. The CNM is apparently a prerequisite for star formation. Although difficult to prove, all indications are that at least the outer disk and possibly the inter-arm atomic gas are in the form of WNM, which accounts for about 50\% of the global total. Median line profiles of the CNM display an extremely narrow line core (FWHM $\sim$~6~km~s$^{-1}$) together with broad Lorentzian wings (FWHM $\sim$~30~km~s$^{-1}$). The line core is consistent with only opacity broadening of a thermal profile. The spatial distribution of CNM linewidths is not random, but instead is extremely rich in structure. High linewidths occur in distinct shell-like structures with diameter of 100's of pc to kpc's, which show some correlation with diffuse H$\alpha$ shells. The primary source of ``turbulent'' linewidth in the atomic ISM appears to be organized motions due to localized energy injection on a scale of a few 100 pc. | Over the years a succession of eminent authors has considered the question of which physical processes determine the temperature of neutral atomic gas and which timescales are required to achieve (a local) thermodynamic equilibrium. This began with Field, Goldsmith and Habing (1969) and continued with Draine (1978), Shull and Woods (1985) and most recently Wolfire {\it et al.} (1995). While some of the processes involved are quite straightforward, others, like the photoelectric emission from dust grains, have had to be substantially updated to reflect our growing knowledge of the properties and abundance of interstellar dust. The consensus which has emerged is that heating is in fact dominated by the photoelectric heating from small dust grains over a wide range of conditions and environments. Cooling, on the other hand, is regulated primarily by emission in the [CII] 158~$\mu$m fine structure line at densities in excess of about 1~cm$^{-3}$ and by Ly$\alpha$ emission at lower densities. The relative importance of the various mechanisms is nicely illustrated in Fig.~3 of Wolfire {\it et al.} (1995). In the same figure can be seen the characteristic phase diagram for atomic gas. Two thermodynamically stable phases are found. The first is the so-called Warm Neutral Medium (WNM) which predominates at low densities and pressures, and has a kinetic temperature that rises toward lower pressures from a value of perhaps 5000~K to 10$^4$~K, with an accompanying increase in ionization fraction, $x$~$\sim$~0.1 to 0.9. The second is the Cool Neutral Medium (CNM) which is primarily found at high densities and pressures, with a kinetic temperature that decreases towards higher pressures from a value of some 200~K to perhaps as low as 20~K. Over some range of intermediate pressures thought to be typical of the interstellar medium in the local neighbourhood of the Galaxy (P/k$\sim$ 2000~K~cm$^{-3}$), the two phases can coexist in pressure equilibrium. However, given the strong gradient in the mid-plane thermal pressure which is likely to follow from the combination of a radial exponential stellar disk and a similarly declining gas surface density distribution, we can safely predict that the inner regions of disk galaxies will have predominantly rather cool CNM, while the outer regions of galaxies will eventually be dominated by the WNM. At intermediate radii we might also expect to see a radial increase in the CNM kinetic temperature in response to the declining mid-plane pressure. | 98 | 4 | astro-ph9804320_arXiv.txt |
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9804 | astro-ph9804116_arXiv.txt | Multidimensional cosmologies allow for variations of fundamental physical constants over the course of cosmological evolution, and different versions of the theories predict different time dependences. In particular, such variations could manifest themselves as changes of the proton-to-electron mass ratio $\mu=m_{\rm p}/m_{\rm e}$ over the period of $\sim 10^{10}$ yr since the moment of formation of high-redshift QSO spectra. Here we analyze a new, high-resolution spectrum of the $z = 2.81080$ molecular hydrogen absorption system toward the QSO PKS 0528--250 to derive a new observational constraint to the time-averaged variation rate of the proton-to-electron mass ratio. We find $| \dot{\mu}/ \mu | < 1.5 \times 10^{-14}$ yr$^{-1}$, which is much tighter than previously measured limits. | The possibility of the variability of fundamental physical constants was first put forward by Dirac (1937) in the course of his discussion with Milne (1937). Later it was considered by Teller (1948), Gamow (1967), Dyson (1972) and other physicists. Interest in the problem increased greatly during the last decade, due to new developments in the Kaluza--Klein and supergravity models of unification of all the physical interactions. Chodos \& Detweiler (1980), Freund (1982), Marciano (1984), and Maeda (1988) discussed possibilities of including these multidimensional theories into the cosmological scenario of the expanding Universe and found that the low-energy limits to the fundamental constants might vary over the cosmological time. Variations of the coupling constants of strong and electroweak interactions might then cause the masses of elementary particles to change. Note that an increase of the proton mass by 0.08\% would lead to transformation of protons into neutrons (by electron capture), resulting in destruction of atoms in the Universe. As demonstrated by Kolb, Perry, \& Walker (1987) and Barrow (1987), observational bounds on the time evolution of extra spatial dimensions in the Kaluza--Klein and superstring theories can be obtained from limits on possible variations of the coupling constants. Damour \& Polyakov (1994) have developed a modern version of the string theory which assumes cosmological variations of the coupling constants and hadron-to-electron mass ratios. Therefore the parameters of the theory can be restricted by testing cosmological changes of these ratios. The present value of the proton-to-electron mass ratio is $\mu=1836.1526645\,(57)$ (CODATA, 1997). Obviously, any significant variation of this parameter over a small time interval is excluded, but such variation over the cosmological time $\sim 1.5\times 10^{10}$ yr remains a possibility. This possibility can be checked by analyzing spectra of high-redshift QSOs. The first analysis of this kind has been performed by Pagel (1977), who obtained a restriction $|\dot{\mu}/\mu|< 5\times 10^{-11}{\rm ~yr}^{-1}$ on the variation rate of $\mu$ by comparison of wavelengths of H\,I and heavy-ion absorption lines, as proposed by Thompson (1975). This technique, however, could not provide a fully conclusive result, since the heavy elements and hydrogen usually belong to different interstellar clouds, moving with different radial velocities. In this paper we use another technique, based on an analysis of H$_2$ absorption lines only. One object suitable for such analysis is the $z = 2.811$ absorption system toward PKS 0528--250, in which Levshakov \& Varshalovich (1985) identified molecular hydrogen absorption lines based on a spectrum obtained by Morton et al.\ (1980). Foltz, Chaffee, \& Black (1988) have presented a limit to possible variation of $\mu$ based on their observations of PKS 0528--250. Their analysis did not, however, take into account wavelength-to-mass sensitivity coefficients, hence their result appeared to be not well grounded. Subsequently the spectrum of Foltz, Chaffee, \& Black (1988) was reappraised by Varshalovich \& Levshakov (1993), who obtained $|\Delta\mu/\mu| < 0.005$ at the redshift $z=2.811$, and by Varshalovich \& Potekhin (1995), who obtained $|\Delta\mu/\mu| < 0.002$ at the $2\sigma$ significance level. (Here $\Delta\mu/\mu$ is the fractional variation of $\mu$.) More recently, Cowie \& Songaila (1995) used a new spectrum of PKS 0528$-$250 obtained with the Keck telescope to arrive at the 95\% confidence interval $-5.5 \times 10^{-4} < \Delta\mu/\mu < 7 \times10^{-4}$. Here we present a profile fitting analysis of a new, high-resolution spectrum of PKS 0528$-$250, obtained in November 1991 with the Cerro-Tololo Inter-American Observatory (CTIO) 4 m telescope. We have calculated the wavelength-to-mass sensitivity coefficients for a larger number of spectral lines and employed them in the analysis, which yields the strongest observational constraint yet to possible $\mu$ variation over the cosmological time scale (eq.~[\ref{limits}] below). | We have obtained a constraint to the variation rate of the proton-to-electron mass ratio $\mu$. Two fitting procedures have been used, one of which simultaneously takes into account all observed spectral regions and transitions, while the other is applied to each spectral feature separately. The two techniques, applied to two different sets of spectral intervals, have resulted in similar upper bounds on $\Delta\mu/\mu$, at the level $\,\sim2\times10^{-4}$. The obtained restriction on $\dot{\mu}/\mu$ (\ref{constraint}) is by an order of magnitude more stringent than the limit set previously by Varshalovich \& Potekhin (1995), who used a spectrum with a lower spectral resolution. Moreover, it is much more restrictive than the estimate of Cowie \& Songaila (1995), based on high-resolution Keck telescope observations. There are two reasons for the higher accuracy of the present estimate. First, our fitting procedure simultaneously takes into account all observed spectral regions and transitions. This is particularly important because many of the transitions are blended, even at the spectral resolution of the spectrum used by Cowie \& Songaila (1995). A separate analysis of spectral lines leads to larger statistical errors, as we have shown explicitly in Section 4.2. Second, we include a larger number of transitions between excited states of the H$_2$ molecule (83 spectral lines, compared with 19 lines used by Cowie \& Songaila), many of which have higher wavelength-to-mass sensitivity coefficients $K_i$. The larger interval of $K_i$ values results in a higher sensitivity to possible mass ratio deviations. The method used here to determine the variation rate of $\mu$ could be formally less sensitive than the one based on an analysis of relative abundances of chemical elements produced in the primordial nucleosynthesis (Kolb et al., 1986). However, the latter method is very indirect since it depends on a physical model which includes a number of additional assumptions. Therefore the present method seems to be more reliable. Quite recently, Wiklind \& Combes (1997) used a similar method (following Varshalovich \& Potekhin, 1996) in order to infer limits on time variability of masses of molecules CO, HCN, HNC and the molecular ion HCO$^+$ from high-resolution radio observations of rotational lines in spectra of a few low-redshift ($z<1$) quasars. The result reported in this paper constrains the mass of the H$_2$ molecule, and thus the proton mass, at much larger $z$. These constraints may be used for checking the multidimensional cosmological models which predict time-dependences of fundamental physical constants. The described method of the calculation of the sensitivity coefficients can also be used for analyzing any other high-redshift molecular clouds, which may be found in future observations. | 98 | 4 | astro-ph9804116_arXiv.txt |
9804 | hep-ph9804291_arXiv.txt | Mixed dark matter scenario can reconcile the COBE data and the observed large scale structure. So far the massive neutrino with a mass of a few eV has been the only discussed candidate for the hot dark matter component. We point out that the hadronic axion in the so-called hadronic axion window, $f_{a} \sim 10^{6}$~GeV, is a perfect candidate as hot dark matter within the mixed dark matter scenario. The current limits on the hadronic axion are summarized. The most promising methods to verify the hadronic axion in this window are the resonant absorption of almost-monochromatic solar axions from M1 transition of the thermally excited $^{57}$Fe in the Sun, and the observation of the ``axion burst'' in water \v{C}erenkov detectors from another supernova. | The cold dark matter (CDM) dominated universe with scale-invariant primordial density fluctuation has been the standard theory of structure formation. After COBE has found the finite density fluctuation in the cosmic microwave background radiation (CMBR), the standard CDM scenario was found to give too much power on smaller scales. Many modifications to the standard CDM scenario were proposed which solve the discrepancy: by introducing a small Hot Dark Matter (HDM) component~\cite{aph9707285}, by ``tilting'' the primordial density fluctuation spectrum~\cite{tilt}, by assuming a finite cosmological constant~\cite{Lambda-CDM}, or by introducing particles (such as $\nu_{\tau}$) whose decay changes the time of radiation-matter equality~\cite{cdm+mnu}. At this point, there is no clear winner among these possibilities.\footnote{However, a large ``tilt'' is difficult to obtain in many inflationary models. $\tau$CDM can be tested well by $B$-factory experiments in the near future~\cite{hph9709411}. The recent data from high-redshift supernovae prefer $\Lambda$CDM \cite{Saul}, but the possible evolution of supernovae needs to be excluded by more systematic comparison between nearby and high-$z$ supernovae.} In this letter, we revisit the mixed dark matter (MDM) scenario from the particle physics point of view. This scenario has attracted strong interests because there has been a natural candidate for the HDM component: massive neutrino(s). A neutrino with a mass of a few eV can naturally contribute to a significant fraction of the current universe. However, it has not been easy to incorporate the HDM together with other neutrino ``anomalies,'' unless all three generation neutrinos (possibly together with a sterile neutrino) are almost degenerate, and their small mass splittings explain various ``anomalies.'' Such a scenario may be viewed as fine-tuned. Especially, the atmospheric neutrino anomaly is quite significant statistically now thanks to the SuperKamiokande experiment, which suggests the mass squared difference of $\Delta m^{2} = 10^{-3} - 10^{-2}~{\rm eV}^{2}$ between the muon and tau neutrinos. If we view the situation from the familiar hierarchical fermion mass matrices, it suggests the tau neutrino mass of 0.03 -- 0.1~eV, and it appears difficult to accommodate the HDM based on massive neutrinos. We point out that the hadronic axion~\cite{KSVZ} can be an alternative motivated candidate for the HDM component in the MDM model. Axion has been proposed as a solution to the strong CP problem in the QCD, and the hadronic axion (or KSVZ axion) is one version which predicts small coupling of the axion to the electron. There has been known a window of $f_{a} \sim 10^{6}$~GeV allowed by existent astrophysical and cosmological constraints if the axion coupling to photons is suppressed accidentally. This is referred to as the ``hadronic axion window.'' Our main observation is that this window gives exactly the right mass of $m_{a} \sim \mbox{a few eV}$ and the number density of the axion appropriate for the HDM component in the MDM scenario. | In this letter, we have pointed out that the hadronic axion in the hadronic axion window ($f_a\sim 10^6{\rm ~GeV}$) can automatically be a good candidate of the Hot Dark Matter component in the mixed dark matter scenario. In order to evade an astrophysical constraint from the background UV light, axion-photon-photon coupling has to be suppressed in the hadronic axion window, probably by an accidental cancellation. This scenario may be tested by detecting the axion burst from a future supernova in water \v{C}erenkov detectors, or detecting solar axions using resonant absorption. | 98 | 4 | hep-ph9804291_arXiv.txt |
9804 | astro-ph9804071_arXiv.txt | We present CCD BVI photometry of the old open cluster Berkeley 21, one of the most distant clusters in the Galactic anticentre direction, and possibly the lowest metallicity object in the open clusters sample. Its position and metal abundance make it very important for the study of the Galactic disc. Using the synthetic Colour - Magnitude Diagram method, we estimate values for distance modulus \mmm = 13.4--13.6, reddening \ebv = 0.74--0.78 (with possible differential absorption), and age = 2.2--2.5 Gyr. | Old open clusters cover a large range of distances, metallicities, and ages (Friel 1995), and that warrants their use in investigations of the chemical and dynamical evolution of our Galaxy. To study the metallicity and age distribution of open clusters with Galactocentric distance, and avoid unnecessary and dangerous biases, a key requisite is homogenous analysis of very accurate observational data, as discussed by, e.g, Janes \& Phelps (1994, JP94) Carraro \& Chiosi (1994, CC94), Friel (1995), Twarog et al. (1997,TAAT97). This is the fifth paper of a series dedicated to the examination of old open clusters of different ages and metallicities, and located at different Galactic radii: for them we measure in a homogenous way distance, age, reddening and metallicity. These quantities are derived from comparison of the observed colour-magnitude diagrams (CMDs) to synthetic ones generated by a numerical code based on stellar evolution tracks and taking into account theoretical and observational uncertainties (Tosi et al. 1991). These simulations are much more powerful than the classical isochrone fitting method to study the evolutionary status of the analysed region and have been successfully applied both to nearby irregular galaxies (Greggio et al. 1998 and references therein) and to galactic open clusters (NGC2243: Bonifazi et al. 1990; Cr261: Gozzoli et al. 1996; NGC6253: Bragaglia et al. 1997; NGC2506: Marconi et al. 1997). Berkeley 21 (Be21) is located toward the Galactic anticentre, at coordinates RA(1950) = 5:48:42, DEC(1950) = 21:46, and l$_{\rm II}$ =187$^{\circ}$, b$_{\rm II}$ = $-2.5^{\circ}$. It has already been observed by Christian \& Janes (1979, hereafter CJ), but their photographic CMD is very shallow, barely reaching the main sequence Turn-Off (TO). They deduced a substantial reddening (\ebv $\simeq$ 1.0), a large distance modulus (\mmm $\simeq$ 16), and a quite young age ($\sim 10^8$ yr). Much better data have been presented by Phelps et al. (1994, PJM94) in their compilation of old open clusters, providing \mmm = 13.9$\pm$0.2 and an age of 2.8 Gyr, derived on the basis of $\delta V$=1.6 ($\delta V$ being the magnitude difference between TO and clump stars, JP94). The metallicity has been estimated by medium-resolution spectroscopy (Friel \& Janes 1993, FJ93), but its actual value strongly depends on the adopted reddening (\ebv = 0.7$\pm$0.2, Janes 1991), with [Fe/H]= $-0.97^{+0.3}_{-0.1}$ dex. This large uncertainty, given the fact that Be21 defines the lowest metallicity limit of the open clusters sample and is one of the clusters most distant from the Galactic centre, is a further limitation for studies of the (possible) age and distance relations with chemical abundance in the Galactic disc (see also Twarog et al. 1997). In Section 2 we describe the observations and data analysis; in Section 3 we present the derived CMDs involving BVI photometry and discuss the presence of binary stars. In Section 4 we compare observed and synthetic CMDs and derive metallicity, age, distance and reddening. Finally, conclusions will be reviewed in Section 5. \begin{table} \begin{center} \caption{Log of the observations. The cluster field has its centre at RA(2000) = 5:51:46, DEC(2000) = +21:48:45. The off-cluster field has coordinates: RA(2000) = 5:52:08, DEC(2000) = +21:53:53} \begin{tabular}{lcccccccc} \hline\hline \multicolumn{1}{c}{Night} &\multicolumn{1}{c}{Field} &\multicolumn{1}{c}{Tel.} &\multicolumn{3}{c}{Exposure in seconds} &\multicolumn{1}{c}{Seeing} \\ & & & B & V & I & excursion \\ \hline Mar 4 &Centre &Danish & 120 & 120 & 120 & 0.90-1.20\\ Mar 4 &Centre &Danish & - & 120 & 900 & 0.85-1.00\\ Mar 4 &Centre &Danish & - & 600 & - & 0.95-1.10\\ Mar 5 &Centre &Danish & 120 & 120 & 120 & 1.00-1.10\\ Mar 5 &Centre &Danish &1500 & 900 & - & 0.95-1.10\\ Mar 14 &Ext. &Dutch & 120 & 60 & 60 & 1.10-1.30\\ Mar 14 &Ext. &Dutch & 1200 & 480 & 480 & 1.20-1.30\\ Mar 14 &Ext. &Dutch & 1200 & 480 & 480 & 1.15-1.30\\ \hline \end{tabular} \end{center} \label{tab-log} \end{table} \begin{figure*} \vspace{14cm} \special{psfile=figmap.ps vscale=95 hscale=95 voffset=-175 hoffset=-25} \caption{Map of the observed field, taken from our $V,B-V$ photometry. North is up and East left.} \label{fig-map} \end{figure*} | Although the fits of observed and simulated diagrams are not as satisfying for Be21 as they have been for other systems examined with the same method, we have been able to determine a fairly consistent confidence interval for its distance, age, reddening and metallicity (see Table 5) by selecting the most reliable among the models described in the previous section. They place it among the old metal poor open clusters, in a region far from the Galactic centre and of moderately high reddening. \subsection{Distance and reddening} We have derived a distance slightly smaller than previous studies, while our evaluation of the reddening is fairly consistent with past works. No previous indication of differential reddening was given in the literature, but our data definitely show it. JP94 found, for the red clumps of 23 open clusters with \dv$\ge$1.0 (i.e. older than about 1.5 Gyr) a mean absolute magnitude $M_V=0.9 \pm 0.4$, and a mean intrinsic colour $(B-V) = 0.95 \pm 0.10$. In our case, these mean values, when applied to the observed $V$=16.80 and $B-V$=1.55, would imply $(m-M)_V$=15.90, and \ebv=0.60, or $(m-M)_0$=14.04. From our best simulations, we obtain instead $(m-M)_0 \simeq $ 13.5, \ebv=0.76, corresponding to $(m-M)_V$=15.86. In other words, the clump-based distances seem to agree, but the colour of the clump stars seems to be quite different from the mean. Part of this discrepancy may be due to the high reddening affecting Be21, whose differential effect on blue and red stars leads to an apparent shrink by 0.04 of the true colour difference between TO and clump stars (Fernie 1963, Twarog 1998 private communication). On the other side, TAAT97, derive a mean $M_V=0.6 \pm 0.1$ for ten clusters not too metal-rich ranging from NGC7789 to Mel66, i.e. approximately from 1 to 5 Gyr, for which they try to measure the distance in a quite homogenous way. This translates, in our case, to $(m-M)_V$=16.2; since they do not cite a mean intrinsic $(B-V)$ no further comparison with our best choice for the reddening is possible. There are no completely reliable reddening determination for this cluster since the UBV data of CJ do not reach the MS, but our determination and that by Janes (1991) agree well. We have further compared our finding with what is expected from the spatial distribution of interstellar extinction near the Galactic plane. To this end we have considered the studies of FitzGerald (1968, fig. 3h) and Neckel \& Klare (1980, fig. 6i). In both cases, a reasonable estimate deduced from their data for low Galactic latitudes and the right longitude, is \ebv $\sim$ 0.8; FitzGerald's (1968) observations also allow for a lower value, closer to 0.5, but seem to exclude the high values, close to 1, needed by the lower metallicity tracks of any group. Janes (1991) and JPM94 give a distance modulus \mmm=13.9$\pm$0.2, somewhat larger than our results for every set of tracks. Carraro et al. (1998), working on the same data, cite a Galactocentric distance of 14.5 Kpc, also implying a distance modulus \mmm$\simeq$13.9. We have no good explanation for this difference, but we must emphasize that adopting \mmm=13.9 we would be forced to select younger ages, and this would have two major drawbacks: a worse disagreement with literature ages (see next), and a worse reproduction of the MS shape in the synthetic CMDs. \subsection{Age} Also in the case of the age, we seem to have found a value lower than given in literature. We can explain the discrepancy partly by the different techniques adopted, partly by the better quality of our data. The various parameter combinations all converge to a fairly small range of possible ages (2.1 to 2.8 Gyr, with favorite age around 2.2 Gyr). In fact, the only largely discrepant value found in our analysis is for the FRANEC98 Z=0.01 tracks, a value in strong disagreement with the spectroscopically determined metallicity. Despite the uncertainties involved in the age determination with our method, we consider it still more reliable than ages derived by other means. Nonetheless, it is not always feasible to determine the age of a cluster with the proper method of synthetic diagram fitting: to do so, high quality data, both deep and precise, are needed, and the process itself is complex. To apply this technique to all the objects of interest takes a long time, while the properties of the whole sample of open clusters are needed to study the Disc population and evolution. For this reason, several parameterizations of cluster ages, based on a handful of well studied objects, have been proposed. These methods, if uncertain in absolute value for the single cluster, yield a reasonable age ranking for the cluster system. These parameterizations are usually based on a difference, in magnitude and/or in colour, between well recognizable points of the CMD (usually the TO and the red clump), as this is much easier to measure than any absolute quantity. Note though that the precise definition of the two points, and especially of the TO, changes among authors. We will cite here the three following examples: i) Anthony-Twarog \& Twarog (1985 and later works) use the magnitude difference between the red giant clump and {\it the brightest point at the TO} ($\delta V_T$) coupled with the difference in colour between the red giant branch at the position of the clump and the bluest point of the TO ($\delta(B-V)_T$); ii) JP94 use a similar \dv, but measured between the red clump and {\it the inflection point between the MS TO and the base of the giant branch}; iii) CC94 define their $\Delta V$ as JP94, but assume that {\it the reference TO luminosity is 0.25 mag fainter than observed in the CMD}, to take into account the fact that presence of unresolved binaries tends to brighten the TO region. In the case of Be21, we have: \dv=1.8 (our measure), $\delta V_T$=1.2, \dv(JP94)=1.6, $\Delta V$=1.55. All these different definitions try to circumvent the problem that the TO point is not always easily identified in open clusters, due to field stars and binaries contamination and/or paucity of stars. The strength of our kind of analysis is that we do not judge on the basis of the observed CMD alone: we know from the tracks the exact location of the TO in each of our simulated CMDs. Given this, we have chosen to measure the magnitude difference as defined above at what we believe to be the true TO, i.e. at the point corresponding to the hottest MS point in the evolutionary tracks. JP94 correlated \dv ~with cluster ages. The calibration of their Morphological Age Index (MAI, expressed in Gyr) translates for Be21 to an age of 2.8 Gyr (based upon \dv=1.6, PJM94), marginally inconsistent with what we get from the direct comparison with evolutionary tracks (see Section 4). The age difference does not arise from any discrepancy in the two sets of data: we find \dv=1.8, quite consistent with the value given by PJM94 considering that we measure it in a slightly different way. However, we have found in the past (e.g. in the case of NGC2506, Marconi et al. 1997) that the MAI tends to overestimate ages. Anthony-Twarog \& Twarog (1985, revised by Twarog \& Anthony-Twarog 1989) proposed the so called Morphological Age Ratio (MAR), defined as MAR = $\delta V_T / \delta (B-V)_T$ (see above). This index is independent of reddening and almost independent of metallicity (Anthony-Twarog \& Twarog 1985, Buonanno et al. 1989). The calibration of the relation between MAR and ages has changed from: age = 1.4$\times$MAR Gyr (Anthony-Twarog \& Twarog 1985) to: age = 2.0$\times$MAR Gyr (Twarog \& Anthony-Twarog 1989). Applying their definition to our CMD, we find the values given in Table 6, and an age of about 2.7 to 3.9 Gyr, depending on the adopted calibration. With no attempt to give a new calibration of the MAR-age relation, simply adopting for the clusters we studied the parameters and ages in Table 6, we find: age = 2.3 $\times$ MAR -- 2.6 ($r.m.s.$ = 0.9). We did not include Be17 in this computation: it represents an extreme of the interpolation and we felt that the parameters derived from the published diagrams were too insecure. This relation gives for Be21 an age of 1.9 Gyr. Carraro et al. (1998) derive for Be21 an age of 3.1 Gyr, based on the PJM94 data and the synthetic CMD method using the Padova tracks. The difference with our results, obtained employing the same sets of tracks (although we do not interpolate in metallicity as they do), may perhaps be explained simply with the worse quality of the observational data they use. Certainly, in no case are we able to reach self-consistently such a large age. \subsection{Metallicity} We note that our method is unable to solve the problem of the cluster precise metal abundance. The comparison with evolutionary tracks can only give a coarse indication of metallicity. Too many variables are present in tracks computation to discriminate metallicity to such an extent. In fact, tracks nominally closer to the metallicity derived for Be21 from spectroscopic measurements ([Fe/H]=--0.97, or Z $\simeq$ 0.002, FJ93) appear less consistent with the observed CMD than tracks more metal rich, because of the large colour extent of the subgiant branch and, in some cases, of an excessively high reddening required to reproduce the observed colour of the MS. Anyway, what can be said is that the best fits are obtained for the slightly more metal-rich combinations, i.e. for Z=0.006 (\fe $\simeq$ -0.5) or 0.004 (\fe $\simeq$ -0.7) as compared to 0.001 (\fe $\simeq$ -1.3). This would go in favour of a metal abundance slightly higher than measured by FJ93. There is the possibility that the FJ93 scale may be underestimating cluster metallicities. TAAT97 compared it with abundances based on DDO photometry and found the FJ93 values systematically low. Another example may be the couple of clusters examined by Gratton \& Contarini (1994): they observed two giants in each cluster at high-resolution and high S/N (R=30,000, S/N $\simeq$ 100) and found for NGC2243 and Mel66 the values [Fe/H]=--0.48 and --0.38 respectively, to be compared with [Fe/H]=--0.56 and --0.51 (FJ93). FJ93 emphasized the fact that the actual value derived for the cluster metallicity from their spectra is strongly dependent on the adopted reddening: the $\pm$0.2 mag error on reddening in Janes (1991) allows for a formal uncertainty of $\pm$0.3 dex in metallicity. They also find marginal support for a \ebv ~value on the higher side, hence for a metal abundance slightly higher than the [Fe/H]=--0.97 they give. This goes in the same direction suggested by our comparisons, even if we do not find any convincing evidence for a larger reddening. Finally, we have identified the four stars studied by FJ93 (N$_{CJ}$ = 50, 406, 413, 415a, which correspond to N$_{our}$ = 50, 67, 20, 51 respectively), to check if they may be influenced by the differential reddening we found; but we consider it quite improbable, since all the four objects are within 1 arcmin from the centre of the cluster. No conclusive word can be said on Be21 metallicity, which would instead be important to know with high precision, since it could represent the lowest value for the open clusters in our Galaxy. A decisive answer would come from high resolution spectroscopy coupled with fine abundance analysis on the four stars examined by FJ93, already known to be cluster members. \bigskip\bigskip\noindent ACKNOWLEDGEMENTS \noindent We warmly thank A. Chieffi, M. Limongi and, specially, O. Straniero for having not only distributed the new FRANEC tracks in advance of publication, but even in format suitable for our purposes. We also thank J.C. Mermilliod for kindly making available his invaluable BDA open clusters database and for useful comments. We are grateful to the referee (Bruce Twarog) for his comments, extremely useful both to improve the clarity of the paper and for future applications. The bulk of the numerical code for CMD simulations has been provided by Laura Greggio. This research has made use of the Simbad database, operated at CDS, Strasbourg, France. | 98 | 4 | astro-ph9804071_arXiv.txt |
9804 | astro-ph9804137_arXiv.txt | We analyze the data of low--energy cosmic--ray $\bar p$ spectrum, recently published by the BESS Collaboration, in terms of newly calculated fluxes for secondary antiprotons and for a possible contribution of an exotic signal due to neutralino annihilation in the galactic halo. We single out the relevant supersymmetric configurations and discuss their explorability with experiments of direct search for particle dark matter and at accelerators. We discuss how future measurements with the Alpha Magnetic Spectrometer (AMS) on the Shuttle flight may disentangle the possible neutralino--induced contribution from the secondary one. | A recent analysis \cite{bess95} of the data collected by the balloon--borne BESS spectrometer on cosmic--ray antiprotons during its flight in 1995 (hereafter referred to as BESS95 data) has provided the most detailed information on the low--energy cosmic--ray $\bar{p}$'s spectrum currently available: 43 antiprotons have been detected, grouped in 5 narrow energy windows over the total kinetic--energy range $180 ~ {\rm MeV} \leq T_{\bar{p}} \leq 1.4 ~{\rm GeV}$. With this experiment the total number of measured cosmic--ray antiprotons in balloon--borne detectors over a period of more than 20 years \cite{all,hof,mitchell,moiseev,barbiellini} has more than doubled. Most remarkably, the BESS95 data provide a very useful information over the low--energy part of the ${\bar p}$ flux, where a possible distortion of the spectrum expected for secondary $\bar{p}$'s (i.e., antiprotons created by interactions of primary cosmic--ray nuclei with the interstellar medium) may reveal the existence of cosmic--ray antiprotons of exotic origin (for instance, due to pair annihilation of relic particles in the galactic halo \cite{th,noi,mitsui}, to evaporation of primordial black holes \cite{mitsui,kww} or to cosmic strings \cite{strings}). In fact, a possible discrimination between primary (exotic) and secondary $\bar p$'s is based on the different features of their low--energy spectra: in this energy regime ($ T_{\bar p} \lsim $ 1 GeV) interstellar (IS) secondary $\bar p$ spectrum is expected to drop off very markedly because of kinematical reasons \cite{gaisser}, while exotic antiprotons show a milder fall off. However, as will be discussed later on, this discrimination power is somewhat hindered by solar modulation and by some other effects affecting particle diffusion in the Galaxy. In Fig. 1 we report the cosmic--ray $\bar p$ flux at the top of the atmosphere (hereafter referred to as TOA flux) measured by BESS95 \cite{bess95}. For experimental data referring to other measurements with much less statistics see Refs.\cite{all,hof,mitchell,moiseev,barbiellini}. Also displayed in Fig. 1 are the minimal, median and maximal fluxes expected for secondary antiprotons at the time of the BESS95 data taking. These fluxes have been derived with a procedure which is described in detail in Secs. II--V. A comparison of the BESS95 data with the theoretically expected fluxes for secondary $\bar p$'s, as displayed in Fig. 1, leads to the following considerations: i) the experimental data are consistent with the theoretically expected secondary flux, within the experimental errors and the theoretical uncertainties; however, ii) the experimental flux seems to be suggestive of a flatter behaviour, as compared to the one expected for secondaries ${\bar p}$'s. Thus, natural questions arise, such as: a) how much room for exotic ${\bar p}$'s would there be in the BESS95 data, for instance in case the secondary flux is approximately given by the median estimate of Fig. 1, b) how consistent with the current theoretical models would be the interpretation of the BESS95 data in terms of a fractional presence of exotic antiprotons, and c) how this interpretation might be checked by means of independent experiments? In the present note we address these questions within an interpretation of a possible excess of $\bar p$'s at low energies in terms of primary antiprotons generated by relic neutralinos in the galactic halo \cite{note1}. The present analysis \cite{note2} is mostly meant to a clarification of many theoretical points which will be even more crucial, when a much more statistically significant experimental information on low--energy cosmic--ray antiprotons will be made available by forthcoming experiments: AMS on the precursor Shuttle flight in May 1998 and on the International Space Station Alpha (ISSA) in January 2002 \cite{ams}, the satellite--borne PAMELA experiment \cite{pamela} and balloon--borne measurements \cite{spill}. Our paper is organized as follows. In Sec. II we discuss the cosmic--ray IS proton spectrum which will be subsequently employed in deriving the secondary antiprotons. In the same section we also illustrate how we treat the solar modulation to connect the IS spectra to the corresponding TOA fluxes. In Sec. III we discuss the sources of cosmic antiprotons, both of primary and of secondary origins. Cosmic rays diffusion properties are derived in Sec. IV; the TOA $\bar p$ spectra are given in Sec. V. In Sec. VI we compare our theoretical fluxes with the BESS95 data and single out the neutralino configurations which may be relevant for the present problem. Secs. VII and VIII are devoted to an analysis on how these supersymmetric configurations can be explored by direct searches for relic neutralinos and by experimental investigation at accelerators. Conclusions and perspectives in terms of the forthcoming measurements of low--energy cosmic--ray $\bar p$'s are illustrated in Sec. IX. | We have presented a new analysis of the cosmic--ray antiprotons flux, expected on the basis of secondary $\bar p$'s, generated by interactions of cosmic--ray primaries with the interstellar medium, and of a possible exotic primary source of $\bar p$'s, originated by neutralino--neutralino annihilations in the Galactic halo. Improvements over previous calculations of secondaries depend mostly on: i) the use of a two--zone propagation model for diffusion of cosmic rays in the halo instead of the standard leaky box model; ii) the inclusion of an energy--loss effect in the propagation properties of cosmic rays (important for the antiproton low energy range considered in this paper); iii) the use of the new data on primary cosmic--ray proton spectrum, as measured by IMAX \cite{imaxp} and CAPRICE \cite{caprice}. The neutralino--induced $\bar p$ flux has been evaluated in a MSSM at the electroweak scale, which incorporates all current accelerator constraints. Use of supergravity--inspired unification conditions at large energy scale has been avoided in order not to arbitrarily constrain the neutralino phenomenology \cite{bere}. Solar modulation of the antiproton flux has been improved by analyzing the most complete set of data over the solar cycles \cite{papini} and the data on the proton spectrum of Refs. \cite{imaxp,caprice}. We have found that the most statistically relevant data on cosmic--ray antiprotons at low--energy \cite{bess95} leave some room for a possible signal from neutralino annihilation in the galactic halo. We have discussed how the relevant supersymmetric configurations may be explored with direct experiments for particle dark matter search and at accelerators. We have shown how the interplay between measurements of cosmic--ray $\bar p$'s and direct search experiments for relic particles is very intriguing and quite important in view of the significant improvements expected in these two classes of experiments in the near future. The present analysis stresses the great interest for the forthcoming AMS measurements with the Shuttle flight and on the ISSA \cite{ams}, as well as for other future measurements with balloon--borne experiments (IMAX\cite{mitchell}, BESS\cite{moiseev}) and with satellites (PAMELA) \cite{pamela}, for disentangling the secondary $\bar p$ flux from a possible primary signal of exotic nature. As an example, we give in Fig. 18 the distribution of measurements expected for AMS with the Shuttle flight according to two different hypothesis: a) dominance of the secondary contribution (lower sequence of crosses), b) significant contribution due a neutralino--induced signal (upper sequence of crosses). In our evaluation of the expected measurements we have taken into account geomagnetic cut--off effects and the expected AMS overall acceptance \cite{batt}. \begin{center} {\bf Acknowledgments} \end{center} We thank Roberto Battiston and Aldo Morselli for interesting discussions about experimental aspects related to the present paper. P.S. would like to thank the French Programme National de Cosmologie for its support. \newpage \begin{table} \caption{Values of the parameters in the expressions (\ref{eq:energy}) and (\ref{eq:rigidity}) for the IS proton flux and of the solar--modulation parameter $\Delta$. These values are obtained by best--fitting the data of Refs.[19-20] with Eqs. (\ref{eq:energy}) and (\ref{eq:rigidity}), either over the entire energy range or only over the high--energy ($T_p \geq 20$ GeV) range. First and third sets of values refer to 3--parameters fits (with Eqs.(\ref{eq:energy}) and (\ref{eq:rigidity}), respectively), second and fourth sets refer to 2--parameters fits at fixed $\Delta$, (with Eqs.(\ref{eq:energy}) and (\ref{eq:rigidity}), respectively). } \begin{center} \begin{tabular}{|c|c|c|c|} \hline & IMAX & CAPRICE & Comments \\ \hline \hline A & 12,300$\pm$1,700 & 17,600$\pm$500 & entire energy\\ $\alpha$ & 2.67$\pm$0.03 & 2.81$\pm$0.01 & range \\ $\Delta$ & 510$\pm$40 & 390$\pm$ 5 & \\ \hline A & 12,300$\pm$3,000 & 19,600$\pm$ 3,000 & \\ $\alpha$ & 2.67$\pm$0.06 & 2.85$\pm$0.04 & $T_p \geq$ 20 GeV \\ $\Delta$ & 510 (fixed) & 390 (fixed) & \\ \hline B & 16,200$\pm$2,000 & 26,000$\pm$1,200 & entire energy\\ $\gamma$ & 2.73$\pm$0.03 & 2.91$\pm$0.02 & range \\ $\Delta$ & 795$\pm$35 & 710$\pm$10 & \\ \hline B & 13,700$\pm$4,100 & 22,800$\pm$ 3,700 & \\ $\gamma$ & 2.69$\pm$0.06 & 2.88$\pm$0.04 & $T_p \geq$ 20 GeV \\ $\Delta$ & 795 (fixed) & 710 (fixed) & \\ \hline \end{tabular} \end{center} \end{table} \vfill \eject \newpage \begin{center} {\bf Figure Captions} \end{center} {\bf Fig. 1} - TOA antiproton flux as a function of the antiproton kinetic energy. The experimental points are the BESS95 data \cite{bess95}. The curves are the median (solid line), minimal (dotted line) and maximal (dashed line) secondary fluxes calculated in this paper, solar--modulated at the time of the BESS95 measurement. {\bf Fig. 2} - TOA spectra of IMAX (full circles) \cite{imaxp} and of CAPRICE (open circles) \cite{caprice} with our best--fit curves with parametrization of Eq. (\ref{eq:energy}). (The error bars are not shown when they are smaller than the dimension of the circles.) {\bf Fig. 3} - TOA spectra of IMAX (full circles) \cite{imaxp} and of CAPRICE (open circles) \cite{caprice}. The solid, dotted and dashed lines denote the median, minimal and maximal IS proton fluxes, respectively, as discussed in Sec. II. (The error bars are not shown when they are smaller than the dimension of the circles.) {\bf Fig. 4} - The grammage $\Lambda_e$ of the CNO primary elements (solid) as inferred from a two--zone diffusion model of the propagation of cosmic rays in the Galaxy. It is plotted as a function of the kinetic energy per nucleon. The dashed curve features the grammage corresponding to protons while the dotted lines delineate the interval of escape lengths inferred from the Ficenec {\it et al.} \cite{Ficenec} observations on $^{3}$He at TOA energies comprised between 100 MeV/n and 1.6 GeV/n. \label{fig:grammage} {\bf Fig. 5} - IS secondary antiproton spectra as functions of the ${\bar p}$ kinetic energy. Solid, dotted and dashed lines denote the fluxes obtained from the median, minimal and maximal IS primary proton fluxes. The dot--dashed line denotes the median ${\bar p}$ spectrum, when the ${\bar p}$ energy losses are neglected. {\bf Fig. 6} - Coefficient $C_{\rm susy}(T_{\bar p},f)$ as a function of the ${\bar p}$ kinetic energy for different values of the flattening parameter $f$. {\bf Fig. 7} - Time variation of the solar--modulation parameter $\Delta$. Full circles represent the best--fit values to the PGS average fluxes at minima (MIN) and at maxima (MAX) and to the fluxes of IMAX \cite{imaxp} and of CAPRICE \cite{caprice}; the open circle refers to the BESS95 data taking period and the cross denotes the extrapolated value of $\Delta$ at the time relevant for the future AMS Shuttle flight (May 1998). {\bf Fig. 8} - Solar modulation of the IS median secondary antiproton flux calculated in this paper. Solid line is the IS spectrum; dashed (dotted) line is the solar--modulated spectrum at minima (maxima). {\bf Fig. 9} - Solar modulation of the IS antiproton flux, due to neutralino annihilation for the representative neutralino configuration with $m_{\chi} = 62$ GeV, $P = 0.98$ and $\Omega_{\chi} h^2 = 0.11$. Solid line is the IS spectrum; dashed (dotted) line is the solar--modulated spectrum at minima (maxima). {\bf Fig. 10} - TOA antiproton fluxes versus the antiproton kinetic energy. The BESS95 data \cite{bess95} are shown by crosses. The dashed line denotes the median secondary flux, the dotted one denotes the primary flux due to neutralino annihilation in the halo for a neutralino configuration with $m_{\chi} = 62$ GeV, $P = 0.98$ and $\Omega_{\chi} h^2 = 0.11$. Solid line denotes the calculated total flux. {\bf Fig. 11} - Scatter plots for configurations of set $M$ (a) and set $N$ (b) in the P--$\tan \beta$ plane. {\bf Fig. 12} - Scatter plots for configurations of set $M$ (a) and set $N$ (b) in the P--$m_{\chi}$ plane. {\bf Fig. 13} - Scatter plots for configurations of set $R$ in the P--$\tan \beta$ plane (a) and in the P--$m_{\chi}$ plane (b). {\bf Fig. 14} - Scatter plots for configurations of set $R$ in the P--$\tan \beta$ plane (a) and in the P--$m_{\chi}$ plane (b) for a flattening of $f = 0.5$. {\bf Fig. 15} - Scatter plot of the values of $\xi \sigma_{\rm scalar}^{\rm nucleon}$ versus the neutralino mass for the configurations of set $M$ (a) and of set $N$ (b). The open curve denotes the (90 \% C.L.) upper bound obtained from experimental data of Ref. \cite{few}. The region delimited by a closed contour is the one singled out by the experiment of Ref. \cite{dama} as possibly indicative of an annual modulation effect. The total local dark matter density is normalized here to the value $\rho_l = 0.4$ GeV cm$^{-3}$. The dashed line shows the discovery potential in case of an improvement by a factor of 10 in current sensitivities for experiments of direct search for particle dark matter. {\bf Fig. 16} - Correlation between $\xi \sigma_{\rm scalar}^{\rm nucleon}$ and the neutralino relic density $\Omega_{\chi} h^2$ for configurations of set $M$. {\bf Fig. 17} - Scatter plot for configurations of set $M$ in the $m_h$--$\tan \beta$ plane. The region on the left of the dashed line denoted by (a) is excluded by current LEP experimental data \cite{lep182}, the one on the right of the dashed line (b) is theoretically disallowed. The other lines display the LEP reach at luminosity $L = 200$ pb$^{-1}$ and various energies \cite{alt}: (A) discovery potential at ${\sqrt s} = 192$ GeV; (B) discovery potential at ${\sqrt s} = 200$ GeV; (C) exclusion at ${\sqrt s} = 200$ GeV. {\bf Fig. 18} - Expected distribution of measurements with the AMS Shuttle flight according to two different hypothesis: a) dominance of the secondary contribution (lower sequence of crosses), b) significant contribution due a neutralino--induced signal (upper sequence of crosses). The dashed line denotes the secondary flux, the dotted one denotes the primary flux due to neutralino annihilation in the halo for a neutralino configuration with the representative values: $m_{\chi} = 62$ GeV, $P = 0.98$ and $\Omega_{\chi} h^2 = 0.11$. Solid line denotes the calculated total flux. | 98 | 4 | astro-ph9804137_arXiv.txt |
9804 | astro-ph9804301_arXiv.txt | We present a fully covariant and gauge-invariant calculation of the evolution of anisotropies in the Cosmic Microwave Background (CMB) radiation. We use the physically appealing covariant approach to cosmological perturbations, which ensures that all variables are gauge-invariant and have a clear physical interpretation. We derive the complete set of frame-independent linearised equations describing the (Boltzmann) evolution of anisotropy and inhomogeneity in an almost Friedmann-Robertson-Walker (FRW) Cold Dark Matter (CDM) universe. These equations include the contributions of scalar, vector and tensor modes in a unified manner. Frame-independent equations for scalar and tensor perturbations, which are valid for any value of the background curvature, are obtained straightforwardly from the complete set of equations. We discuss the scalar equations in detail, including the integral solution and relation with the line of sight approach, analytic solutions in the early radiation dominated era, and the numerical solution in the standard CDM model. Our results confirm those obtained by other groups, who have worked carefully with non-covariant methods in specific gauges, but are derived here in a completely transparent fashion. | The cosmic microwave background radiation (CMB) occupies a central role in modern cosmology. It provides us with a unique record of conditions along our past lightcone back to the epoch of decoupling (last scattering), when the optical depth to Thomson scattering rises suddenly due to Hydrogen recombination. Accurate observations of the CMB anisotropy should allow us to distinguish between models of structure formation and, in the case of non-seeded models, to infer the spectrum of initial perturbations in the early universe. Essential to this programme is the accurate and reliable calculation of the anisotropy predicted in viable cosmological models. Such calculations have a long history, beginning with Sachs \&\ Wolfe (1967) who investigated the anisotropy on large scales $(\gtrsim 1^{\circ})$ by calculating the redshift back to last scattering along null geodesics in a perturbed universe. On smaller angular scales one must address the detailed local processes occurring in the electron/baryon plasma prior to recombination, and the effects of non-instantaneous last scattering. These processes, which give rise to a wealth of structure in the CMB power spectrum on intermediate scales and damping on small scales (see, for example, Silk (1967, 1968)), are best addressed by following the photon distribution function directly from an early epoch in the history of the universe to the current point of observation. This requires a numerical integration of the Boltzmann equation, and has been carried out by many groups, of which Peebles \&\ Yu (1970), Bond \&\ Efstathiou (1984, 1987), Hu \&\ Sugiyama (1995), Ma \&\ Bertshcinger (1995), Seljak \&\ Zaldarriaga (1996) is a representative sample. The calculation of CMB anisotropies is simple in principle, but in reality is plagued with subtle gauge issues (Stoeger, Ellis, \&\ Schmidt 1991; Stoeger et al.\ 1995; Challinor \&\ Lasenby 1998). These problems arise because of the gauge-freedom in specifying a map $\Phi$ between the real universe (denoted by $S$) and the unperturbed background model(denoted by $\bar{S}$)~\cite{ellis89a}, which is usually taken to be a Friedmann-Robertson-Walker (FRW) universe. The map $\Phi$ identifies points in the real universe with points in the background model, thus defining the perturbation in any quantity of interest. For example, for the density $\rho$ as measured by some physically defined observer, the perturbation at $x\in S$ is defined to be $\delta \rho(x)\equiv \rho(x) - \bar{\rho}(\bar{x})$, where $\bar{\rho}$ is the equivalent density in the background model, and $x$ maps to $\bar{x}$ under $\Phi$. The map $\Phi$ is usually (partially) specified by imposing coordinate conditions in $S$ and $\bar{S}$. Any residual freedom in the map $\Phi$ after the imposition of the coordinate conditions (gauge-fixing) gives rise to the following gauge problems: (i) the map cannot be reconstructed from observations in $S$ alone, so that quantities such as the density perturbation, which depend on the specific map $\Phi$, are necessarily not observable; (ii) if the residual gauge freedom allows points in $\bar{S}$ to be mapped to physically inequivalent points in $\bar{S}$ in the limit that $S=\bar{S}$, then unphysical gauge mode solutions to the linearised perturbation equations will exist. There are several ways to deal with the gauge problems described above. In the earliest approach~\cite{lifshitz46}, one retains the residual gauge freedom (in the synchronous gauge) but keeps track of it so that gauge mode solutions can be eliminated. Furthermore, the final results of such a calculation must be expressed in terms of the physically relevant, gauge-invariant quantities. Although there is nothing fundamentally wrong with this approach if carried out correctly, it suffers from a long history littered with confusion and errors. The need to express results in terms of gauge-invariant variables suggests that it might be beneficial to employ such variables all along as the dynamical degrees of freedom in the calculation. A further advantage of such an approach is that gauge modes are automatically eliminated from the perturbation equations when expressed in terms of gauge-invariant variables. This is the approach adopted by Bardeen (1980), who showed how to construct gauge-invariant variables for scalar, vector and tensor modes in linearised perturbation theory, by taking suitable linear combinations of the gauge-dependent perturbations in the metric and matter variables. This approach has been used in several calculations of the CMB anisotropy (see, for example, Abbott \&\ Schaefer (1986) and Panek (1986)). However, the Bardeen variables are not entirely satisfactory. The approach is inherently linear, so that the variables are only defined for small departures from FRW symmetry. Furthermore, the approach assumes a non-local decomposition of the perturbations into scalar, vector and tensor modes at the outset, each of which is then treated independently. As a result, the Bardeen variables are only gauge-invariant for the restricted class of gauge-transformations that respect the scalar, vector and tensor splitting. Finally, although the Bardeen variables are gauge-invariant, they are not physically transparent, in that, in a general gauge, they do not characterise the perturbations in a manner that is amenable to simple physical interpretation. An alternative scheme for the gauge-invariant treatment of cosmological perturbations was given by Ellis \&\ Bruni (1989) (see also, Ellis, Hwang, \&\ Bruni (1989)) who built upon earlier work by Hawking (1966). In this approach, which is derived from the covariant approach to cosmology/hydrodynamics of Ehlers and Ellis (Ehlers 1993; Ellis 1971), the perturbations are described by gauge-invariant variables that are covariantly defined in the real universe. This ensures that the variables have simple physical interpretations in terms of the inhomogeneity and anisotropy of the universe. Since the definition of the covariant variables does not assume any linearisation, exact equations can be found for their evolution, which can then be linearised around the chosen background model. Furthermore, the covariant approach does not employ the non-local decomposition into scalar, vector or tensor modes, at a fundamental level. If required, the decomposition can be performed at a late stage in the calculation to aid solving the equations. Even if one denies that working with gauge-invariant variables is a significant advantage, the key advantage of the covariant approach, however, is that one is able to work exclusively with physically relevant quantities, satisfying equations that make manifest their physical consequences. The covariant and gauge-invariant approach has already been applied to the line of sight calculation of CMB anisotropies under the instantaneous recombination approximation (Dunsby 1997; Challinor \&\ Lasenby 1998), and has been used to obtain model-independent limits on the inhomogeneity and anisotropy from measurements of the CMB anisotropy on large scales~\cite{maartens95}. In this paper, we extend the methodology developed in these earlier papers, to give a full kinetic theory calculation of CMB anisotropies valid on all angular scales. Our motivation for reconsidering this problem is two-fold. Firstly, it is our belief that the covariant and gauge-invariant description of cosmological perturbations provides a powerful set of tools for the formulation of the basic perturbation equations, and their subsequent interpretation, which are superior to the techniques usually employed in such calculations for the reasons discussed above. In particular, by applying covariant methods for the problem of the generation of CMB anisotropies, we can expect the same advantages of physical clarity and unification that have already been demonstrated in other areas, (Ellis et al.\ 1989; Bruni, Ellis, \&\ Dunsby 1992; Dunsby, Bruni, \&\ Ellis 1992; Dunsby, Bassett, \&\ Ellis 1996; Tsagas \&\ Barrow 1997). The approach described here brings the underlying physics to the fore, and can only help to consolidate our rapidly growing understanding of the physics of CMB anisotropies. Furthermore, although we only consider the linearised calculation here, the extension of these methods to the full non-linear case is quite straightforward (Maartens, Gebbie, \&\ Ellis 1998). Our second motivation is to perform an independent verification of the results of other groups (for example, Ma \&\ Bertschinger (1995)), with a methodology that is free from any of the gauge ambiguities that have caused problems and confusion in the past. Given the potential impact on cosmology of the next generation of CMB data, we believe that the above comments provide ample justification for reconsidering this problem. For definiteness we consider the cold dark matter (CDM) model, although the methods we describe are straightforward to extend to other models. We have endeavoured to make this paper reasonably self-contained, so we begin with a brief overview of the covariant approach to cosmology and define the key variables we use to characterise the perturbations in Section~\ref{sec_cov}. We then go on to present a complete set of frame-independent equations describing the evolution of the matter components and radiation in Section~\ref{sec_eqs} in an almost-FRW universe (with arbitrary spatial curvature). These equations, which employ only covariantly defined, gauge-invariant variables, are independent of any harmonic analysis; they describe scalar, vector and tensor perturbations in a unified manner. Many of the equations have simple Newtonian analogues, and their physical consequences are far more transparent than the equations that underlie the metric-based approaches. Equations pertinent to scalar modes, see Section~\ref{sec_scal}, and tensor modes, see Section~\ref{sec_tens}, can be obtained from the full set of equations with very little effort, and are useful at this late stage in the calculation as an aid to solving the linearised equations. A significant feature of this approach is that a covariant angular decomposition of the distribution functions is made early on in the calculation, before any splitting into scalar, vector and tensor modes. This allows scalar, vector and tensor modes to be treated in a more unified manner. In particular, the azimuthal dependence of the moments of the distribution functions does not have to be put in by hand (after inspection of the azimuthal dependence of the other terms in the Boltzmann equation), as happens in most metric-based calculations. This is particularly significant for tensor modes where the required azimuthal dependence is non-trivial and is different for the two polarisations of gravitational waves. We consider the equations for scalar modes in considerable detail. We present the integral solution of the Boltzmann multipole equations in a $K=0$ almost-FRW universe, and discuss the relation between line of sight methods (which employ lightlike integrations along the lightcone) and the Boltzmann multipole approach (where a timelike integration is performed). We derive analytic solutions for scalar modes in the early radiation dominated universe, that are used as initial conditions for the numerical solution of the scalar equations, the results of which we describe in Section~\ref{sec_num}. In Section~\ref{sec_tens} we give a brief discussion of the tensor equations in the covariant approach. The covariant angular decomposition naturally gives rise to a set of variables that describe the temperature anisotropy in a more direct manner than in the conventional metric-based approaches. This is particularly apparent for tensor perturbations, where the CMB power spectrum at a given multipole $l$ is determined by the $l-2$, $l$ and $(l+2)$-th moments of the conventional decomposition of the photon distribution function, which obscures the physical interpretation of these moments. Finally, we end with our conclusions in Section~\ref{sec_conc}. Ultimately, our results confirm those of other groups (for example, Ma \&\ Bertschinger (1995)) who have performed similar calculations by working carefully in specific gauges, but are obtained here with a unified methodology that is more physically transparent and less prone to lead to confusion over subtle gauge effects. We employ standard general relativity and use a $(+---)$ metric signature. Our conventions for the Riemann and Ricci tensors are fixed by $[\nabla_{a},\nabla_{b}] u^{c} = -{\clr_{abd}}^{c} u^{d}$, and $\clr_{ab} \equiv {\clr_{acb}}^{c}$. Round brackets around indices denote symmetrisation on the indices enclosed, and square brackets denote antisymmetrisation. We use units with $c=G=1$ throughout, and a unit of distance of $\mpc$ for numerical work. | \label{sec_conc} We have shown how the full kinetic-theory calculation of the evolution of CMB anisotropies and density inhomogeneities can be performed in the covariant approach to cosmology (Ehlers 1993; Ellis 1971), using the gauge-invariant perturbation theory of Ellis \&\ Bruni (1989). Adopting covariantly-defined, gauge-invariant variables throughout ensured that our discussion avoided the gauge ambiguities that appear in certain gauges, and that all variables had a clear, physical interpretation. We presented a unified set of equations describing the evolution of photon and neutrino anisotropies and cosmological perturbations in the CDM model, which were independent of a decomposition into scalar, vector or tensor modes and the associated harmonic analysis. Although we only considered the case of linear perturbations around an FRW universe here, it is straightforward to extend the approach to include non-linear effects (Maartens et al. 1998), which should allow a physically transparent discussion of second-order effects on the CMB. Indeed, the ease with which one can write down the exact equations for the physically relevant variables is one of the major strengths of the covariant approach. The linear equations describing scalar modes and tensor modes were obtained from the full set of equations in a straightforward and unified manner, highlighting the advantage of having the full equations (independent of the decomposition into scalar, vector and tensor modes) available. For the scalar case, the Boltzmann multipole equations for the moments of the distribution functions obtained here were equivalent to those usually seen in the literature. However, for tensor modes, the covariant approach led naturally to a set of moment variables that more conveniently describe the temperature anisotropy than those usually employed. For scalar modes, we discussed the solution of the perturbation equations in detail, including the integral solution of the Boltzmann multipole equations and the relation between the timelike integrations performed in the multipole approach to calculating CMB anisotropies, and the lightlike integrations of the line of sight approach. The numerical solution of the scalar equations in a $K=0$, almost-FRW, CDM universe were also discussed. Our numerical results provide independent confirmation of those of other groups, (see, for example, Ma \&\ Bertschinger (1995) and Seljak \&\ Zaldarriaga (1996)), who have obtained their results by employing non-covariant methods in specific gauges. Typically, these methods require one to keep careful track of all residual gauge-freedom, both to enable identification of any gauge-mode solutions, and to ensure that the final results quoted are gauge-invariant (and hence observable). Fortunately, the isotropy of the photon distribution function in an exact FRW universe ensures that the CMB power spectrum, as calculated from the gauge-dependent perturbation to the distribution function, is gauge-invariant for $l\geq 1$. We hope to have shown the ease with which the covariant approach to cosmology can be applied to the problem of calculating CMB anisotropies. The covariant and gauge-invariant method discussed here frees one from the gauge problems that have caused confusion in the past, and focuses attention on the physically relevant variables in the problem and the underlying physics. Future work in this area will include the discussion of non-linear effects (Maartens at al. 1998), the inclusion of polarisation, and the effects of hot dark matter, all of which can be expected to bring the same advantages of physical clarity and transparency that we hope to have demonstrated here. | 98 | 4 | astro-ph9804301_arXiv.txt |
9804 | astro-ph9804184_arXiv.txt | This paper describes the first optical spectroscopic survey of class I sources (also known as embedded sources and protostars) in the Taurus-Auriga dark cloud. We detect 10 of the 24 known class I sources in the cloud at 5500--9000 \AA. All detected class I sources have strong H$\alpha$ emission; most also have strong [O~I] and [S~II] emission. These data -- together with high quality optical spectra of T Tauri stars in the Taurus-Auriga cloud -- demonstrate that forbidden emission lines are stronger and more common in class I sources than in T Tauri stars. Our results also provide a clear discriminant in the frequency of forbidden line emission between weak-emission and classical T Tauri stars. In addition to strong emission lines, three class I sources have prominent TiO absorption bands. The M-type central stars of these sources mingle with optically visible T Tauri stars in the HR diagram and lie somewhat below both the birthline for spherical accretion and the deuterium burning sequence for disc accretion. | Examining the earliest phases of low mass stellar evolution requires observations of protostars deeply embedded in the dense cores of nearby molecular clouds. These ``class I'' sources (\cite{lad87}) have blackbody-like spectral energy distributions that peak at wavelengths of 30--100 \mum~and bolometric luminosities, $L_{\rm b} \approx$ 0.1--100 \lsun~(\cite{ada87}; \cite{mye87}; \cite{wil89}; \cite{ken90}; \cite{gre94}). Despite many detailed studies of their circumstellar environments (see, for example, \cite{ta91a},b; \cite{and93}; \cite{ter93}; \cite{mor92}, 1995; \cite{bon96}; \cite{gom97}; \cite{hog97}; \cite{whi97}), understanding the stellar physics of these objects has proved elusive. Comparisons of observed bolometric luminosity functions with models is straightforward but controversial (\cite{wil89}; \cite{ken90}; \cite{fl94a}, 1994b). The apparent lack of photospheric features in many objects has led several groups to abandon the HR diagram as a means for testing stellar evolutionary tracks of the youngest stars. These groups have proposed the bolometric temperature (\cite{mye98} and references therein), the submillimeter flux (\cite{sar96}), and the visual extinction (\cite{ada90}) to replace effective temperature and have developed models to place evolving pre--main-sequence stars in their modified evolutionary diagrams. The accuracy of these techniques remains uncertain, because the methods are new and incompletely tested. In this paper, we report an optical spectroscopic survey designed to detect photospheric absorption features from the central stars of class I sources in the Taurus-Auriga cloud. Although the line-of-sight extinction to the central star is large, $A_V \approx$ 30--60 mag (\cite{whi97}), large ground-based telescopes can detect optical light scattered off cavities in the infalling envelopes of many objects. Optical data also provide the best measure of spectral types for pre--main-sequence stars. In general, I-band and J-band data are least contaminated by emission from an accretion disc and its associated boundary layer or magnetic accretion column (\cite{kh90}). However, the very large continuum veiling detected on near-IR spectra of some class I sources (\cite{cas92}; \cite{gr96a}, 1996b) favors I-band spectra, because the J-band veiling can be large if the disc extends to the stellar photosphere (\cite{kh90}; \cite{ken96}). Finally, optical spectra of class I sources allow an unambiguous comparison with optically brighter T Tauri stars, which have known spectral types in a well-calibrated system (see, for example, Kenyon \& Hartmann 1995; KH95 hereafter). Our results provide the first optical detection of M-type absorption features in an embedded protostar. We identify TiO absorption bands in three Taurus-Auriga class I sources; one other star may have TiO features and a fifth star may have K-type absorption features. We use optical spectra of T Tauri stars to calibrate the spectral types of class I sources and then construct a complete HR diagram for the Taurus-Auriga cloud. These data, coupled with new evolutionary tracks for protostars accreting from discs and two spectral types derived from near-IR spectra (\cite{gr96b}), show that class I sources in Taurus-Auriga mingle with T Tauri stars and lie below the birthline in the HR diagram. We also detect strong emission lines on the spectra of all protostars. Forbidden emission from [N~II] and [S~II] is much more common among class I sources than older, optically brighter stars having the same bolometric luminosity. The fluxes of forbidden emission lines also seem stronger among class I sources than other pre--main-sequence stars in the cloud. We find no evidence that the permitted emission lines, such as H$\alpha$ and He~I, are more common or stronger than in T Tauri stars. These results extend and confirm previous conclusions that jet activity declines as a pre--main-sequence star contracts to the main-sequence. We describe our observations in Sec. 2, explain our results in Sec. 3, and conclude with a brief discussion in Sec. 4. | In the previous sections, we have described the first optical spectroscopic survey of class I, embedded sources in a single molecular cloud. We supplemented these data with high quality optical spectra of a representative sample of older and optically brighter T Tauri stars. The combined set of spectra show that the optical spectra of class I sources qualitatively resemble the optical spectra of T Tauri stars. Our analysis further reveals common physical properties and substantial differences between class I sources and T Tauri stars, as summarized below. Our data provide the first indication that the distribution of stellar spectral types among class I sources may not be very different from that of WTTS and CTTS. Of the five class I sources with strong optical continua, one (L1489 IRS) is a continuum + emission source, three are M-type stars, and another (04264+2433) may have an M-type central star. To the best of our knowledge, {\it these are the first low mass protostars with measured optical spectral types.} This sample is too small for a meaningful comparison with the distribution of spectral types among more evolved pre--main-sequence stars in the cloud. We note, however, that the median spectral type for WTTS and CTTS is K7-M0 and that the frequency of continuum + emission sources is $\sim$ 5\%--10\% (KH95). Published observations indicate other similarities between class I sources and older pre--main-sequence stars in Taurus-Auriga. First, class I sources have the same intrinsic near-IR colors as do CTTS. Whitney \etal (1997) show that the observed near-IR colors of class I sources can be modeled as a CTTS surrounded by an infalling envelope with an optical extinction, $A_V \approx$ 30--60 mag. This analysis leads to the conclusion that the radiation from the star and inner disc of a class I source is similar to that of a T Tauri star (see also \cite{gr96b}; Calvet \etal 1997 reach a different conclusion). Second, the bolometric luminosity distributions of class I sources, CTTS, and WTTS are indistinguishable (KH95). All three groups of pre--main-sequence stars have median luminosities of $L_{\rm b} \approx$ 0.5--0.8 \lsun. This unexpected result is supported by the positions of class I sources in the HR diagram. Our data show that class I sources have luminosities and effective temperatures very similar to those of CTTS and WTTS in the cloud. These conclusions are surprising, because a class I source should have a larger luminosity once it has accreted nearly all of its final mass, and this luminosity should decline with time as the star approaches the main sequence (see, for example, \cite{sta83}, 1988; \cite{pal93}; \cite{har97}; Fig. 8). The current sample, however, is too small to test stellar models in detail. The errors in luminosity and effective temperature are also too large. Observations with the next generation of ground-based telescopes will undoubtedly expand the sample, reduce the errors, and provide better tests of protostellar accretion theory. One feature that distinguishes class I sources is their strong forbidden-line emission. As a group, class I sources are much more likely to have forbidden-line emission than CTTS or WTTS (Fig. 7). This result confirms previous conclusions from imaging data (e.g., \cite{gom97}) and indicates that class I sources are more likely to drive outflows than CTTS or WTTS (see also \cite{bon96}; \cite{mor92}, 1994). The equivalent widths of the forbidden lines are also larger in class I sources than in CTTS or WTTS. Although some large equivalent widths may be due to very weak optical continua, the [S~II] equivalent widths in HH31 IRS2 -- a class I source with a prominent TiO absorption band -- are larger than observed in {\it any} CTTS in our sample (see Tables 1--2). Deeper optical spectra of our sample and other class I sources would clarify this point. Our sample is not large enough to test whether class I sources also have more prominent {\it permitted} emission lines than CTTS. The median H$\alpha$ equivalent width for class I sources, $\sim$ 90 \AA, is much larger than the median equivalent width for CTTS, $\sim$ 30--40 \AA. This difference is roughly what we expect if class I sources have larger continuum veiling than CTTS (\cite{cas96}; \cite{gre97}) and if the H$\alpha$ equivalent width correlates with veiling (\cite{har95} and references therein). However, the frequency of He~I $\lambda\lambda$5876, 6678 emission among class I sources is roughly comparable to that among CTTS. We measure a He~I emission frequency of 50\% among 6 class I sources with reasonable signal-to-noise at 6000 \AA, 57\% among 14 flat-spectrum sources, and 65\% among 46 class II sources. For both emission lines, the class I sample is probably biased against small equivalent widths, because class I sources without emission lines are probably fainter than sources with emission lines. A deeper survey with a larger telescope could enlarge the sample of class I sources with high quality optical spectra. These data would provide a good test for differences in the distribution of H$\alpha$ equivalent widths between class I sources and CTTS. These results fit into the general picture of Taurus-Auriga class I sources developed in KH95 and in Kenyon \etal (1990). In this picture, class I sources are envelopes of gas and dust falling into the central star-disc system at rates of a few $\times~10^{-6}~\msunyr$ (see also \cite{ada87}; \cite{ke93a}, 1993b; \cite{whi97}). Bell \& Lin (1994) show that the stable accretion rate through the disc onto the central star is either very low -- $\lesssim$ a few $\times~10^{-7}~\msunyr$ -- or very high -- $\gtrsim$ a few $\times~10^{-5}~\msunyr$ -- compared to the infall rate. The disc spends most of its time in the low accretion rate state; the disc mass then slowly increases with time until it reaches a critical level and evolves to the high accretion rate state. This model explains the low observed luminosities of nearly all class I sources as well as the occasional high luminosity of a source such as L1551 IRS5. Models with time-dependent disc accretion also qualitatively account for the evolution of forbidden and permitted emission lines in pre--main-sequence stars. We expect the time-averaged accretion rate through the disc to decline as the envelope disperses. If the H$\alpha$ and other permitted emission lines of class I sources form in the accretion region of the inner disc as in CTTS, then the median H$\alpha$ equivalent width should decline as a pre--main sequence star evolves from a class I source into a CTTS and then into a WTTS. Most models for jet formation link the mass loss rate in the jet to the mass accretion rate through the disc (see, for example, \cite{cab90}; \cite{naj94}; \cite{sh94a}, 1994b), so we expect forbidden emission to decline as well. Explaining the observations of emission line equivalent widths with a quantitative model of a dispersing envelope and evolving disc, however, is currently beyond our reach. Finally, our results further demonstrate the advantages of optical spectra. Recent surveys of larger samples of class I sources using near-IR spectroscopy have yielded only two spectral types each in Taurus-Auriga (\cite{cas96}; \cite{gr96b}) and $\rho$ Oph (\cite{gr96b}, 1997). Casali \& Eiroa (1996; see also \cite{cas92}; \cite{gr96b}, 1997) conclude that continuum emission from dust in a circumstellar disc or envelope veils photospheric absorption features on near-IR spectra of class I sources. Preliminary results further suggest that this veiling is larger in class I sources than in CTTS or WTTS (\cite{cas96}; \cite{gr96b}, 1997). Dust emission is much weaker relative to a normal stellar photosphere at shorter wavelengths, $\lesssim 1~\mu$m, so optical spectra may yet provide the best measure of spectral types in class I sources. \vskip 6ex We thank the staffs of the MMT, Palomar, and Whipple Observatories for assistance with our observations. Fred Chaffee kindly acquired several spectra of the class II sources listed in Table 1. Susan Tokarz reduced the FAST spectra and graciously assisted with the reduction of the MMT and Palomar spectra. We also thank M. Geller, M. G\'omez, C. Lada, A. Mahdavi, and B. Whitney for advice and comments. The suggestions of an anonymous referee improved our presentation. Observations at the Palomar Observatory were made as part of a continuing collaborative agreement between Palomar Observatory and the Jet Propulsion Laboratory. Portions of this research were supported by the National Aeronautics and Space Administration through grant NAGW-2919 and by the Space Telescope Science Institute through grant GO-06132.01-94A. C.A.T. thanks the Royal Society and the Hungarian Academy of Sciences for an exchange fellowship during the majority of his contribution to this project. \vfill \eject | 98 | 4 | astro-ph9804184_arXiv.txt |
9804 | astro-ph9804229_arXiv.txt | We present and discuss 25 spectra obtained in November 1996, covering all phases of the CAL 87 binary system. These spectra are superior both in signal-to-noise and wavelength coverage to previously published data so that additional spectral features can be measured. Photometry obtained on the same nights is used to confirm the ephemeris and to compare with light curves from previous years. Analysis of the color variation through the orbital cycle has been carried out using archival MACHO data.\footnote{This paper utilizes public domain data obtained by the MACHO Project, jointly funded by the US Department of Energy through Lawrence Livermore National Laboratory under contract W7405-Eng-48, the National Science Foundation through the Center for Particle Astrophysics of the University of California undercooperative agreement AST-8809616, and the Mount Stromlo and Sidings Springs Observatory by the Bilateral Science Technology and Regional Development.} When a barely resolved red field star is accounted for, there is no ($V-R$)-color variation, even through eclipse. There have been substantial changes in the depth of minimum light since 1988; it has decreased more than 0.5 mag in the last several years. The spectral features and radial velocities are also found to vary not only through the 0.44-day orbit but also over timescales of a year or more. Possible interpretations of these long-term changes are discussed. The 1996 spectra contain phase-modulated Balmer absorption lines not previously seen, apparently arising in gas flowing from the region of the compact star. The changes in emission-line strengths with orbital phase indicate there are azimuthal variations in the accretion disk structures. Radial velocities of several lines give different amplitudes and phasing, making determination of the stellar masses difficult. All solutions for the stellar masses indicate that the companion star is considerably less massive than the degenerate star. The Balmer absorption-line velocities correspond to masses of $\sim$1.4M$_{\odot}$ for the degenerate star and $\sim$0.4M$_{\odot}$ for the mass donor. However, the strong He II emission lines indicate a much more massive accreting star, with M$_X>$4M$_{\odot}$. | CAL 87 has long been known as one of only a small number of luminous X-ray binaries in the Large Magellanic Cloud (Long, Helfand, \& Grabelsky 1981, Pakull et al.\ 1988). Its optical spectrum, with He II and H emission lines on a very blue continuum, shows it to be similar to galactic low-mass X-ray binaries. However, its X-ray spectrum reveals CAL 87 to be one of the rare, very luminous (L$_{bol}\geq10^{38}$ erg s$^{-1}$) supersoft sources (SSS) which have little or no radiation above $\sim$0.5 keV (e.g. Tr\"umper et al.\ 1991, Greiner 1996). The SSS are widely thought to be binaries in which a white dwarf is undergoing rapid accretion from a more massive companion, leading to steady nuclear burning on the surface of the white dwarf (van den Heuvel et al.\ 1992). CAL 87 is unique among supersoft sources in having both optical and X-ray eclipses (Callanan et al.\ 1989, Cowley et al.\ 1990: CSCH) which provide extra information about the disk structure and in principle help to constrain the stellar masses. In the original spectroscopic data of CSCH, He II 4686\AA\ emission was shown to move with K$=$40 km s$^{-1}$ and proper phasing with respect to the eclipse so that it was interpreted as due to orbital motion of the compact star. CSCH concluded that the compact star had a mass $\ge$6M$_{\odot}$, and hence these data implied the presence of a black hole. However, optical spectra taken a few years later, combined with velocities from lines in the far UV (Hutchings et al.\ 1995), showed a quite different behavior. The velocity amplitude was larger and the phasing was very different, indicating that at times the velocities are not entirely due to orbital motion and the line-formation regions change. One problem is that spectroscopic determination of the masses is not entirely straightforward since the nearly edge-on view of the system ($i\sim78^{\circ}$) causes complications in the line profiles due to motions in the accretion disk. In this paper we report on a series of spectroscopic data with improved signal-to-noise (S/N) and new, concurrent photometry which we obtained at CTIO in 1996 in order to conduct a more thorough investigation of CAL 87. | We have discovered that the photometric eclipse depth has changed over recent years, and this variation may even be cyclic. The increase of light observed at mid-eclipse may be due to a geometrical change in the disk structure, without a significant change in total luminosity -- such as a bright structure above or below the disk plane. Alternatively, the disk may grow in size in its plane, accompanied by extra absorption due to the outflowing material, along lines-of-sight near the center. The emergence of the Balmer absorption in the 1996 spectra suggests the latter scenario. In 1994 November we obtained three optical spectra of CAL 87 which were reported by Hutchings et al.\ (1995) in their discussion of eight ultraviolet spectra observed by HST in 1995 January. We have remeasured the optical spectra, although they are of considerably lower quality than the new data and only show the strongest lines. In 1994--5 He II 4686\AA\ and 1640\AA\ emission-line velocities show a very different phasing (maximum velocity near $\phi\sim0.9$) and higher amplitude (K$\sim$70 km s$^{-1}$) compared to both the CSCH data (K=40 km s$^{-1}$) and present data (K$\sim$30 km s$^{-1}$). The equivalent width of 4686\AA\ was lower in the 1994 spectra than in 1996 by a factor of nearly two. In two of the three 1994 optical spectra weak absorption features are visible at H$\beta$ and H$\gamma$, similar to those seen in 1996, but the low S/N make the velocities unreliable. Thus, it appears that there are long-term spectral changes that may include significant non-orbital motions. The O VI lines might give a clean measure of orbital motion, as they arise in the inner disk, but they are very weak, resulting in a large velocity scatter. The measured velocity amplitude (K=35 km s$^{-1}$) leads to the same conclusion as CSCH, that the compact star must be massive, with M$_X>$4M$_{\odot}$. If instead the velocity amplitude of the compact star is shown by the Balmer absorption lines (K=73 km s$^{-1}$), then the resulting mass diagram is almost identical to that of the supersoft binary SMC 13 (Crampton et al.\ 1997) and lower masses are determined. We have argued above that the measured He II velocity amplitude may be somewhat lower than the actual velocity of the compact star, because the phasing of absorptions will decrease the apparent velocity extremes, while non-orbital motions within the system add at other phases. We point out that the He II velocities in CSCH, when the eclipse was deeper, gave K=40 km s$^{-1}$. The Balmer-line velocities are also likely to contain some non-orbital motions. Thus K=73 km s$^{-1}$ is a reasonable upper limit for the motion of the degenerate star. Figure~\ref{mass} shows the resulting masses corresponding to these two extreme K values. The masses are constrained by the fairly well-known inclination ($i\sim70^{\circ}-80^{\circ}$) and by the requirement that the mass-losing star fills its Roche lobe. No main sequence star fills the Roche lobe defined by these plots, for orbital inclinations larger than $\sim$30$^o$. Thus, the mass-losing star must have evolved off the main sequence to cause mass-exchange. The lowest mass that can thus evolve within a Hubble time is $\sim$ 0.4 M$_{\odot}$, and these are marked in the diagram at the appropriate inclination value. These therefore correspond to the lowest X-ray star masses under these assumptions, and just allow a white dwarf in the K=73 km s$^{-1}$ case. It may be possible that the mass-losing star has very low mass by having lost most of it in some earlier event, and is still filling its Roche lobe now. However, in any of these cases the X-ray star is more massive and the resulting masses are not those expected by the `standard' model (e.g. van den Heuvel et al.\ 1992) in which it is assumed that the compact star is a $\sim$1M$_{\odot}$ white dwarf and the donor star has a mass of $\sim$2.0 M$_{\odot}$. The spectrum in eclipse shows no signs of a late type spectrum (see Figures~\ref{med},~\ref{binned},~\ref{bin2}), consistent with the absence of an evolved more massive secondary. Unfortunately, the complex changes in the disk spectrum make it difficult to be more definitive than this at present, but it appears advisable to revisit the evolutionary scenarios to accomodate a mass-losing star that is less massive than the compact star, not only in CAL 87 but also for the other supersoft X-ray binaries. We discuss separately (Cowley et al.\ 1998) the overall mass determinations for a number of supersoft binary systems. Some of the supersoft X-ray binaries have been found to have highly displaced lines indicating the presence of `jets'. These are most noticable in two low-inclination systems RX J0513$-$69 (Crampton et al.\ 1996) and CAL 83 (Crampton et al.\ 1987) where the displacements are several thousand kilometers per second. In the intermediate inclination system RX J0019$+$22 the jets show a velocity of $\sim\pm$800 km s$^{-1}$ from the central emission line. The jet lines move with the same phase and velocity amplitude as the central line, suggesting that the motion of He II 4686\AA\ emission is indeed orbital. Since CAL 87 is viewed nearly edge-on, any line emission coming from such jets would have a very low radial velocity and thus not be separated from the central emission profile. The asymmetry seen in the He II 4686\AA\ line of CAL 87 could arise from a pair of shifted lines which have larger amplitude (or possibly slightly different phase) than the central line. This would imply that the central line motion is not entirely orbital, but without better resolution this suggestion cannot be verified. | 98 | 4 | astro-ph9804229_arXiv.txt |
9804 | astro-ph9804190_arXiv.txt | We report on Westerbork 1.4 GHz radio observations of the radio counterpart to $\gamma$-ray burst GRB~970508, between 0.80 and 138 days after this event. The 1.4 GHz light curve shows a transition from optically thick to thin emission between 39 and 54 days after the event. We derive the slope $p$ of the spectrum of injected electrons (d$N$/d$\gamma_{\rm e}\propto\gamma_{\rm e}^{-p}$) in two independent ways which yield values very close to $p=2.2$. This is in agreement with a relativistic dynamically near-adiabatic blast wave model whose emission is dominated by synchrotron radiation and in which a significant fraction of the electrons cool fast. | The peak luminosities of $\gamma$-ray bursts (GRBs) are highly super-Eddington and require relativistic outflows (Paczy\'{n}ski 1986; Goodman 1986). Paczy\'{n}ski and Rhoads (1993) pointed out that radio emission is expected as a result of the interaction between such a relativistic outflow and an external medium, as is, e.g., observed in extragalactic jet sources (see also Katz 1994; M\'{e}sz\'{a}ros and Rees 1997). They estimated that the strongest GRBs may be followed by transient ($\sim 10$ mJy) radio emission at intervals ranging from minutes (for a distance $d$ $\sim$ 10$^{5}$ pc) to several weeks ($d$ $\sim$ 10$^{9}$ pc). However, searches for radio counterparts through follow-up observations (Frail et al. 1994,1997a; Koranyi et al. 1995; Galama et al. 1997a,b) were without success until recently. With the rapid accurate location capability of the Wide Field Cameras (WFCs; Jager et al. 1995) onboard the Italian-Dutch X-ray observatory BeppoSAX (Piro et al. 1995) it has recently become possible to detect fading X-ray, optical and radio counterparts to GRBs (Costa et al. 1997a; Piro et al. 1997a; Groot et al. 1997a,b; Van Paradijs et al. 1997; Galama et al. 1997c,1998a; Sahu et al. 1997; Bond 1997; Metzger et al. 1997; Frail et al. 1997b,c; Bremer et al. 1998; Halpern et al. 1997). These observations have settled the discussion on the GRB distance scale (`galactic halo' versus `cosmological', see e.g. Fishman and Meegan 1995, Lamb 1995, Paczy\'nski 1995): GRBs occur at Gpc distances. GRB 970508 is the first GRB to be detected in the radio (Frail et al. 1997b,c); the radio source position coincides with that of the optical (Bond 1997) and X-ray (Piro et al. 1997b) afterglow sources. Assuming that the variations of the source at 4.86 and 8.46 GHz are due to interstellar scintillation (ISS), their damping with time is consistent with a highly relativistically expanding shell passing a diameter of $\approx 3\mu$as (Frail et al. 1997b). VLBI observations show that the source is unresolved ($<$ 0.3 mas, Taylor et al. 1997). We here report on the results of 1.4 GHz radio observations of GRB 970508, made with the Westerbork Synthesis Radio Telescope (WSRT) between 0.80 and 138 days after the burst occurred. | } The observed optical spectral slope $\alpha$ and the optical power law decay of the light curve $F_{\nu} \propto t^{\delta}$ is not consistent with the expected relation for the simplest blast wave model ($\delta = 3\alpha/2$; e.g. Wijers, Rees and M\'esz\'aros 1997). The observed power law decay value, $\delta = -1.141 \pm 0.014$ ($t > 2$ days, Galama et al. 1998b; see also Pedersen et al. 1998, Castro-Tirado et al. 1998, Sokolov et al. 1998) would imply $\alpha = -0.761 \pm 0.009$, while in the optical passband $\alpha = -1.12 \pm 0.04$ is observed (Galama et al. 1998c, from here on Paper II; see also Sokolov et al. 1998). In the following we show that this may be explained by rapid cooling of a significant fraction of the electrons. A population of electrons with a power-law distribution of Lorentz factors $\gamma_{\rm e}$ (d$N$/d$\gamma_{\rm e}\propto\gamma_{\rm e}^{-p}$) above some minimum value $\gamma_{\rm m}$ emits a power law synchrotron spectrum above the frequency $\nu_{\rm m}$ (corresponding to radiating electrons with $\gamma_{\rm m}$; e.g. Rybicki \& Lightman 1979). Independently, above some Lorentz factor $\gamma_{\rm c}$ the electrons may cool rapidly, and an extra break in the spectrum is expected at the corresponding frequency $\nu_{\rm c}$ (Sari, Piran, \& Narayan 1998). Beyond a certain time $t_0$ ($t_0$ is small $\sim$ 500 sec; see Paper II) the evolution of the blast wave is adiabatic (Sari et al. 1998 and see Paper II); then $\nu_{\rm m} < \nu_{\rm c}$, and the spectrum varies as $F_{\nu} \propto \nu^{-(p-1)/2}$ from $\nu_{\rm m}$ up to $\nu_{\rm c}$; above $\nu_{\rm c}$ it follows $F_{\nu} \propto \nu^{-p/2}$ and below $\nu_{\rm m}$ it follows the low frequency tail, $F_{\nu} \propto \nu^{1/3}$ (Sari et al. 1998). The evolution in time of the GRB afterglow is determined by the evolution of these break frequencies: $\nu_{\rm c} \propto t^{-1/2}$ and $\nu_{\rm m} \propto t^{-3/2}$ (both decrease with time). The decay part of the optical R$_{\rm c}$ (Coussins R) band light curve (in the optical passband $\nu > \nu_{\rm c}$ for $t\gsim$ 1.2 days; Paper II) goes as $F_{\nu} \propto t^{(2-3p)/4}$, while the spectrum is then $F_{\nu} \propto \nu^{-p/2}$ (Sari et al. 1998). This allows us to make two independent measurements of $p$: using $F_{\rm R_{\rm c}}\propto t^{-1.141 \pm 0.014}$ we find $p= 2.188 \pm 0.019$ and using $\alpha_{\rm opt} = -1.12 \pm 0.04$ gives $p = 2.24 \pm 0.08$. The excellent agreement between the values of $p$ supports that a significant fraction of the electrons cool rapidly and that the evolution of the GRB remnant is adiabatic. Additional evidence for rapid cooling of a significant fraction of the electrons is given in Paper II. Observations by Bremer et al. (1998) with the IRAM Plateau de Bure Interferometer (PdBI) at 86 GHz show a maximum around $\sim$ 12 days. We identify this maximum with the break frequency $\nu_{\rm m}$ passing 86 GHz at $t_{\rm m,86 GHz} \sim$ 12 days (Paper II). We expect, the 8.46 and 1.4~GHz emission to peak at $t_{\rm m,8.46 GHz} \sim$ 55 days and $t_{\rm m,1.4 GHz} \sim$ 180 days, respectively ($\nu_{\rm m} \propto t^{-3/2}$). Near day 55, a shallow maximum can be seen in the 8.46 GHz light curve (Frail et al. 1997b). Unfortunately our 1.4 GHz light curve cannot be used to test the presence of the maximum at that frequency, both due to low signal to noise and because it ends 150 days after the burst, i.e. before the predicted maximum. Before $\nu_{\rm m}$ passes 8.46 GHz at $t_{\rm m,8.46 GHz}$ we expect the 8.46 GHz spectrum to follow the low frequency tail $F_{\nu} \propto \nu^{1/3}$, while after $t_{\rm m,8.46 GHz}$ it is expected to be $F_{\nu} \propto \nu^{-(p-1)/2} = \nu^{-0.6}$ (Sari et al. 1998 and we have used $p$ = 2.2). Thus, we predict a gradual transition between $t_{\rm m,8.46 GHz} \sim$ 55 days and $t_{\rm m,4.86 GHz} \sim$ 80 days (when also at 4.86 GHz $\nu_{\rm m}$ has passed) from $\alpha$ = 1/3 to $\alpha = -0.6$. We note that this expectation is different from blast wave models that do not include the effect of rapid cooling of a significant fraction of the electrons ($F_{\nu} \propto \nu^{-1.1}$ similar to the optical slope; see e.g. Wijers et al. 1997). Also the decays at 8.46 GHz (after $t_{\rm m,8.46 GHz} \sim$ 55 days) and 4.86 GHz (after $t_{\rm m,4.86 GHz} \sim$ 80 days) are expected to be different from that in the optical and X-ray passbands, $F_{\nu} \propto t^{3(1-p)/4} = t^{-0.9}$; where we have used $p$ = 2.2). These predictions can be tested with the continued monitoring at the VLA at 4.86 and 8.46 GHz by Frail et al. (1998). The radio afterglow light curves of GRB 970508 (Frail et al. 1997b and this Letter) show a much more gradual evolution than expected (see e.g. the fit to the 8.46 and 4.86 GHz data by Waxman, Kulkarni and Frail 1998). Also a constant self-absorption frequency was expected (e.g. Waxman et al. 1997) while we here show that a transition from optically thick to thin emission occurred around $\sim$ 45 days. For $t < t_o$ Sari et al. (1998) predict a decrease with time of the self-absorption frequency, $\nu_{\rm a}$, while for $t >$ $t_0$ the self-aborption frequency $\nu_{\rm a}$ remains constant. The transition from optically thick to thin 1.4 GHz radiation then suggests that $t_0$ $\sim$ 45 days. Also the 8.46 GHz light curve (Frail et al. 1997b) suggests that $t_0$ cannot be much smaller than 10 days, i.e. 10 days $\lsim$ $t_0$ $\lsim$ 55 days (we have extrapolated backwards in time from the 8.46 GHz peak at $t_{\rm m}$ $\sim$ 55 days with the expected dependence $F_{\nu} \propto t^{1/2}$ for times $t < t_{\rm m}$). This is not in agreement with the finding that $t_0 \sim 500$ sec (Paper II). However, the absence of a break in the smooth power law decay of the optical light curve from 2 to 60 days after the burst (Pedersen et al. 1998; Castro-Tirado et al. 1998; Sokolov et al. 1998; Galama et al. 1998c) shows that there is no important transition in that period. This does imply that some additional ingredient is needed; for example, Waxman et al. (1998) argue that the transition from ultrarelativistic to mildly relativistic expansion of the blast wave may explain the decrease in the self-absorption frequency $\nu_{\rm a}$ with time and the slow time dependence of the early radio light curves. The excellent agreement in the derived value for $p$ ($p$ = 2.2) from the decay of the optical light curve and the optical spectral slope support an adiabatic dynamical evolution of the GRB remnant and an extra break in the synchrotron spectrum at the frequency $\nu_{\rm c}$ above which the radiation is from electrons which cool rapidly compared to the remnant's expansion time. We predict a transition in the radio spectral index $\alpha_{\rm 4.86-8.46 GHz}$ from 1/3 to --0.6, between 55 and 80 days; the light curves are predicted to decay as $F_{\nu} \propto t^{-0.9}$ after 55 days at 8.46 GHz and 80 days at 4.86 GHz. | 98 | 4 | astro-ph9804190_arXiv.txt |
9804 | astro-ph9804159_arXiv.txt | We present observations of \htwo\ fluorescence at wavelengths between 1000 and 1200 \AA\ from the bright reflection nebula IC 63. Observations were performed with the Berkeley spectrograph aboard the \orf-II mission \cite{Hetal98}. To the best of our knowledge, this is the first detection of astrophysical \htwo\ fluorescent emission at these wavelengths (excluding planetary atmospheres). The shape of the spectrum is well described by the model of \citeN{S89}. The absolute intensity, however, is fainter than an extrapolation from observations at longer ultraviolet wavelengths \cite{WSBB89} by a factor of ten. Of the mechanisms that might help reconcile these observations, optical depth effects in the fluorescing \htwo\ itself are the most promising (or at least the most difficult to rule out). | In many environments, the equilibrium abundance of the hydrogen molecule (\htwo) is determined by a balance between formation on dust grains and photodissociation pumped by photons at wavelengths below about 1108 \AA\ \cite{SB82}. Photons between this threshold energy and the Lyman limit populate a large number of rovibrational states in the $B^1 \Sigma_u^+$ and $C^1 \Pi_u$ electronic levels. En route back to the $X^1 \Sigma_g^+$ level, the excited \htwo\ may radiate through any of several bound-bound and/or bound-free channels leading to a complex fluorescence emission spectrum. Located about 20$\arcmin$ or 1.3 pc from $\gamma$ Cas, dense gas in the reflection nebula IC 63 is illuminated by a bright UV radiation field. Line widths of a variety of molecular species are quite narrow \cite{JVDB94}, suggesting that shock excitation is comparatively unimportant in this environment. The nebula provides an excellent test case for models of \htwo\ fluorescence and other photochemical processes \cite{JVDBSS95}. The fluorescence model described in \citeN{S88} and \citeN{S89} has been successfully applied in the interpretation of infrared \htwo\ fluorescence from the reflection nebulae IC 63 and IC 59 \cite{LLBJF97}, and to the ultraviolet fluorescence in the band near 1600 \AA\ from IC 63 \cite{S89} and the Taurus molecular cloud \cite{H94}. In this work we test the model in a previously unexplored band deep in the ultraviolet where the significant majority of the radiated fluorescent power is expected to emerge. | \label{conclusions} We have detected \htwo\ fluorescence at wavelengths between 1000 and 1200 \AA\ from the bright reflection nebula IC 63 with the Berkeley spectrograph aboard the \orf-II mission \cite{Hetal98}. The wavelengths and relative strengths of the fluorescent features within the ORFEUS band agree well with the predictions of the model of \citeN{S89}. The absolute fluorescent intensity is fainter than an extrapolation from observations at longer ultraviolet wavelengths \cite{WSBB89} by a factor of ten. The measurements can not be reconciled by differential extinction from a foreground slab of dust (presumably associated with neutral gas) nor by absorption from quiescent \htwo. Optical depth effects in the fluorescing \htwo\ itself, predicted by \citeN{WSBB89}, remain the most plausible mechanism to explain our observations. | 98 | 4 | astro-ph9804159_arXiv.txt |
9804 | astro-ph9804315_arXiv.txt | We present contemporary infrared and optical spectra of the plateau type~II SN~1995V in NGC~1087 covering four epochs, approximately 22 to 84 days after shock breakout. The data show, for the first time, the {\it infrared} spectroscopic evolution during the plateau phase of a typical type~II event. In the optical region P~Cygni lines of the Balmer series and of metals such as Sc~II, Fe~II, Sr~II, Ca~II and Ba~II lines were identified. The infrared (IR) spectra were largely dominated by the continuum, but P~Cygni Paschen lines and Brackett~$\gamma$ lines were also clearly seen. The other prominent IR features are confined to wavelengths blueward of 11000~\AA\ and include Sr~II 10327, Fe~II 10547, C~I 10695 and He~I 10830~\AA. Helium has never before been unambiguously identified in a type~IIp supernova spectrum during the plateau phase. We demonstrate the presence of He~I 10830~\AA\ on days 69 and 85. The presence of this line at such late times implies re-ionisation. A likely re-ionising mechanism is $\gamma$-ray deposition following the radioactive decay of $^{56}$Ni. We examine this mechanism by constructing a spectral model for the He~I 10830~\AA\ line based on explosion model s15s7b2f of Weaver \& Woosley (1993). We find that this does not generate the observed line owing to the confinement of the $^{56}$Ni to the central zones of the ejecta. In order to reproduce the He~I line, it was necessary to introduce additional upward mixing or ``dredge-up'' of the $^{56}$Ni, with $\sim$10$^{-5}$ of the total nickel mass reaching above the helium photosphere. In addition, we argue that the He~I line-formation region is likely to have been in the form of pure helium clumps in the hydrogen envelope. The study of He~I 10830~\AA\ emission during the photospheric phase of core-collapse supernovae provides a promising tool for the constraint of initial mixing conditions in explosion models. | Type~IIp (plateau) supernovae form the classic subgroup of the core-collapse supernovae. They are believed to arise from massive stars (12-25 M$_{\odot}$) during the red supergiant phase. Early theoretical work by Falk and Arnett (1973) showed that hydrodynamical instabilities should appear in explosions of such massive stars. As the shock wave propagates through the stellar envelope it sets up density and pressure profiles which can in some cases result in the formation of Rayleigh-Taylor~(RT) instabilities (Chevalier 1976). In particular, RT instabilities are expected to grow at the interface of the core and the hydrogen envelope because of the large entropy (and density) jump that occurs there (Weaver \& Woosley 1980). A direct consequence of these instabilities is that chemical mixing in the ejecta takes place (Bandiera 1984). Herant \& Woosley (1994) have studied 2-D simulations of red supergiant explosions over a wide mass range, and found that the growth of hydrodynamic instabilities is highly likely in all cases. They showed that as the explosion (outgoing) shock plows into the hydrogen envelope, a reverse (ingoing) shock is formed, and between them, RT instabilities grow. Bubbles of hydrogen formed by these instabilities are violently dragged towards the centre of the star by the reverse shock. Simultaneously, compact helium and oxygen clumps advance out into the hydrogen envelope and bubbles of $^{56}$Ni are formed and distributed in the outer parts of the ejecta. It has also been realised that strong dredge-up should result from the neutrino-driven convection close to the neutron star surface (Herant \& Benz 1992) which has been invoked to account for the conversion of the core-collapse to explosion. This should also produce fast-moving blobs or ``fingers'' of radioactive material which eventually penetrate the outer layers of the supernova. So far, only SN~1987A has provided us with clear observational evidence for dredge-up in a core-collapse event. This includes the shape of the light curve, the early detection of X-rays and $\gamma$-rays and the width of the iron lines in the infrared. However SN~1987A was, of course, only a single event, and a rather unusual one in that it arose from a blue supergiant progenitor. Therefore, we cannot simply assume that similar dredge-up occurs in all other type~II events. Indeed, simulations have shown that differences in progenitor structure can lead to significantly modified hydrodynamical evolution (Herant \& Benz 1992; Herant \& Woosley 1994). Clearly, to establish whether or not deep dredge-up is typical of all core-collapse supernovae, a major step would be to demonstrate dredge-up and mixing in the most-common of all core-collapse events, the type~IIp supernova. A powerful demonstration of the occurrence of deep dredge-up would be the appearance of radioactive material at the surface at early times. Helium lines arising in the supernova envelope can be used as a tracer of the upwardly-mixed radioactive material. Helium lines are of high excitation. During the first week of the supernova, recombination maintains the populations of the excited levels and so He~I lines are seen. However after about 10 days, the conditions in the type~IIp atmosphere are such that all the helium will have recombined and de-excited to the ground state. But, if dredge-up occurs during the explosion, radioactive $^{56}$Ni may reach the outer parts of the supernova envelope at early times. If it does, the $\gamma$-rays from its decay ($^{56}$Ni $\Rightarrow$ $^{56}$Co $\Rightarrow$ $^{56}$Fe) will excite or re-ionise the helium. Thus the detection of helium lines during the plateau phase (20-120~d post explosion) should imply upward mixing of radioactive material from the core. Unfortunately no optical He~I lines have ever been unambiguously identified during the plateau phase. However, there are two well-known strong lines in the infrared viz. He~I 10830~\AA\ (2s$^{3}$S--2p$^{3}$P$^{0}$) and 20580~\AA\ (2s$^{1}$S--2p$^{1}$P$^{0}$) which offer the prospect of testing for dredge-up of radioactive material. This technique was applied by Graham (1988) and Chugai (1991) using the 10830~\AA\ line in the SN 1987A at early times. Lucy (1991) invoked upward mixing of $^{56}$Ni to account for strong optical He I lines in the early-time spectra of type Ib supernovae. To investigate dredge-up in type~IIp supernovae we began in 1995 a programme of infrared and optical spectroscopy of this type of supernova. The data we present here comprise an extensive set of IR/optical spectroscopic observations of the type~IIp SN~1995V, spanning epochs of 22 to 85 days post explosion. The observations are described in section 2. In section 3 we compare the observations with a simple spectral synthesis model and discuss the line identifications, especially in the IR. In section 4 we describe the method we used to estimate the amount of dredge-up of $^{56}$Ni. In section 5 we present the results from the comparison of the model with the data, and our estimations for the amount of dredge-up. In section 6 we discuss the implications of this work for our understanding of dredge-up. | Reproduction by explosion model s15s7b2f of the observed He~I emission is achieved only by invoking substantial additional dredge-up of the $^{56}$Ni from the core. As expected, the case with no helium clumping requires the greatest dredge-up. However, we argue that the no-clumping case is probably highly unrealistic, since it is difficult to see how one could achieve such a large amount of $^{56}$Ni dredge-up and yet have no pure helium clumps in the same environment. We deduce, therefore, that there must exist some pure helium clumps in the hydrogen envelope. As we have shown, with fairly modest clumping of $\chi_{He}=$0.1 or 0.2, the uniform central core coincides, to within the errors with that of the unmodified s15s7b2f model. The addition of a steep, power-law density component to this will bring sufficient $^{56}$Ni to the surface to account for the He~I 10830~\AA\ emission. Nevertheless, the fraction of the $^{56}$Ni mass which must be dredged up beyond the uniform core is quite substantial. For SN~1995V we conclude that a) a small amount of $^{56}$Ni ($\sim$10$^{-6}$ M$_{\odot}$) must have been dredged up to the helium photosphere (v$\sim$4,250~km/s), and b) clumps of pure helium must have also existed in this region. High velocities for the decay products of $^{56}$Ni ($\sim$ 3000 km/s) were also observed in the ejecta of SN 1987A (e.g. Meikle et al. 1993). As shown by Herant \& Benz (1992), if the $^{56}$Ni is located at the base of the ejecta at t$\sim$300s it is impossible to accelerate even a small fraction to about 3000 km/s during subsequent instabilities. In order to achieve such high velocities, it is necessary to invoke outward mixing of the nickel at even earlier times, such as might be caused by neutrino convection. If this occurred, then the later instabilities would carry the nickel to still higher velocities. In order to match the observations of SN1987A, Herant \& Benz had to premix nickel out to 1.5 M$_{\odot}$ above the mass cut. Herant \& Woosley (1994) studied shock propagation, mixing and clumping in the explosion of red supergiants. In order to take into account the pre-mixing of $^{56}$Ni during the initiation of the explosion, they diluted the nickel by a factor $\sim$4 above the mass cut. They then followed the shock propagation and the growth of RT instabilities. For all progenitors their simulations showed that extensive RT instabilities develop in the ejecta in the wake of the reverse shock from the H/He interface. In contrast to the blue supergiant studies, these instabilities have ample time in which to evolve and completely reshape the ejecta. In spite of this, in all the explosions simulated, nickel did not reach velocities higher than $\sim$1500 km/s. Similarly, helium did not exceed velocities higher than 2500 km/s. Our results, therefore, indicate that a higher degree of pre-mixing may be required than is invoked in the Herant \& Woosley models. Recently, Bazan \& Arnett (1997) have simulated mixing in core-collapse events. Their simulations include both RT and Richtmeyer-Meshkov instabilities. These produce much higher velocities for $^{56}$Ni than do R-T instabilities alone. Velocities as high as $\sim$4000~km/s are predicted. The mass and profile index of the upwardly-mixed $^{56}$Ni derived above could be of considerable value in constraining the initial parameters of these instabilities. | 98 | 4 | astro-ph9804315_arXiv.txt |
9804 | astro-ph9804123_arXiv.txt | s{ The X-ray emission from clusters of galaxies is one of the most pursued observational probe to investigate the distribution of dark matter and the related density parameter $\Omega_0$. The crucial link to derive the statistics of observables from a dynamical theory is constituted by the physics for the diffuse baryons (or ICP) responsible of the X--ray emission. Here we present a physical model for the ICP which leads to a definite $L$--$T$ relation. Then we perform a physically based cosmological test, pointing out three cold dark matter universes: a Tilted critical CDM, a flat CDM with $\Omega_0=0.3$, and an Open CDM with $\Omega_0=0.5$, which are discussed on the basis of the RDCS survey. } | Groups and clusters of galaxies constitute cosmic structures sufficiently close to equilibrium and with sufficient density contrast ($\delta\approx 2\, 10^2$ inside the virial radius $R$) as to yield definite observables. They are dominated by dark matter (hereafter DM), while the baryon fraction is observed to be less than $20$\%. The great majority of these baryons are in the form of {\sl diffuse plasma} (ICP) with densities $n\sim 10^{-3}$ cm$^{-3}$ and virial temperatures $k\,T\sim 5$ keV, and are responsible for powerful X--ray luminosities $L\sim 10^{44}$ erg/s by optically thin thermal bremsstrahlung. As the plasma is a good tracer of the potential wells, much better than member galaxies, the X--ray emission is a powerful tool to investigate the mass distribution out to moderate and high redshifts. The ICP temperature directly probes the height of the potential well, with the baryons in the role of mere tracers; on the other hand, the luminosity with its strong dependence on density ($L\propto n^2$) reliably probes the baryonic content and distribution. Statistically, an average $L$--$T$ correlation is observed along with substantial scatter, and this provides the crucial link to relate the X-ray luminosity functions with the underlying statistics of the DM. A physical model for the diffuse baryons is difficult to achieve. In fact, the simple self similar model (Kaiser 1986), which assumes the ICP amount to be proportional to the DM's at all $z$ and $M$, leads to a relation $L\propto T^2$, conflicting with the observed correlation for rich clusters. The latter is close to $L\propto T^{3.5}$ (David et al. 1993; Mushotzky \& Scharf 1997). Here we propose a physical model for baryons, which leads to a prediction for the $L$--$T$ relation (see Cavaliere, Menci \& Tozzi 1997, CMT97) and allows a non parametrical approach to the search for cosmological parameters. The results, presented in \S 3, are a synthesis from Cavaliere, Menci \& Tozzi (1998, CMT98). | We presented a physically based approach to cosmological tests with clusters of galaxies. We describe the X--ray emission from clusters with a specific model for the diffuse baryons. In this sense this approach is alternative to the parametrical approach (see Borgani this meeting). The results of our model depend on two parameters, the external temperature $T_1$ and density $n_1$, which are {\sl not} free. Specifically, we use for $T_1$ the range $0.1\div 0.8 $ keV provided by the stellar preheating. The value of $n_1$ for rich clusters is related to the DM density by the universal baryonic fraction. Thus we compute the expression of the bolometric luminosity for a given temperature. The average of the square of the density jump factor $\langle g^2 \rangle$ over the merging histories coupled with $\beta(T)$ is what gives to the statistical $L-T$ correlation the curved shape shown in fig. \ref{lt}b. In addition, our approach predicts an intrinsic {\it variance} of dynamical origin due to the different merging histories, and built in the factor $g^2$. With the ICP state so described, we proceeded to constrain the cosmological parameters. After the observations by Rosati et al. (1998), we have computed the X-ray observables for groups and clusters of galaxies. On the basis of local data, the set of acceptable CDM universes is restricted to three disjoint domains: $\Omega=1$ for the Tilted CDM with high baryon content; $\Omega_o\simeq 0.5$ for standard CDM; $\Omega_o\approx 0.3$ for CDM in flat geometry. However, only the TCDM and the $\Lambda$CDM universes give acceptable faint counts. As an overall remark, a common feature of all the above universes is constituted by some excess in the counts. This may indicate some non--trivial incompleteness in the canonical hierarchical clustering, worth keeping under scrutiny. We recall that in the adiabatic models for the ICP (Evrard \& Henry 1991, Kaiser 1991) the evolution of the $L-T$ relation is reduced or even negative, thus alleviating the excess. However, the anti--evolution required in OCDM would be very difficult to justify (the adiabatic models are largely discussed in CMT98). Now the question is: to what extent enlarging the data base on X--ray clusters will help in further constraining cosmology? We argue that the variance intrinsic to the hierarchical clustering, and amplified by the ICP emissivity, sets an effective limitation. Richer, faint surveys will hardly provide a sharper insight into cosmology unless one reduces both the uncertainty concerning $\sigma_8$ and the larger one concerning $L_{o}$. However, we stress that such efforts will find soon a more proper aim than constraining $\Omega_o$. This is because MAP, and subsequently PLANCK, will accurately measure on still linear scales not only the perturbation power spectrum but also directly $\Omega_o$. Once the cosmological framework has been fixed, the study of groups and clusters in X-rays will resume its proper course, that is, the physics of systems of intermediate complexity which is comprised of the DM and of the ICP component. | 98 | 4 | astro-ph9804123_arXiv.txt |
9804 | astro-ph9804065_arXiv.txt | We present an empirical method which measures the distance to a Type Ia supernova (SN~Ia) with a precision of $\sim$ 10\% from a single night's data. This method measures the supernova's age and luminosity/light-curve parameter from a spectrum, and the extinction and distance from an apparent magnitude and color. We are able to verify the precision of this method from error propagation calculations, Monte Carlo simulations of well-sampled SNe~Ia, and the Hubble diagram of scarcely observed supernovae. With the reduction in telescope time needed, this method is three to four times more efficient for measuring cosmological parameters than conventional light-curve based distance estimates. | The explosion of a Type Ia supernova (SN~Ia) is a catastrophic phenomenon veiled in layers of complexity. Recent efforts to monitor these events have led to an increased ability to predict, if not fully understand, the stages of SN Ia evolution. The model for the photometric history of SNe~Ia has been refined from a homogeneous description \cite[]{brunophd,branch_miller,ts95} to one which characterizes a relation between peak luminosity and light-curve shape \cite[]{philm15,hametal95,hametal96,rpk95,kim_stretch98}. The slower, broader light-curves are intrinsically brighter at peak than the faster, narrower light-curves. Recognizing and exploiting such relations has led to a renaissance in the use of SNe~Ia as extragalactic distance indicators. Extending luminosity/light-curve relations to multiple passbands separates the competing effects of dust, intrinsic differences, and distance on the light of SNe~Ia \cite[]{rpk96}. Distances with 5-10\% uncertainty can be obtained using the light-curve shapes of well-observed supernovae. The optical spectra of SNe~Ia are rich in information [see \cite{fil97} for a review]. Many of the elements synthesized and ejected in the explosion have been identified despite the blending of their high-velocity profiles \cite[]{bran81b,nugphd}. In addition, the relative strengths of some spectral features have been shown to correlate with SN~Ia peak luminosity \cite[]{nugseq95}. As the supernova evolves, predictable casts of features appear and disappear, illuminated by the photosphere's recession through the synthesized layers. The temporal evolution of these features is sufficiently reliable to be used as a clock to determine the current age of a SN~Ia to a precision of 1-2 days \cite[]{mink39,riess_age97}. Unfortunately, supernovae occur without warning, making it difficult to collect the observations necessary to measure their distances. Observing an unscheduled event in up to four filters many times over the course of $\sim$ 100 days is a time consuming and logistically formidable task. The observing record for a typical SN~Ia is quite fragmentary. Following this process, it will be many years of work to gather the number of SN~Ia distances necessary to put strong limits on cosmological parameters. Even at high redshifts ($z \geq 0.3$), where a strategy for batch detections of multiple supernovae has made it possible to schedule supernova discoveries and their follow up \cite[]{perl97}, difficulties arise. Since these observations require the largest telescopes, the light-curves are typically more poorly sampled than the nearby ones caught at a similar phase. Yet, from a single night's observations, a SN~Ia's spectrum and photometric magnitudes can reveal its age, intrinsic luminosity, extinction, and apparent brightness. From this information one can estimate the distance to a supernova without further observations (except for the possible need of a galaxy image to subtract the host's light). Here we explore this possibility with two independent sets of SNe~Ia. We describe this technique in \S 2 and its expected uncertainty in \S 3. In \S 4 we apply it to randomly selected snapshots of extensively observed SNe~Ia to determine the precision of such distance estimates. In \S 5 we construct the Hubble diagram of ``cast-off'' SNe~Ia: objects which were observed only once or a few times. We extend the application of this method in \S 6 to SNe~Ia with $0.2 \leq z \leq 0.83$. In \S 7 we discuss variations of this technique and its leverage on estimating cosmological parameters. | In principle, enough information can be garnered from a single supernova spectrum and photometric epoch to estimate the distance to a SN~Ia. In practice, the results of \S 3, 4, and 5 suggest this method produces distances having a precision of $\sim$ 10\%, with variations that are a function of the quality of the data and the age of the supernova. Depending on the amount of host galaxy contamination, it may be necessary to obtain spectra and images of the host galaxy after the supernova has faded. The snapshot distance method employs the same luminosity and extinction corrections used in the MLCS method of Riess, Press, and Kirshner 1996a (and more recently updated in Riess et al. 1998b). We find no significant offset between the distance estimates of the two methods and only a modest reduction in precision for the snapshot distance method. There are three more limited versions of a distance method using single epoch SN Ia observations which reveal the utility of luminosity and extinction corrections (see Table~\ref{discomp_tab} and Figure~\ref{show}). These variants employ the SFA measurement but lack the luminosity correction, the extinction correction, or both. Disregarding individual luminosities and light-curve shapes predicted by the spectral ratios, we fit homogeneous, fiducial templates to the photometric epoch. Further, we discarded our estimate of the extinction from the color excess. We measured the resulting ``standard candle'' distances to SNe~Ia using the Monte Carlo technique described in \S 4. As seen in Figure~\ref{show}, the distribution of dispersions has a mean of 0.35 mag, a value consistent with previous SNe~Ia distance estimates which assume SN Ia light curve homogeneity and do not correct for extinction \cite[]{st93,ts95,hametal95,hametal96,rpk95,rpk96,branch_miller,vauetal95}. These distances are also 15\% {\it greater} in the mean (or smaller in the implied Hubble constant) than either MLCS or snapshot distances, consistent with other comparisons of distance estimates which assume homogeneity instead of heterogeneity of SNe~Ia \cite[]{ts95,hametal95,hametal96,rpk95,rpk96}. To simulate the effect of single filter information, we used our luminosity correction without an extinction correction. This procedure results in a mean dispersion of 0.25 mag and distances which are only 3\% greater than those obtained from MLCS. This result, though better than the standard candle method, is still worse than the complete snapshot procedure. A final variant is to disregard a luminosity/light-curve shape correction but maintain an extinction correction from the color excess. Such a method using light-curves was proposed by \cite{vdb95} to account for both intrinsic luminosity differences as well as absorption by dust. This method takes advantage of the coincidental near-agreement between the standard reddening law and the relation between intrinsic color and luminosity to correct for both extinction and luminosity differences. \cite{rpk96} have noted that while both sources of luminosity variation affect the SN color in the same direction, the specific ratios of the luminosity difference to color difference are not precisely the same for extinction and intrinsic SN Ia variation. Monte Carlo simulations of this method combined with a SFA measurement give a mean dispersion of 0.19 mag and distances which are 5\% greater in the mean than those of the MLCS method. Despite the low dispersion of this method, we are suspicious of the distances it predicts. The distribution of dispersions obtained from our Monte Carlo simulation (see Figure~\ref{show} and Table~\ref{discomp_tab}) is more skewed than any other, including an asymmetric ``tail'' encompassing dispersions greater than 0.3 mag. We believe that for SNe~Ia with only moderate amounts of extinction or whose luminosities are similar to those of typical SNe~Ia, this method has merit. Yet for very red SNe~Ia, this method can predict distances which are systematically and considerably in error due to the inability to distinguish between absorption by dust and intrinsic variation. The snapshot method predicts distances which agree in the mean with only a moderate reduction in precision from light-curve shape methods. Yet because of the greatly diminished expense in data collection, this method can be more effective for problems which benefit equally well from a high {\it quantity} of SNe~Ia distances as from the {\it quality} of those distances. Two such applications are mapping the nearby peculiar velocity field and determining the cosmological parameters which dictate global geometry. Recent attempts to map the cosmic velocity field with SNe~Ia \cite[]{riess_beta97} suffer from dilute spatial sampling. Replete peculiar velocity maps could reveal the influence of matter fluctuations and constrain the matter content of the local Universe. Nearby, many more SNe~Ia are discovered than can be regularly monitored. By decreasing the observational requirements of each SN Ia, it should be possible to increase the sampling of the local velocity field. The light-curves employed by \cite{hametal96} and \cite{rpk96} are typically sampled for 10 to 15 epochs. The observational demands increase as the supernova rapidly fades. With the telescope time invested in a single set of SN Ia multi-color light-curves, sufficient data for 15 to 25 SNe~Ia snapshot distances could be gathered. Accounting for the inherent distance uncertainties, telescope time spent collecting snapshot distance data is 3 to 4 times more efficient than time spent collecting light-curves for distance estimates. Efforts to measure the cosmological parameters $\Omega_M$ and $\Omega_\Lambda$ from distant SNe~Ia could also profit from snapshot distances. Systematic searches for SNe~Ia at z $\geq 0.3$ have yielded a plethora of objects \cite[]{perliau95a,schetal95}. Combining a new generation of large telescopes with $\sim$ 1 degree fields of view and multi-fiber spectrometers with the snapshot method could allow SNe~Ia distances to be gathered in batch at an unprecedented rate. At the current rate of discovery, a night spent searching five 1-degree fields followed by a night collecting spectra of the candidates with a multi-fiber spectrometer could yield $\sim$ 50 SNe~Ia distances \cite[]{schmidt97,rate_96}. Repeating this process every new moon could yield up to $\sim$ 600 distances a year. At this rate of accumulation, it should be possible to convincingly separate the effects of various sources of energy density on the redshift-magnitude relation \cite[]{omol_95}. A more optimal method for measuring SN Ia distances would employ both the predictive power of SN Ia light and color curve shapes with that of SN Ia spectra. Such a method would replace the distinction between a snapshot distance and a light-curve distance with a distance estimate which makes the most economical use of all available SN Ia observations. Using the tools described in Riess, Press, and Kirshner 1996a, Nugent et al. 1995, Riess et al. 1997, Riess et al. 1998b, and this paper, such a method appears to be quite feasible. \bigskip This work was supported by the NSF through grant AST--9417213 to A.V.F. and AST-9528899 and AST-9617058 to R.P.K., by the Miller Institute for Basic Research in Science through a fellowship to A.G.R., and by the Director, Office of Computational and Technology Research, Division of Mathematical, Information, and Computational Sciences of the U.S. DoE under contract number 76SF00098 to P.E.N. Some of the calculations presented in this paper were performed at the National Energy Research Supercomputer Center (NERSC), supported by the U.S. DoE. We thank Stephan Benetti and George Djorgovski for allowing us to use their spectra of SNe Ia prior to publication, and Bruno Leibundgut for suggestions that helped improve this paper. \appendix | 98 | 4 | astro-ph9804065_arXiv.txt |
9804 | hep-ph9804285_arXiv.txt | The expected proton and neutrino fluxes from decays of massive metastable relic particles are calculated using the HERWIG QCD event generator. The predicted proton spectrum can account for the observed flux of extremely high energy cosmic rays beyond the Greisen-Zatsepin-Kuzmin cutoff, for a decaying particle mass of ${\cal O}(10^{12})$~GeV. The lifetime required is of ${\cal O}(10^{20})$~yr if such particles constitute all of the dark matter (with a proportionally shorter lifetime for a smaller contribution). Such values are plausible if the metastable particles are hadron-like bound states from the hidden sector of supersymmetry breaking which decay through non-renormalizable interactions. The expected ratio of the proton to neutrino flux is given as a diagonistic of the decaying particle model for the forthcoming Pierre Auger Project. | It has been known for some time that interactions on the 2.73 K blackbody cosmic microwave background (CMB) will severely degrade the energies of cosmic ray nucleons with energies beyond $\sim5\times10^{19}\ev$ --- the Greisen-Zatsepin-Kuzmin (GZK) cutoff \cite{gzk}. It was therefore very surprising when the Fly's Eye atmospheric fluorescence detector reported the observation of an extremely high energy cosmic ray (EHECR) event with an energy of $(3.0\pm0.9)\times10^{20}\ev$ \cite{flyseye}. This was followed by the detection of a $(1.7-2.6)\times10^{20}\ev$ event by the AGASA air shower array \cite{agasa}. These discoveries substantiated earlier claims from the Volcano Ranch \cite{vr}, Haverah Park \cite{hp} and Yakutsk \cite{yak} air shower arrays that cosmic rays do exist beyond the GZK cutoff. About a dozen such events are now known. Detailed accounts of the data may be found in recent reviews \cite{reviews}. In Figure~\ref{fig1} we show the EHECR spectrum for energies exceeding $10^{18}\ev$ \cite{spec}; note that the fluxes have been multiplied by $E^3$. It is believed that cosmic rays with energies up to $\sim5\times10^{18}\ev$, the so-called `ankle', are predominantly of galactic origin, possibly accelerated by the Fermi mechanism in supernova remnants \cite{books}. Above this energy, the spectrum flattens and the composition changes from being mostly heavy nuclei to mostly protons. Such a correlated change in the spectrum and composition was first established by the Fly's Eye experiment \cite{flyseye} and Figure~\ref{fig1} shows their suggested two-component fit to the data. The new component which dominates at energies beyond $\sim5\times10^{18}\ev$ is isotropic and therefore cannot possibly originate in the galactic disk \cite{agasa2,lc95}. However it also extends well beyond the GZK cutoff raising serious problems for hypothetical extragalactic sources. Because of the rapid energy degradation at these energies through photo-pion production on the CMB, such sources must exist within $\sim500\mpc$, in fact within $\sim50\mpc$ for the highest energy Fly's Eye event \cite{cronin}. For heavy nuclei, the energy loss is less severe according to a revised calculation \cite{heavy} so the range may extend upto $\sim100\mpc$. General arguments \cite{greisen,hillas} provide correlated constraints on the magnetic field strength and spatial extent of the region necessary to accelerate particles to such high energies and these requirements are barely met by likely astrophysical sites such as active galactic nuclei and the `hot spots' of radio galaxies \cite{cr}. Moreover there are few such sources close to us and no definite correlations have been found between their locations and the arrival directions of the most energetic events \cite{source,agasa2}. It has been speculated that gamma-ray bursts which too are isotropically distributed, may be responsible for EHECRs \cite{grb}. However since these are at cosmological distances, one would expect to see the GZK cutoff in the cosmic ray spectrum contrary to observations (cf. ref.\cite{grbtest}). Some of the above arguments may be evaded if the EHECR events are due not to nucleons but neutral particles such as photons and neutrinos. Although high energy photons also suffer energy losses in traversing the CMB and the extragalactic radio background, there is no threshold effect which would cause a cutoff near the GZK value \cite{photon}. However the observed shower profile of the highest energy Fly's Eye event \cite{flyseye} argues against the primary being a photon since it would have interacted on the geomagnetic field and started cascading well before entering the atmosphere \cite{notphoton}. The observed events are also unlikely to be initiated by neutrinos as they all have incident angles of less than $40^\circ$ from the zenith and thus too small a path length in the atmosphere for interactions \cite{gqrs96}. This argument may be evaded if neutrinos become strongly interacting at high energies due to new physics beyond the Standard Model \cite{neutrino,hongmo}, but such proposals are found not to be phenomenologically viable \cite{bhg98} (although this is disputed \cite{nu}). (Alternatively, the propagating high energy neutrinos could annihilate on the relic cosmic neutrino background, assumed to have a small mass of ${\cal O}(0.1)$~eV, to make hadronic jets within the GZK zone \cite{weiler}.) Other exotic possibilities have been suggested, e.g. monopoles \cite{wk96}, stable supersymmetric hadrons \cite{farrar} and loops of superconducting cosmic string (`vortons') \cite{bp97}. However these possibilities have many phenomenological problems \cite{mn98,ber} and we do not discuss them further. Thus one is encouraged to seek `top-down' explanations for EHECRs in which they originate from the decay of massive particles, rather than being accelerated up from low energies. The most discussed models in this connection are based on the annihilation or collapse of topological defects such as cosmic strings or monopoles formed in the early universe \cite{hill,td,bs95,bv97}. When topological defects are destroyed their energy is released as massive gauge and Higgs bosons which are expected to have masses of ${\cal O}(10^{16})\gev$ if such defects have formed at a GUT-symmetry breaking phase transition. The decays of such particles can generate cascades of high energy nucleons, $\gamma$-rays and neutrinos. A more recent suggestion is that EHECRs arise from the decays of metastable particles with masses $m_X\sim10^{13}-10^{16}\gev$ which constitute a fraction of the dark matter \cite{bkv97}. These authors suggest that such particles can be produced during reheating following inflation or through the decay of hybrid topological defects such as monopoles connected by strings, or walls bounded by strings. The required metastability of the particle is ensured by an unspecified discrete symmetry which is violated by quantum gravity (wormhole) effects. Another suggestion is that the long lifetime is due to non-perturbative instanton effects \cite{kr97}. In ref.\cite{fkn97}, a candidate metastable particle is identified in a $SU(15)$ GUT. A generic feature of these `top-down' models is that the EHECR spectrum resulting from the decay cascade is essentially determined by particle physics considerations. Of course the subsequent propagation effects have astrophysical uncertainties but since the decays must occur relatively locally in order to evade the GZK cutoff \cite{bkv97}, they are relatively unimportant. Thus although the proposal is speculative, it is possible, in principle, to make reliable calculations to confront with data. In this work we consider another possible candidate for a relic metastable massive particle \cite{ben98} whose decays can give rise to the observed highest energy cosmic rays. First we discuss (\S~\ref{crypton}) why this candidate, which arises from the hidden sector of supersymmetry breaking, is perhaps physically better motivated than the other suggested relics. We then undertake (\S~\ref{decay}) a detailed calculation of the decay cascade using a Monte Carlo event generator to simulate non-perturbative QCD effects. This allows us to obtain a more reliable estimate of the cosmic ray spectrum than has been possible in earlier work on both topological defect models \cite{td} and a decaying particle model \cite{bkv97}. We confront our results with observations and identify the mass and abundance/lifetime required to fit the data. We conclude (\S~\ref{concl}) with a summary of experimental tests of the decaying particle hypothesis. | We have investigated the hypothesis that the highest energy cosmic rays, in particular those observed beyond the GZK cutoff, arise from the decay of massive metastable relic particles which constitute a fraction of the dark matter in the galactic halo. To simplify computations (using the HERWIG Monte Carlo event generator) we have considered only decays into $q\bar{q}$ pairs with unit branching ratio. Comparison with experimental data indicates that a decaying particle mass of ${\cal O}(10^{12})\gev$ is required to fit the spectral shape while the absolute flux requires a lifetime of ${\cal O}(10^{20})\yr$ if such particles contribute the critical density. The predicted decay spectra may be somewhat altered if 3-body decays and other final states (e.g. supersymmetric particles \cite{bk97}) are considered. However our conclusions regarding the preferred mass and relic abundance/lifetime of the decaying particle are unlikely to be affected. In particular it would appear that the approximations used to calculate the particle spectra in previous studies of decaying topological defects \cite{td} and hypothetical massive particles \cite{bkv97} were not sufficiently accurate. Our work indicates that the topological defect model is disfavoured unless the mass of the decaying gauge bosons is less than about $10^{13}\gev$, which is well below the unification scale of $\sim10^{16}\gev$. (A similar conclusion is arrived at by independent arguments in refs.\cite{bss97,vah98}.) By contrast, cryptons from the hidden sector of supersymmetry breaking have a mass of the required order, as well as a decay lifetime which is naturally suppressed. However their relic abundance is difficult to estimate reliably, although we have argued that it may be cosmologically interesting. The primary intention of this work is to attempt to quantify the decaying particle hypothesis in a manner which is of interest to experimentalists. We have therefore computed the expected neutrino to proton ratio as a function of energy since this is an important test of competing hypotheses for forthcoming experiments, in particular the Pierre Auger project \cite{auger}. Of course our cleanest prediction is that the cosmic ray spectrum should cut off just below the mass of the decaying crypton, at $\sim3\times10^{20}\ev$. Moreover, with sufficient event statistics it should be possible to identify the small anisotropy which should result from the distribution of the decaying particles in the Galactic halo \cite{dt98}. Thus although the hypothesis investigated here is very speculative, it is nevertheless testable. Perhaps Nature has indeed been kind to us and provided a spectacular cosmic signature of physics well beyond the Standard Model. | 98 | 4 | hep-ph9804285_arXiv.txt |
9804 | astro-ph9804175_arXiv.txt | s{We are contructing an interferometric telescope, the Very Small Array, to study the cosmic microwave background on angular scales 0.2--$4.5^{\circ}$. The physical layout and electronic design of the telescope are optimised to give maximum protection from systematic effects, while still providing sufficient sensitivity to make high signal-to-noise images. A prototype single baseline is currently being tested, with scientific results expected during 2000.} | It is widely accepted that the majority of the cosmological results that might be obtained from the cosmic microwave background (CMB) will come from accurate measurements of the CMB power spectrum over the region of the accoustic peaks, that is in the range $100 < l < 2000$ (where $l$ is the spherical harmonic multipole). We have therefore been constructing an instrument, the Very Small Array (VSA), to measure the power spectrum (and provide images) in precisely this angular range. We have argued previously\cite{jones} that interferometers are very well suited for ground-based CMB measurements, given their relative immunity to the atmosphere and other systematic effects compared with switched-beam experiments. Here we will describe some of the more detailed design considerations of the array, and review the progress of the project so far. | 98 | 4 | astro-ph9804175_arXiv.txt |
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9804 | astro-ph9804205_arXiv.txt | We present HST WFPC2 observations in three bands (F450W=B, F467M and F814W=I) of a group of three galaxies at $z=2.8$ discovered in a ground-based narrow-band search for \lya emission near the $z=2.8$ quasar \pks. One of the galaxies is a damped \lya (DLA) absorber and these observations bear on the relation between the DLA clouds and the Lyman-break galaxies and the stage in the evolution of galaxies they represent. We describe a procedure for combining the undersampled WFPC2 images pointed on a sub-pixel grid, which largely recovers the full sampling of the WFPC2 point spread function (psf). These three galaxies have similar properties to the Lyman-break galaxies except that they have strong \lya emission. The three galaxies are detected in all three bands, with average $m_B\sim26$, $m_I\sim25$. Two of the galaxies are compact with intrinsic (i.e. after correcting for the effect of the psf) half-light radii of $\sim0.1$ arcsec ($0.4 h^{-1}$ kpc, $q_{\circ}=0.5$). The third galaxy comprises two similarly compact components separated by 0.3 arcsec. The HST images and a new ground-based \lya image of the field provide evidence that the three galaxies are more extended in the light of \lya than in the continuum. Combined with the evidence from the \lya line widths, previously measured, this suggests that we are measuring the size of the surface of last scattering of the escaping resonantly-scattered \lya photons. The measured impact parameters for this DLA galaxy (1.17 arcsec), for a second confirmed system, and for several candidates, provide a preliminary estimate of the cross-section-weighted mean radius of the DLA gas clouds at $z\sim 3$ of $<13 h^{-1}$ kpc, for $q_{\circ}=0.5$. The true value is likely substantially smaller than this limit as DLA clouds at small impact parameter are harder to detect. Given the observed sky covering factor of the absorbers this implies that for $q_{\circ}=0.5$ the space density of DLA clouds at these redshifts is more than five times the space density of spiral galaxies locally, with the actual ratio probably considerably greater. For $q_{\circ}=0.0$ there is no evidence as yet that DLA clouds are more common than spiral galaxies locally. We summarise evidence that filamentary structures occur in the distribution of galaxies at high redshift. | By studying typical galaxies at high redshift, here $z>2$, we can record the origins of normal galaxies such as our own Milky Way. Schmidt \cite{sc65} was the first to observe a high$-$redshift galaxy when he obtained a spectrum of 3C9, a radio$-$loud quasar of redshift $z=2.01$. Normal star$-$forming high$-$redshift galaxies are about 1000 times fainter, and remained undiscovered until the present decade. A small number of candidates (some of which may be active galaxies) have been identified with 4$-$metre telescopes, through the detection of Ly$\alpha$ emission (e.g. Steidel, Sargent, and Dickinson 1991, Lowenthal et al. 1991, M\o ller and Warren 1993, Pascarelle et al. 1996, Francis et al. 1995). Another approach has been to employ deep broad$-$band imaging to identify candidates by the expected Lyman$-$limit discontinuity in their spectra (Steidel and Hamilton 1992). The brightest Lyman$-$break galaxies have $m_R\sim 24$ and spectroscopic confirmation has only become feasible with the completion of the Keck 10$-$metre telescope. The report by Steidel et al. (1996) on the results of spectroscopic observations of Lyman$-$break candidates really marked the beginning of the statistical study of high-redshift galaxies. They were able to confirm redshifts for 15 star$-$forming galaxies in the range $3.0<z<3.5$. They discovered that Lyman$-$break galaxies generally show weak Ly$\alpha$ emission, but that the rest$-$frame ultra$-$violet spectra may be recognised by the presence of absorption lines characteristic of the spectra of nearby star$-$forming galaxies. Madau et al. (1996) applied the same methodology to the Hubble Deep Field data to extend the results to fainter magnitudes, and their analysis currently provides the best lower limits to the integrated star formation rate in galaxies at $z>2$. Complementary to the searches for starlight from high$-$redshift galaxies have been the analyses of absorption lines in the spectra of high$-$redshift quasars. These studies have provided measurements of the mass density of neutral hydrogen in the universe (e.g. Wolfe et al. 1986, Lanzetta et al. 1991), and the metallicity of the gas (e.g. Pettini et al 1994), and how these quantities have changed with redshift. The analysis of Pei and Fall (1995) of the absorption--line data reconstructs the global history of star formation, gas consumption, and chemical enrichment, accounting in a self--consistent way for the effects of the progressive extinction of the background quasars due to dust as star formation proceeds. The advantage of the absorption--line approach to the history of star formation is that it is global in nature as all the neutral gas at any redshift is directly observed. With the deep imaging studies it is necessary to extrapolate below the survey flux limit to measure the total star formation rate. The weakness of the global approach however is that it tells us nothing about the sizes and morphologies of the galaxies in which the star formation occurs. Deep imaging of DLA absorbers provides the missing information. In other words if we imaged the absorbers we could combine the data of faint galaxy surveys (luminosities, sizes, shapes) with the information on the gas obtained from the spectroscopic studies (column densities, chemistry), allowing a more detailed comparison with theories of galaxy formation. Unlike for MgII absorbers (e.g. Bergeron and Boiss\'e 1991, Steidel, Dickinson, and Persson 1994a) little progress has been made in such a programme of imaging of DLA systems. There are only two published unambiguous detections, discussed below, as well as a small number of candidate counterparts of DLA absorbers (at both low redshift $-$ Steidel et al. 1994b, Le Brun et al. 1997 $-$ and high-redshift $-$ Arag\'{o}n-Salamanca et al. 1996). In this paper we report on the results of 30 orbits of imaging observations with the {\it Hubble Space Telescope} (HST) of a group of three galaxies at $z=2.81$, named S1, S2, S3, one of which (S1) is a damped \lya absorber. The three galaxies lie in the field of the quasar \pks, and were detected in the light of \lya emission by M\o ller and Warren (1993, hereafter Paper I). A detailed discussion of their nature was presented by Warren and M\o ller (1996, hereafter Paper II \footnote{We note here that recent spectroscopic observations by Ge et al. (1997) confirm our conclusion that the \lya emission from the DLA absorber is due to star formation rather than photoionisation by the quasar.}). The HST images provide measures of their sizes, magnitudes, and colours. The other high$-$redshift damped system that has been successfully imaged is the absorber at $z=3.150$, towards the quasar $2231+131$, observed by Djorgovski et al. \cite{dj96}. In this paper we intercompare the measured properties of these two absorbers, the two companion galaxies S2 and S3, and the population of Lyman$-$break galaxies, to draw conclusions about the space density of DLA absorbers, their structure, and the relation between the DLA absorbers and the Lyman-break galaxies. The layout of the rest of the paper is as follows. The HST observations of the field towards \pks are described in Section 2, as well as new ground$-$based narrow$-$band observations of the field. The three galaxies are very small, and the HST images were dithered with a half-integer pixel step in order to improve the image sampling. The algorithm for combining the images is outlined in this section. In Section 3 we present the reduced HST images. In Section 4 we provide the results of aperture photometry and profile fitting of the HST and ground-based images. Finally, in Section 5 we discuss these results and their implications for our understanding of the nature of DLA absorbers and Lyman-break galaxies. | \subsection{Continuum emission} The three galaxies S1, S2, S3 are similar in their properties to members of the recently detected population of Lyman-break galaxies, except that they have strong \lya emission, with mean restframe ${\rm EW=60\AA}$. For example the B and I magnitudes for our three sources $m_B\sim 26$, $m_I\sim 25$, are within the range of magnitudes of the Lyman-break galaxies of redshift $z\sim3$ observed by Steidel et al. (1996) and Lowenthal et al. (1997). The average half-light radius of the four objects listed in Table 4 is 0.1 arcsec. This is smaller than the value quoted for the Lyman-break galaxies (Giavalisco et al. 1996, Lowenthal et al. 1997), by a factor of two to three. However we have broken S3 into two sub components and measured the radii for each subcomponent. In addition the radii quoted in Table 4 have been corrected for the effect of the psf. Young star-forming galaxies should have approximately constant flux per unit frequency over the restframe ultraviolet and optical regions of the spectrum. The expected colour of a source at $z=2.8$ with a power-law continuum varying as $f_{\nu}\propto\nu^0$, after accounting for typical absorption in the \lya forest, is $m_B-m_I=0.66$. As seen in Table 3 the $m_B-m_I$ colours for the three sources S1, S2, S3 are consistent with this value, for both the small and large apertures. The fact that S1 has properties similar to the Lyman-break galaxies suggests that DLA absorbers and Lyman-break galaxies may commonly be associated with each other. Lowenthal et al (1997) have also made a connection between the two populations by suggesting that the broad \lya absorption line seen in several cases in the spectra of Lyman-break galaxies is a damped line from neutral hydrogen in the galaxy. However the absorption lines seen in their average spectrum and in the relatively high S/N spectrum presented by Ebbels et al. (1996) are not optically thick in the line centre. The absorption--line profiles are difficult to interpret because the background source is extended. It is even possible that in some cases the absorption line is stellar absorption in the spectra of late B stars and that the galaxies are observed in a post-burst phase, rather than actively forming stars (see Valls-Gabaud 1993) i.e. for some cases this may be the explanation for the weakness of the \lya emission line in the Lyman-break galaxies, as opposed to dust. High-redshift galaxies with strong \lya emission, like the three sources near \pks, may therefore be younger than many of the Lyman-break galaxies. \subsection{Extended \lya emission} In this subsection we summarise the evidence that the \lya emission from the three sources is more extended than the continuum emission. The evidence for S2 and S3 comes from an examination of the $m_B-m_M$ and $m_B-m_N$ colours, summarised in Table 3. The bottom row in Table 3 provides the predicted $m_B-m_M$ colours for each source computed, as detailed in Section 4.1.2, from the B and N total magnitudes. The small aperture $m_B-m_M$ colours are smaller than the predicted values, as would be the case if the \lya emission is more extended than the continuum. For S2 the difference is significant at $2.8\sigma$, and for S3 at $2.3\sigma$. Notice also that the large aperture $m_B-m_M$ colours are larger than the small aperture colours, which is consistent with the hypothesis of extended \lya emission. A broad absorption line in the M passband could only explain part of the discrepancy between the measured and predicted colours. For example for S2, simply removing $50{\rm\AA}$ of the continuum flux within the M band changes the $m_B-m_M$ colour by only 0.1 mag. The total B magnitude for S1 (i.e. the large-aperture measurement from the HST image, Table 4) is not reliable because of the uncertainty associated with the psf subtraction. This prevents us from usefully making the same comparison between predicted and measured $m_B-m_M$ colours for S1. However there is an indication that S1 may also be more extended in \lya. The measurements of the half-light radius of S1 from the ground-based \lya image yielded values of $0.51^{+0.08}_{-0.08}$ arcsec for the exponential profile and $0.66^{+0.25}_{-0.19}$ arcsec for the de Vaucouleurs profile (Section 4.2.2). These values are significantly larger than the half light radius of the region of continuum emission of $0.13\pm0.06$ arcsec (Table 4), measured from the HST image. Such a diffuse contribution of \lya to the M-band flux would be hard to detect, and could have been partially subtracted in attempting to remove large-scale residuals from the psf subtraction (Section 4.2.1). However we cannot rule out the possibility that there is also a similar diffuse contribution to the continuum flux. The argument for extended \lya is made visually in Figure 4. Here we compare the M-band images after subtraction of the computed contribution of the continuum (i.e. showing the \lya emission only), against how the objects would have appeared if the \lya light profile were the same as the continuum profile. \begin{figure} \vspace{8.5cm} \special{psfile=Reffig_1.ps angle=0 voffset=-52 hoffset=-14 vscale=40 hscale=40} \caption[ ]{Illustration of the evidence for extended \lya emission for, from top to bottom, S1, S2, S3. Each panel shows the same 2 arcsec by 2 arcsec region as Figure 3, here smoothed by a 0.15 by 0.15 arcsec boxcar filter to enhance faint features. The left column shows the summed image from all three filters. The middle column shows the M-band images with the computed contribution of the continuum subtracted. This image, then, contains only the \lya line component. In the right column we show how the \lya-only image would have appeared if the \lya profile were identical to the continuum profile.} \end{figure} To summarise, there is evidence that the regions of \lya emission from these three sources are more extended than the regions of continuum emission. Although the colour and size differences quoted above are only marginally significant, these results accord with our earlier conclusion (Paper II) that the cause of the relatively broad \lya emission lines for S1 and S3 is resonant scattering, as the photons escape though a high column density of HI. In this picture the \lya image records the last scattering surface, which will be larger than the region of star formation. We note in passing that this implies that the observed strength of \lya absorption and emission in spectra can depend on the slit width used. \subsection{The sizes of the DLA gas clouds} In this section we make a summary of the measured impact parameters $b$ of DLA absorbers. Using this we estimate the cross-section-weighted mean radius $\bar{R}_{DLA}$ of the gas clouds at $z>2$. Combining this information with the measured line density of DLA absorbers $dn/dz$ we are able to infer the ratio of the comoving space density of DLA absorbers at $z>2$ to the local space density of spiral galaxies. \begin{table*} \begin{flushleft} \caption[]{Measured impact parameters $b$ of DLA absorbers} \begin{tabular}{lcccccc} \hline\noalign{\smallskip} Quasar&$z_{abs}$&log(N(HI))&$b$&$b$& confirmed & reference \\ & & &($q_{\circ}=0.0$)&($q_{\circ}=0.5$)& & \\ & &cm$^{-2}$ &kpc&kpc& & \\ \hline $0454+0393$ & 0.86 & 20.8 & $4.1h^{-1}$& $3.3h^{-1}$ & N & 1 \\ $0302-223$ & 1.01 &$\leq20.0\:\:\:\:\:$& $6.2h^{-1}$& $4.9h^{-1}$ & N & 1 \\ $1331+170$ & 1.78 & 21.2 & $4.7h^{-1}$& $3.1h^{-1}$ & N & 1 \\ $0151+048$A & 1.93 & 20.4 & $7.7h^{-1}$& $5.0h^{-1}$ & Y & 2 \\ $1215+333$ & 2.00 & 21.0 & $8.4h^{-1}$& $5.3h^{-1}$ & N & 3 \\ $0841+129^a$& 2.37 & 21.3 & $8.0h^{-1}$& $4.7h^{-1}$ & N & 3,4 \\ $0841+129^a$& 2.48 & 21.0 & $8.0h^{-1}$& $4.7h^{-1}$ & N & 3,4 \\ $0528-250^b$& 2.81 & 21.3 & $8.1h^{-1}$& $4.5h^{-1}$ & Y & 5 \\ $2231+131$ & 3.15 & 20.0&$15.7h^{-1}\:\,$& $8.2h^{-1}$ & Y & 5 \\ \noalign{\smallskip} \hline \end{tabular} \\ $^a$ the detected galaxy could be the counterpart to one or other of the two DLA absorbers listed \\ $^b$ impact parameter taken from HST image (this paper) \\ References: 1. Le Brun et al. (1997), 2. Fynbo et al. (1997), 3. Arag\'{o}n-Salamanca et al. (1996), 4. Wolfe et al. (1995), 5. M\o ller and Warren (1993), 6. Djorgovski et al. (1996). \end{flushleft} \end{table*} There have been many searches for optical counterparts of the DLA absorbers. In Table 6 we have listed the small number of DLA absorbers, with redshifts $z>0.8$, for which likely candidate or confirmed optical counterparts have been detected. Listed are the quasar name, the absorber redshift and column density, the measured impact parameter (i.e. the projected physical separation between the optical counterpart and the line of sight to the quasar), and whether or not the counterpart has been confirmed, i.e. the redshift of the counterpart has been measured to be the same as the redshift of the absorber. In Figure 6 we plot impact parameter against column density for the absorbers listed in Table 6, for $q_{\circ}=0.5$. The detection of the counterpart to a DLA absorber by broad-band imaging, which requires the digital subtraction of the quasar image, is more difficult for small impact parameters where the photon noise from the quasar image is greater. Therefore it is very likely that the measured impact parameters of the few absorbers for which counterparts have been detected are larger than the average impact parameter for DLA absorbers. The vertical line in Figure 6 at ${\rm log(N(HI))=20.3}$ marks the conventional lower limit of the column density of DLA absorbers. The distribution of points in the figure is consistent with the expected anticorrelation between impact parameter and column density. We are interested in the mean impact parameter $\bar{b}_{DLA}$ that would be measured if counterparts were detected for all DLA absorbers in a large unbiased survey. From inspection of Figure 6 we suggest that it is safe to conclude that for $z>2$, $q_{\circ}=0.5$, $\bar{b}_{DLA}$ is less than $7h^{-1}$ kpc. For $q_{\circ}=0.0$ we suggest $\bar{b}_{DLA}<13h^{-1}$ kpc. The actual mean impact parameters are probably substantially smaller than these limits. We relate the mean impact parameter to the cross-section-weighted mean radius $\bar{R}_{DLA}$ of the absorbers by $\bar{b}_{DLA}=\alpha\bar{R}_{DLA}$. The edge of the DLA is the point at which the column density falls below $N(HI)=2\times10^{20}$ cm$^{-2}$, and we effectively assume that the column density falls off sharply at larger radii. For face on disks $\alpha=\frac{2}{3}$, and for a disk that is nearly edge on $\alpha=\frac{4}{3\pi}$, so we take $\alpha=0.55$ as an average value for randomly inclined disks. On this basis we infer the following limits on the cross-section-weighted mean radius of DLA absorbers at high redshift: $$\bar{R}_{DLA}<23.6h^{-1} {\rm kpc}\:\:\:\: (z>2, q_{\circ}=0.0) $$ $$\bar{R}_{DLA}<12.7h^{-1} {\rm kpc}\:\:\:\: (z>2, q_{\circ}=0.5). $$ Using these limits and the measured line density of absorbers $dn/dz$ we can compute limits to the ratio of the comoving space density of DLA absorbers to the the local space density of spiral galaxies. We follow essentially the methodology employed by Wolfe et al. (1986), and Lanzetta et al. (1991, hereafter LTW). For spiral galaxies locally they adopted a galaxy luminosity function $\Phi(\frac{L}{L_*})$ of Schechter form, with power-law index $s$, a power-law (Holmberg) relation between radius and luminosity, of index $t$, and a ratio between gas radius and optical (Holmberg) radius $\xi$, independent of luminosity. They found that the incidence of DLA absorbers per unit redshift $dn/dz$ at $z\sim 2.5$ was considerably higher than expected, by a factor $F\sim 5$, on the basis of no evolution in galaxy cross section or luminosity function normalisation. We now allow for evolution by supposing that the space density of DLA absorbers is higher than the local space density of spirals by a factor $E_{\Phi}(z)$, and that the gas radii of galaxies are larger at high redshift by a factor $E_r(z)$. Since $dn/dz$ is proportional to the product of the space density and the galaxy cross section $\sigma$, we have that $F=E_{\Phi}(z)E^2_r(z)$. By comparing the expected value of the cross-section-weighted radius $\bar{R}$ to the measured limits to $\bar{R}_{DLA}$, we obtain limits to $E_r(z)$. Then from the measured values of $F$ we determine limits to $E_{\Phi}(z)$, which is our goal. Under the above assumptions the cross-section-weighted average radius of DLA absorbers is given by: $$\bar{R}(z)=\frac{\int_0^{\infty} R(\frac{L}{L_*})\sigma(\frac{L}{L_*}) \Phi(\frac{L}{L_*})d(\frac{L}{L_*})} {\int_0^{\infty}\sigma(\frac{L}{L_*})\Phi(\frac{L}{L_*})d(\frac{L}{L_*})}$$ $$=\frac{E_r(z)\xi R_*\Gamma(1+3t-s)}{\Gamma(1+2t-s)}$$ where $R_*$ is the optical radius of a local $L_*$ spiral galaxy. Following LTW we adopt the following values of the parameters: $t=0.4$, $s=1.25$, $\xi=1.5$, $R_*=11.5h^{-1}$ kpc. This leads to $\bar{R}(z)=11.0h^{-1}E_r(z)$. Comparing with the above measured limits to $\bar{R}_{DLA}$ we obtain the following limits to the growth factor of galaxy disks: $$E_r(z)<2.15\:\: (z>2, q_{\circ}=0.0)$$ $$E_r(z)<1.16\:\: (z>2, q_{\circ}=0.5)$$ For their sample D2, of which at least 30 out of the 38 candidate DLA absorbers have been confirmed, LTW found $F=3.8$, for $q_{\circ}=0.0$, and $F=7.1$, for $q_{\circ}=0.5$. These results then imply that the ratio of the comoving space density of DLA absorbers at high redshift to the local space density of spiral galaxies is given by: $$E_{\Phi}(z)=\Phi_*(DLA)/\Phi_*(spiral)>0.8\:\: (z>2, q_{\circ}=0.0) $$ $$E_{\Phi}(z)=\Phi_*(DLA)/\Phi_*(spiral)>5\:\: (z>2, q_{\circ}=0.5) $$ Because the actual average impact parameters of DLA absorbers are very likely substantially smaller than the quoted limits, the actual comoving space density ratios are probably considerably greater than the limits quoted above. \begin{figure} \vspace{11.5cm} \special{psfile=Fig_bvsN.ps angle=0 voffset=-60 hoffset=-35 vscale=50 hscale=50} \caption[ ]{Relation between impact parameter and column density for all confirmed or candidate counterparts of DLA absorbers of redshift $z>0.8$ (listed in Table 6). Different symbols correspond to different redshift ranges: squares $z>2.5$; triangles $2.5\geq z>1.5$; circles $1.5\geq z>0.8$. Absorbers for which the counterpart has been confirmed by spectroscopy are shown as filled symbols, and are labeled. The two possible DLA absorbers for the quasar $0841+129$ are joined by a dotted line. The hatched area shows the densely populated region from the simulation of Katz et al. (1996), and the dashed lines are the boundaries enclosing all the points in the simulation.} \end{figure} In Figure 5 we show also the relation between impact parameter and column density measured by Katz et al (1996) from a hydrodynamic simulation of a CDM universe with $q_{\circ}=0.5$. Given the limited spatial resolution of the simulation, and the fact that star formation and consequent feedback were not treated, it would be premature to draw any conclusions about the apparent good agreement between the results of the simulation and the observations. Nevertheless this plot and the above calculation underscore the importance of measuring impact parameters for a large sample of DLA absorbers, a goal we are pursuing by imaging with the STIS instrument on HST. \subsection{The structure of the DLA absorbers} The observations summarised above provide some indication of the typical structure of a DLA absorber. The following sketch is suggested. At the centre of the gas cloud, corresponding in projection to the highest column densities, there may be a region of star formation, of diameter $\sim 1h^{-1}$ kpc. This central source would be observed as a Lyman-break galaxy. The gas column density decreases outwards, as suggested by the anticorrelation between impact parameter and $N_{HI}$. The size of the region over which the column density is $>2\times10^{20}$cm$^{-2}$ is several kpc. The region of star formation would be surrounded by a zone of ionised gas. The \lya photons are resonantly scattered in escaping to the surface of the surrounding cloud of neutral gas. The size of the observed region of \lya emission is larger than the region of star formation, but smaller than the diameter of the gas cloud. The latter indicates a preferred direction of escape, implying that the gas resides in a flattened structure. In reality the gas cloud is likely to be irregular rather than smooth, and might contain several knots of star formation, surrounded by HII regions of different sizes, and the \lya photons will escape preferentially where the column density of neutral gas is lowest, possibly along a complex network of tunnels. Merging clouds will be observed as galaxies with irregular structure. \subsection{Filamentary structure in the distribution of galaxies at $z=3$} We have previously suggested (Paper II) that the approximately linear arrangement of S1, S2, S3, as well as other known groups at high-redshift, may correspond to the filamentary arrangements of galaxies and galaxy sub-units that are found in computer simulations of the high-redshift universe. Figure 6 summarises the observational situation, and shows the spatial arrangement of four high-redshift groups of galaxies discovered in \lya searches (from top to bottom: Francis et al., 1995, considering the \lya absorber as a galaxy; Pascarelle et al., 1996; Paper I; Le F\`{e}vre et al., 1996). All groups have been rotated to a horizontal baseline. Clearly each of the structures is elongated. \begin{figure} \vspace{9.5cm} \special{psfile=Fig_align.ps angle=0 voffset=-145 hoffset=-78 vscale=58 hscale=58} \caption[ ]{Relative positions of galaxies in four high-redshift groups discovered in \lya searches. Each group has been rotated so that the principal axis is horizontal in the figure, and the groups have been offset from one another vertically. Note the different scale of the lower plot.} \end{figure} To quantify the degree of alignment of each group we have measured the smallest internal angle $\delta_L$ of the triangle defined by the three objects in each group (for the PKS 0528-250 field at z=2.81 we used the position of S1 for the DLA absorber, rather than the position of the quasar sightline through the absorber). The measured angles are, from top to bottom in the figure, $\delta_L = 7.0^{\circ}, 1.0^{\circ}, 5.8^{\circ}, 2.1^{\circ}$. The average value of the angle $\delta_L = 4.0^{\circ}$ is certainly very suggestive of filamentariness in the distribution of galaxies at high redshift. It is also notable that the sizes of the groups range over an order of magnitude. In the computer simulations filaments are seen at all scales (e.g. Evrard et al. 1994). We do not suggest that every high-redshift group of three or four galaxies discovered will exhibit a similar degree of alignment, but rather that filamentary networks akin to those seen in the simulations may become visible as larger samples of high-redshift galaxies with denser sampling become available. The phenomenon might provide a useful discriminant for models of structure formation. | 98 | 4 | astro-ph9804205_arXiv.txt |
9804 | astro-ph9804033_arXiv.txt | Recent observations supported by theoretical models have lead to the view that giant and supergiant stars are over abundant, and/or a high metallicity component may be present, in the stellar populations at the centres of active galaxies. Here we attempt to quantify these effects by observing the strengths of the stellar absorption lines of Mg~b, NaI, CaII triplet as well as molecular bands such as CN and TiO. Using long-slit spectroscopic data we are able to separate the stellar populations in and around the nucleus, for a sample including, normal, LINER, starburst and Seyfert galaxies. In this paper we present the data, namely spectra of the nucleus and of a number of circum-nuclear regions. Comparisons reveal gradients in both the reddening and the stellar population within the central regions of most galaxies. Detailed stellar population synthesis is presented in a companion paper. | A crucial unsolved question is whether the stellar populations in the nuclear regions of Active Galaxies differ from those of non-active galaxies of the same Hubble type. Correlations between near IR CO indices, far-infrared and X-ray luminosities of active galaxies may indicate that the more powerful monsters reside in more actively star-forming host galaxies (Yamada, 1994). The accelerated star formation caused by dynamic instabilities which trigger and/or fuel the nuclear activity could result in an overabundance of giant, supergiant and super metal rich (SMR) stars (Scoville, 1992 and references therein). Terlevich et al. (1990) observed that in some Active Galactic Nuclei (AGN) the near IR CaII triplet absorption features are as strong, or even stronger than those of normal non-active galactic nuclei. They have suggested that the ``featureless" blue continuum previously thought to be non-stellar in origin, actually arises from the unresolved continuum from a young cluster of stars containing red supergiants. The presence of a ``truly'' featureless continuum is then required only in the case of some Seyfert~1 galaxies which show dilution of the CaII triplet lines. Based on independent evidence we know that in some galaxies a starburst region surrounds the unresolved nucleus which emits the broad lines e.g. NGC~7469 (Wilson et al., 1991). Nonetheless this does not prove a causal connection between the starburst and AGN. An alternative interpretation, plausible within the framework of our knowledge of AGN, is that the stellar lines are superimposed on a non-stellar nuclear continuum, and are formed in a super metal rich population. Indeed, one could expect abundance anomalies due to the intense star formation in the metal rich environment of the nucleus. From our detailed study of the Seyfert~1 nucleus in the galaxy NGC~3516 (Serote Roos et al., 1996), we find that the stellar population exhibits a noticeable dilution by a featureless continuum in the wavelength range 5000-9800\AA\ {\it as well as} a high proportion of super metal rich stars. Thus, before dismissing the possibility of the presence of any non-thermal component in the near infrared continuum of Seyfert galaxies, we must first define the stellar population in the nucleus and the surrounding regions for a sample of AGN of all levels of activity. Until now most studies tackling this subject have made use of only a small spectral domain around a few spectral features. In order to make further progress it is necessary to extend these studies to cover more stellar absorption features, providing signatures of different stellar populations e.g. MgI $\lambda$5175, NaI~D $\lambda$5896 and 8196\AA\ and the numerous TiO and CN bands. The NaI lines, for example, help to distinguish between the effects of overmetallicity rather than a supergiant dominated population. Indeed, the ambiguity between metallicity and supergiant excess can be resolved by comparison of the CaII triplet and NaI absorption strengths (e.g. Zhou, 1991) although the interstellar contribution to NaI introduces some uncertainty, and also by comparison with the strength of MgI (Couture \& Hardy, 1990). In this paper we present long-slit spectroscopy, in the range 5000-9800\AA\, of a sample of galaxies with different levels of activity. These observations are used to estimate radial gradients in the stellar population and to extract the stellar spectrum of the very nucleus. The sample is selected to include various classes of active galaxies, i.e. Seyferts 1 and 2, LINERs and starbursts, in order that the strength of the stellar features may be correlated with the general properties of each class. We also present data for two normal galaxies which will be used as comparison templates. This is not a complete sample in any statistical sense, but it does provide insights into the diversity of bulge stellar populations in and around AGN of different classes of activity. The detailed analysis of the stellar populations for this sample using a new spectral synthesis programme (Pelat, 1997), is to be found in Paper II (Serote Roos et al., 1997). | The nuclei of our sample galaxies are generally redder than the outer regions (this excludes Seyfert 1s as their nuclear spectra are dominated by strong broad emission lines). This colour gradient could either be due to dust or to stellar population gradients. In addition to dust/population gradients, the presence of a featureless continuum is inferred in a number of AGN. This component might be of non-thermal origin, plausible in Seyfert galaxies, or of stellar origin in the case of nuclear starburst activity. The extra component dilutes the strengths of the stellar absorption lines in the spectra of the nuclear regions of AGN. A detailed study of the spectral shape of the diluting component will allow us to differentiate between the two hypotheses. A full population synthesis analysis for most of the galaxies in our sample, using a new method developed to determine unique solutions, is presented in Paper II. | 98 | 4 | astro-ph9804033_arXiv.txt |
9804 | astro-ph9804080_arXiv.txt | Ascertaining the core collapse supernova mechanism is a long-standing problem in astrophysics. The current paradigm begins with the collapse of a massive star's iron core and the generation of an outwardly propagating shock wave that results from core rebound. Because of nuclear dissociation and neutrino losses, the shock stagnates. This sets the stage for a shock reheating mechanism whereby neutrino energy deposition via electron neutrino and antineutrino absorption on nucleons behind the shock reenergizes it (Bethe \& Wilson 1985; Wilson 1985). The shock reheating phase is essential to the supernova's success, but it is precisely this phase that is difficult to simulate realistically. During shock reheating, core electron neutrinos and antineutrinos are radiated from their respective neutrinospheres, and a small fraction of this radiated energy is absorbed in the exterior shocked mantle. The shock reheating depends sensitively on the electron neutrino and antineutrino luminosities, spectra (best characterized by the {\small RMS} energies), and angular distributions in the region behind the shock (e.g., see Burrows \& Goshy 1993, Janka \& M\"{u}ller 1996, Mezzacappa \etal 1998). These, in turn, depend on the neutrino transport in the semitransparent region encompassing the neutrinospheres, necessitating a neutrino transport treatment that is able to transit accurately and seamlessly between neutrino-thick and neutrino-thin regions. Various neutrino transport approximations have been implemented in simulating core collapse supernovae. The most sophisticated approximation, which naturally has been used in realistic one-dimensional simulations, is multigroup flux-limited diffusion ({\small MGFLD}; e.g., Bowers \& Wilson 1982, Bruenn 1985, Myra \etal 1987). {\small MGFLD} closes the neutrino radiation hydrodynamics hierarchy of equations at the level of the first moment (the neutrino flux) by imposing a relationship between the flux and the gradient of the neutrino energy density (the zeroth moment). For example, \begin{equation} F_{\nu}=-\frac{c\Lambda}{3}\frac{\partial U_{\nu}}{\partial r}+..., \label{eq:mgfld} \end{equation} \begin{equation} \Lambda = \frac{1}{1/\lambda + |\partial U_{\nu}/\partial r|/3U_{\nu}}, \label{eq:lambda} \end{equation} \smallskip \noindent where $\lambda$ is the neutrino mean free path, and $U_{\nu}$ and $F_{\nu}$ are the neutrino energy density and flux (Bruenn 1985). [Other forms for the flux-limiter $\Lambda$ can be found in Bowers \& Wilson (1982), Levermore \& Pomraning (1981), and Myra \etal (1987).] Whereas the limits $\lambda \rightarrow 0$ and $\lambda \rightarrow \infty$ produce the correct diffusion and free streaming fluxes, it is in the critical intermediate regime where the {\small MGFLD} approximation is of unknown accuracy. Unfortunately, the quantities central to the postshock neutrino heating, i.e., the neutrino {\small RMS} energies, luminosities, and mean inverse flux factors, are determined in this regime, and given the sensitivity of the neutrino heating to these quantities, it becomes necessary to consider more accurate transport schemes. Moreover, in detailed one-dimensional simulations that have implemented elaborate {\small MGFLD} neutrino transport (e.g., see Bruenn 1993, Wilson \& Mayle 1993, and Swesty \& Lattimer 1994), explosions were not obtained unless the neutrino heating was boosted by additional phenomena, such as convection. This leaves us with at least two possibilities to consider: (1) Failures to produce explosions in the absence of additional phenomena, such as convection, have resulted from neutrino transport approximation. (2) Additional phenomena may be essential in obtaining explosions. | Comparing three-flavor {\small MGBT} and three-flavor {\small MGFLD} in postbounce supernova environments, we find that {\small MGBT} leads to a significant increase/decrease in the {\it net} heating/cooling rate, particularly just above/below the gain radius. The {\small MGBT} net heating rate can be as much as 2 times the {\small MGFLD} net heating rate above the gain radius, with net cooling rates that are typically 0.8 times the {\small MGFLD} rate below. These differences stem primarily from differences in the neutrino luminosities and mean inverse flux factors; the heating rate is linearly proportional to both these quantities, and differences in both add to produce a significant difference in the net heating rate. We also observe that the differences in the net heating rate are greatest at earlier postbounce times for a given progenitor mass, and at a given postbounce time, greater for greater progenitor mass. This is illustrated in Table 1. The enhancement in heating with increased progenitor mass suggests that the net heating enhancement from {\small MGBT} is potentially robust and self-regulated. In closing, our results are promising, and their ramifications for core collapse supernovae, and in particular, for the postbounce neutrino-heating, shock-revival mechanism, await one- and two-dimensional dynamical simulations with {\small MGBT} coupled to the core hydrodynamics. One-dimensional simulations are currently underway, and we plan to report on them soon. \begin{table}[t] \vspace{-24pt} \begin{center} \footnotesize\rm \caption{Maximum Net Heating/Cooling Rates \label{t1}} \begin{tabular}{lccc} \topline {Progenitor Mass [${\rm M}_{\odot}$]} & {$t_{pb}$ [ms]} & {Maximum Net Heating Ratio} & {Maximum Net Cooling Ratio} \\ \midline 15& 106& 2.0& 0.8 \\ & 233& 1.3& 0.8 \\ 25 & 156& 2.0& 0.8 \\ \bottomline \end{tabular} \end{center} \vspace{-12pt} \end{table} | 98 | 4 | astro-ph9804080_arXiv.txt |
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9804 | astro-ph9804049_arXiv.txt | We investigate in detail the hypothesis that low surface brightness galaxies (LSB) differ from ordinary galaxies simply because they form in halos with large spin parameters. We compute star formation rates using the Schmidt law, assuming the same gas infall dependence on surface density as used in models of the Milky Way. We build stellar population models, predicting colours, spectra, and chemical abundances. We compare our predictions with observed values of metallicity and colours for LSB galaxies and find excellent agreement with {\it all} observables. In particular, integrated colours, colour gradients, surface brightness and metallicity match very well to the observed values of LSBs for models with ages larger than 7 Gyr and high values ($\lambda > 0.05$) for the spin parameter of the halos. We also compute the global star formation rate (SFR) in the Universe due to LSBs and show that it has a flatter evolution with redshift than the corresponding SFR for normal discs. We furthermore compare the evolution in redshift of $[Zn/H]$ for our models to those observed in Damped Lyman $\alpha$ systems by \scite{Pettini+97} and show that Damped Lyman $\alpha$ systems abundances are consistent with the predicted abundances at different radii for LSBs. Finally, we show how the required late redshift of collapse of the halo may constrain the power spectrum of fluctuations. | Late type low surface brightness galaxies are found to be a significant fraction of the total galaxy population in the Virgo cluster (\pcite{Impey+88}), in the Fornax cluster (\pcite{Irwin+90}), and in the field (\pcite{Mcgaugh+95}; \pcite{Sprayberry+96}; \pcite{Sprayberry+97}). A theory of galaxy formation should therefore account for the existence of LSBs. On the other hand, LSBs are particularly useful because they are simpler objects than high surface brightness (HSB) galaxies: they have relatively little present day star formation and little dust (reddening). Therefore it is easier to model their stellar population and star formation history. They are also more dark matter dominated than HSBs, and therefore particularly suitable for probing the structure of dark matter halos. In order to compare theories of galaxy formation with observational surveys of galaxies it is important not only to quantify observationally the abundance of LSBs, but also to understand their nature. For instance, if LSBs are made of very young stellar populations, as often claimed in the literature, they may be irrelevant in the picture of the Universe at high redshift, while they could be a substantial component of that picture if their stellar populations are old. The existence of LSB galaxies can be readily understood if it is assumed, following \scite{Fall+80}, that the specific angular momentum of the baryons is approximately conserved during their dissipation into a rotationally supported disk, and that the disk length--scale is therefore related to the angular momentum of the dark matter halo. The low surface brightness is the consequence of the low surface density of the disk, which is due to the larger spin parameter of the dark matter halos of LSBs relative to the spin parameter of the halos of HSB disks (e.g. \pcite{Dalcanton_disc+97}; \pcite{Mo+97}; \pcite{Jimenez+Heavens+97}). However, while the surface density of a disk can be simply related to the mass and spin parameter of its dark matter halo, its surface brightness, and therefore its mass to light ratio ($M/L$), depends on the stellar population and could in principle vary with galaxy type and age. For instance, $M/L$ must have a significant time dependence due to the history of star formation in the disk and to the continuous death of massive stars. Moreover, $M/L$ in a given photometric band changes with time also as a consequence of the colour evolution of the stellar population, which is sensitive also to the history of the chemical enrichment of the disk. Therefore, a comparison between a simple model of disk formation and the observations cannot be made without an appropriate model for the stellar population in the disk that includes {\it self-consistently} the chemical evolution of the population. \begin{figure} \centering \leavevmode \epsfxsize=1. \columnwidth \epsfbox{surf_1p0.eps} \caption[]{Initial surface density for an isothermal sphere and 3 different masses. The solid lines correspond to $\lambda=0.1$ to 0.01 (from top to bottom) in steps of 0.01. The circular velocity of the halo (dashed line) is also plotted.} \label{f1} \end{figure} In this paper we model self-consistently the chemical and photometric evolution of the stellar population of LSB disks, formed in dark matter halos described by isothermal spheres. The aim of this paper is twofold: \begin{enumerate} \item show that only one parameter, namely the {\it spin parameter $\lambda$} of the dark halo, can explain the surface brightness of LSBs. \item use LSB disk galaxies to study the formation and evolution of all disk galaxies. \end{enumerate} The second point is motivated by the first, that is by the fact that LSB disks are in fact so similar to HSB disks. On the other hand, LSBs are more dark matter dominated than HSBs, and therefore better described by a very simple model where only the gravitational field of an isothermal halo is considered. Our assumptions are as follows: \begin{enumerate} \item{The specific angular momentum of baryonic matter is the same as for dark matter} \item{Gas settles until centrifugally supported, in a given dark matter halo} \item{The star formation rate is given by the Schmidt law} \item{The gas infall rate is assumed to be the same function of total surface density as used in models of the Milky Way} \end{enumerate} \begin{figure*} \centering \leavevmode \epsfxsize=1.6 \columnwidth \epsfbox{surf_o.eps} \caption[]{Contour plot of the value of the central surface density (in M$_{\odot}$pc$^{-2}$) for different values of the halo mass and spin parameter. Left panel: halo formation redshift $z=0.7$. Right panel: halo formation redshift $z=2.1$.} \label{f2} \end{figure*} We first use the density profile of the halo to compute the surface density of the settling disk. We then use the Schmidt formation law and an infall rate that reproduces the observed properties of the Galaxy in conjunction with the chemical evolution models by \scite{Matteucci+89} to compute the star formation rate at several radii of the disk and the evolution of several chemical species (H, D, He, C, N, O, Ne, Mg, Si, S, Ca, Fe and Zn). We then proceed to compute spectra, integrated colours, colour profiles and surface brightness for two different values of the spin parameter of the halo. We first present in the next section the simple non--self--gravitational disk model and the non--singular halo model. We describe in section 3 the set of synthetic stellar population models used to predict the spectra and colours of the stellar population. In section 4 we model the star formation in the disk with a Schmidt law and a gas infall law, assuming the same model of the chemical evolution of the Galaxy as \scite{Matteucci+89}. Finally, in section 5 we compute synthetic spectra and photometric properties of the stellar population in the disk, using the stellar evolution tracks of JMSTAR15, stellar atmosphere models by \scite{Kurucz_92} and J{\o}rgensen (private communication) and the chemical evolution models built for LSBs (see section 4). We discuss the results in section 6. The most important results of this work are: \begin{enumerate} \item Observed colour profiles, chemical abundances and surface brightness profiles for LSBs are well fitted if they are assumed to have a spin parameter for its halo higher than HSBs. \item LSBs are not young objects, as often claimed in the literature, since their colours are well fitted by old ($> 7$ Gyr) stellar populations. \item There is discrepancy between the photometric age of the galaxies and the age of formation of their halo, which indicates that the star formation can start about 2~Gyr before the halo is formed. This discrepancy is reduced to 1~Gyr if the Hubble constant is assumed to be $H_0=65$~kms$^{-1}$Mpc$^{-1}$ and completely removed if the Universe is open ($\Omega < 0.3$) or has a significant vacuum energy contribution ($\Lambda > 0.6$). \item The earliest stellar populations of LSBs could be present in the high redshift universe in a similar proportion to HSBs as they are now. Their colours at high redshift (as the colours of all disks) are about 1~mag bluer than at low redshift. \end{enumerate} In a previous paper (\pcite{Padoan+Jimenez+Antonuccio97a}) we have shown that LSBs are not necessarily young and un-evolved systems (e.g. \pcite{McGaugh+Bothun94,deBlok_phot+95}). We used the simplest possible model, a burst of star formation, to determine a lower limit to the age of LSBs studied by \scite{deBlok_phot+95}. In the present work we considerably improve our previous model because we use a more realistic continuous star formation process, we include self--consistently the chemical evolution, and we connect the disk model to the cosmological scenario using the spherical collapse model (\pcite{Gunn_Gott_72}). This more detailed model confirms our previous result that the blue LSB disks in the sample by \scite{deBlok_phot+95} are not un-evolved objects collapsed at late times from low initial over-densities (\pcite{McGaugh+Bothun94,Mo+94}), but rather normally evolved disk galaxies. Throughout the paper the present day value of the Hubble constant is assumed to be $H_0=75$~kms$^{-1}$Mpc$^{-1}$. | \subsection{The Nature of LSBs} \begin{figure*} \centering \leavevmode \epsfxsize=1.4 \columnwidth \epsfbox{lsbabund.eps} \caption[]{The redshift evolution for some elements is plotted. The top-left panel shows the evolution of $[O/H]$ and is compared with \scite{McGaugh_oxigen94} measures. Our model reproduces perfectly the peak in $[O/H]$ among LSBs found by \scite{McGaugh_oxigen94} at $[O/H]=-0.9$. The top-right panel shows the analogous to the previous one but for $[Zn/H]$ compared with \scite{Pettini+97} measures. It transpires from the figure that the predicted $[Zn/H]$ for LSBs is consistent with the spread found for DLAs.} \label{f17} \end{figure*} In this work we have shown that the photometric properties of the bluest galaxies in the sample by \scite{deBlok_phot+95} can be reproduced assuming that they form and evolve as normal disk galaxies, with relatively high spin parameter. That the spin parameter alone could explain the low surface brightness was already shown by \scite{Dalcanton_disc+97}, but they assumed a given $M/L$ ratio. Here we have strengthened the point by also explaining the colours, the colour gradients, and the chemical abundances. We have also provided synthetic spectra that could be compared with the observations, in the effort to prove that LSB disk galaxies are indeed normal galaxies with high spin parameter. The stellar populations of LSBs are rather evolved, especially in the central part of their disks. Even the bluest galaxies in the LSB sample by \scite{deBlok_phot+95} are older than 7~Gyr, and typically star formation in these galaxies starts about 9~Gyr ago. This confirms the results of \scite{Padoan+Jimenez+Antonuccio97a}, where the galaxies in the same sample were found to be at least 7~Gyr old. LSB and HSB disks are therefore hosted in similar dark matter halos that differ only for their spin parameter, and they evolve in a similar way, apart from the fact that LSB disks are less concentrated than HSB disks. This means that the present day (or low redshift) abundance of LSBs relative to HSBs should be about the same at any redshift. Moreover, the conclusions about the epoch of star formation and halo formation (see the next section), and about the colour evolution must be valid for disk galaxies in general. In particular, Fig.~\ref{f11} shows that the colour evolution of LSBs is very strong. The U-B, B-V and V-I colours of the central part of the disk at high redshift are about 1~mag bluer than they are at low redshift, and the B-R colour even 1.5~mag bluer. This must be a property of disks in general, and not only of LSB disks. \subsection{Star Formation and Galaxy Formation} We have shown in section 2.1 that the size of LSB disk galaxies is such that their halos are formed in the redshift range $0.1<z<1.2$. We have also been able to estimate the formation redshift of the halo of each galaxy in the sample, and we have found a mean value $z=0.7$. In an Einstein--de--Sitter universe with $H_0=75$~kms$^{-1}$Mpc$^{-1}$, a redshift $z=0.7$ corresponds to a look--back time of $6.8$~Gyr. On the other hand, the photometric properties of LSB galaxies, in particular their colours, are such that star formation must have started approximately $9$~Gyr. Therefore, the star formation starts about 2~Gyr before the galactic halos are formed. These numbers are likely to be similar for HSB disk galaxies too, since they have colours and sizes very similar to LSBs, also suggesting a rather low formation redshift for their halos. The discrepancy between the epoch when star formation starts in a galaxy and the epoch when the halo forms means that the process of star formation starts before the galactic halo has assembled, as predicted for example by bottom--up galaxy formation scenarios, characterized by power spectra with more power on sub--galactic scales than on galactic ones. In Fig.~\ref{f18} we show the evolution with redshift of the SFR. The upper panel shows the fraction of gas with time that is converted into stars at every redshift for the nucleus and the disk. The continuous thick line shows the average of the nucleus and the disk as a representative of the whole LSB. It is worth noticing that most of the gas in the nucleus is converted into stars between $2 < z < 4$, while most of the star formation in the disk takes place at $ z < 2$. The lower panel shows the global SFR in the Universe for LSBs assuming that they have the same comoving number density as HSBs and that the typical mass for a $L_*$ LSB is a factor 5 smaller than the typical mass of a normal $L_*$ galaxy. The data points are taken from \scite{Madau_97} and trace the global SFR in the Universe for HSBs. LSBs lie below the observed points because they have typical masses smaller than HSBs. The important point to notice here is that the global SFR for LSBs has a flatter evolution than the observed points. \begin{figure} \centering \leavevmode \epsfxsize=1.0 \columnwidth \epsfbox{sfr.eps} \caption[]{Top panel shows the gas fraction with time in our LSB model that is converted into stars $vs$ redshift for the disk and the nucleus. The thick line corresponds to the average of the nucleus and the disk and should be a fair representative of the whole LSB. The bottom panel shows the predicted global SFR in the Universe due to LSBs assuming that the mass for a $L_*$ LSB is 5 smaller than the mass for a $L_*$ HSB and that the comoving number density for LSBs is the same as for HSBs (see text). The observed points are taken from \scite{Madau_97}. Notice that the evolution with redshift is flatter than the one measured by \scite{Madau_97}.} \label{f18} \end{figure} \subsection{Discussion} In order to explain the rotational properties of LSBs, we require halos with velocity dispersions $\approx 100$ km s$^{-1}$, or masses $\approx 10^{11}$ M$_\odot$. This is also consistent with the required gas masses and a baryon fraction of a few percent. To obtain the right sizes of LSBs, we additionally require a relatively late halo collapse, $z \simeq 0.7$. Clearly this may present a problem for galaxy formation models which have a lot of power on the appropriate scale. An obvious example is COBE-normalised CDM, with shape parameter $\Gamma=0.5$ \cite{FW91}, for which the characteristic collapse redshift for $100$ km s$^{-1}$ halos is around $z=6$. Reducing the amplitude to that of standard CDM (bias parameter 2.5) alleviates the problem to a certain extent, but the collapse redshift is still $z=3$. It is more promising to consider mixed dark matter models, for which the power spectrum is in any case a better match to the observed galaxy spectrum. For a neutrino fraction of 30\%, the comoving number density of halos with mass exceeding $1.5 \times 10^{11}$ M$_\odot$ reaches 0.01 $h^3$ Mpc$^{-3}$, comparable to the density of large galaxies, at a redshift of about 1.7 \scite{KBHP95}. This is still a little high, but it should be remembered that the collapse of halos of given mass does take place over a reasonably wide range of redshifts, and some fraction of halos would form after $z=0.7$. In the model considered by \scite{KBHP95}, about a quarter of halos exceeding this mass form after $z=0.7$. Analytically, there is expected to be a weak correlation between high spin and late collapse \scite{HP88}, and numerical simulations by \scite{Ueda94} show an anticorrelation between spin and density. We postulate therefore that the LSBs form in halos with high spin which form relatively late, in the context of a hierarchical model. The indications are that a model with relatively little small-scale power, such as mixed dark matter, would do best here, but there are enough uncertainties in the collapse redshift, both in the simplifications of the halo model, and in physical processes which might delay collapse (e.g. \pcite{BabulRees92}) that this conclusion is not strong. In summary, the most important results of this work are: \begin{enumerate} \item Observed colour profiles, chemical abundances and surface brightness profiles for LSBs are well fitted if they are assumed to have a spin parameter for its halo higher than HSBs. \item LSBs are not young objects, as often claimed in the literature, since their colours are well fitted by old ($> 7$ Gyr) stellar populations. \item There is discrepancy between the photometric age of the galaxies and the age of formation of their halo, which indicates that star formation can start about 2~Gyr before the halo is formed. This is perfectly acceptable in hierarchical models of galaxy formation, and, as discussed in the text, the discrepancy can be reduced or removed if the cosmological model we have assumed is incorrect. If an open ($\Lambda$) cosmology had been adopted, the above discrepancy would have been completely removed. \end{enumerate} | 98 | 4 | astro-ph9804049_arXiv.txt |
9804 | astro-ph9804339_arXiv.txt | The variable star population of the galactic globular cluster NGC 1851(C0512-400) has been studied by CCD photometry, from observations made in the B, V and I bands during 1993-1994. Light curves are presented for 29 variables, seven of which are new discoveries. The behavior of the RR Lyraes in the period-temperature diagram appears normal when compared to clusters which bracket the NGC 1851 metallicity. Reddening and metallicity are re-evaluated, with no compelling evidence being found to change from the values of $E(B-V) = 0.02$ and $[Fe/H] = -1.29$ (Zinn scale) adopted in recent studies of the cluster. Photometry is provided for stars in an annulus with radii 80 and 260 arcsec centered on NGC 1851. To at least $V=18.5$ there is excellent agreement with the extensive earlier photometry for the brighter NGC 1851 stars, with systematics less than 0.02 mag in all colors. Instability strip boundary positions for several clusters shows a trend for the red boundary to move to redder colors as the metallicity increases. keywords: globular clusters: individual (NGC1851) - RR Lyrae variable | The galactic globular cluster (GC) NGC 1851 (C0512-400) is rich, centrally-condensed and belongs to the small group of clusters which display bimodal horizontal-branch (HB) morphology, defined (Catelan et al 1998) as having fewer RR Lyrae stars than either blue or red HB stars. Canonical theory, that is considering a GC as a population characterized by a single age, constant abundance and a red-giant branch (RGB) mass loss parameter that has a narrow Gaussian distribution (typically $\sigma_M \sim 0.02 M_{\sun}$), cannot explain such unusual HB morphology. In order to account for the bimodality, attention has focussed recently on scenarios which can alter the mass loss parameter, such as tidal stripping of red-giant envelopes in dense environments, rapid rotation, stellar encounters, and binary interactions. Sosin et al (1997) discuss these various options in the context of the most extreme example known of a GC with bimodal HB, NGC 2808, which displays a blue HB with multiple gaps that extends to below the main sequence turnoff in the $V, B-V$ color-magnitude diagram (CMD), $M_V \sim 5$. They conclude that none of the present explanations are a satisfactory match to the observations. However, Sweigart \& Catelan (1997) have modeled the unusual HB morphology of the metal rich GC's NGC 6338 ($[Fe/H] = -0.60$) and NGC 6441 ($[Fe/H = -0.53$) for which Rich et al (1997) have obtained CMD's using the Hubble Space Telescope. Both clusters have HB's which slope upwards (brighter) with decreasing $B-V$, and have extended blue tails. Models with high helium abundance, rapid rotation, and helium mixing into the envelope are all able to produce a sloping HB morphology, and sometimes a bimodal distribution. The helium mixing alternative is particularly interesting given the observed heavy-element abundance variations in globular cluster red-giant stars (Kraft 1994). Mixing deep enough to produce enhanced aluminium, as observed in some stars, will also dredge up helium. Extensive deep mixing might be expected to destroy the sharp boundary corresponding to the deepest penetration of the convective zone, and thus prevent the observational pile-up of stars on the RGB near the level of the HB. NGC 1851 in fact appears to have quite a prominent such clump, thus suggesting that deep mixing has not taken place on the RGB at a luminosity less than that of the clump. Notwithstanding, the number of possible options still available to explain the peculiar NGC 1851 HB is considerable. The CMD of NGC 1851 has most recently been studied by Walker (1992) (hereafter W92) in the $B$ and $V$ bands, in the UV by Parise et al (1994) and in the $V$ and $I$ bands by Saviane et al. (1997) (hereafter S97), where references to earlier work can be found. In both optical studies the bimodal HB is interpreted as a consequence of differing efficiencies of mass loss as the stars evolve up the RGB. W92 suggested that a unimodal mass distribution might be able to produce a bimodal HB stellar distribution when the detailed shape of the evolutionary tracks was taken into account, however S97 do not find good agreement when comparing with the Bertelli et al. (1994) tracks. They prefer a bimodal mass loss distribution, and indeed find some evidence to suggest that the radial distributions of the blue and red HB stars differ, pointing towards some as yet unexplained interaction between the dynamical evolution and the stellar evolution of these stars. On the other hand, Catelan et al (1998), using updated (Sweigart 1997) Sweigart and Gross (1976, 1978) models, find they can reproduce the NGC 1851 HB morphology with a unimodal, albeit very wide, mass distribution having characteristics $<M_{HB}> = 0.665 M_{\sun}$ and $\sigma_M = 0.055 M_{\sun}$. Stetson et al. (1996) and Sarajedini et al. (1997) review the question of the relative ages of globular clusters, to reach very different conclusions. In both cases the HB level of NGC 1851 (W92) is used to link the second-parameter pair NGC 288 and NGC 362. Critical to these arguments is the V magnitudes of the reddest BHB stars and the RR Lyraes; the latter will be provided here for the first time. Lying in the region of the HB which is sparsely populated, the RR Lyraes may provide important clues to help explain the reason for the bimodal HB, given the constraints that the pulsation properties place on stellar evolutionary status. There have been suggestions (Catelan 1997) that the 1851 variables are peculiar with respect to their behavior in the period-temperature diagram and that the photographic studies, as detailed below, also show that several of the RRab stars have light curve amplitudes near 2 magnitudes, much larger than normal. Sawyer Hogg (1973) lists 10 variables in NGC 1851, from discoveries by Bailey (1924) and Laborde and Fourcade (1966). Preliminary periods for some of these stars, and an additional four new discoveries, resulted from a short observing campaign by Liller (1975) using photographic photometry at the CTIO 1.0-m and 1.5-m telescopes. She noted that V2 and V8 appeared to be constant, and V9 was very red. Wehlau et al. (1978) (hereafter W78) measured an additional 57 plates, almost all taken with the 1.0-m Swope telescope at Las Campanas, and analysed them along with the Liller (1975) plates. Periods were derived and light curves presented for a total of 19 RR Lyrae stars, the mean period of the RRab stars being 0.573 days, and the ratio of the number of RRc to RRab variables was found to be 0.36, both are values typical of an Oosterhoff type I system. The photometric zeropoint calibration for these observations was very uncertain, due to the lack of a definitive photometric sequence in the field. Stetson (1981) found four additional variables and two apparently constant stars lying in the instability strip. Wehlau et al. (1982) (hereafter W82) studied these stars using their original plate material supplemented with another 18 plates taken in 1970, whereupon three stars were found to be RR Lyraes, a fourth was classified as a probable field W UMa star, while a fifth is a red variable. W78 and W82 determined accurate periods and approximate mean $<B>$ magnitudes for all the stars they identified as RR Lyraes, classifying 15 stars as RRab and seven as RRc. The light curves display the $0.1-0.2$ mag scatter typical for the photographic technique. No modern studies of the NGC 1851 variables appear in the literature, however S97 identify an additional seven candidate RR Lyrae variables from a comparison of their photometry with that of W92, selecting stars with deviant photometry and $V \sim 16, V-I \sim 0.4$. Here we present new CCD photometry for the NGC 1851 RR Lyrae variables and compare with results for RR Lyraes in other GC's in this program (IC 4499, M68, M72, NGC 6362). | 98 | 4 | astro-ph9804339_arXiv.txt |
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9804 | astro-ph9804027_arXiv.txt | The Hubble Deep Field-South was chosen to have a QSO (RA 22:33:37.6 Dec $-$60:33:29 J2000 and B=17.5) in the field to allow for studies of absorption systems intersecting the sight line to the QSO. To assist in the planning of HDF-S observations we present here a ground-based spectrum of the QSO. We measure a redshift of $z=2.24$ for the quasar and find associated absorption in the spectrum at $z=2.204$ as well as additional absorption features. | Unlike the original Hubble Deep Field, the Hubble Deep Field South (HDF-S) was chosen specifically to contain a $z>2$ QSO suitable for studying the relationship between the high redshift galaxies identified in the HDF-S and the absorption lines in the spectrum of the HDF-S QSO. The QSO was found on a UK Schmidt Telescope objective prism plate scanned by Mike Irwin using the Automated Plate Measuring facility in Cambridge analysed by Paul Hewett and then confirmed by observations at the Anglo-Australian Telescope (Boyle 1997). To aid future observations of the HDF-S we present here a low-resolution spectrum of the QSO; the data are available at http://bat.phys.unsw.edu.au/$\sim$kms/hdfs/. | 98 | 4 | astro-ph9804027_arXiv.txt |
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9804 | astro-ph9804211_arXiv.txt | We use the measures of Li and rotational velocities in F Hyades stars to assess the role of the wind-driven meridian circulation and of shear turbulence in the transport of angular momentum in stars of different masses. Our models include both element segregation and rotation-induced mixing, and we treat simultaneously the transport of matter and angular momentum as described by Zahn (1992) and Maeder (1995). We show that the hot side of the Li dip in the Hyades is well explained within this framework, which was also successfully used to reproduce the C and N anomalies in B type stars (Talon et al. 1997). On the cool side of the dip, another mechanism must participate in the transport of angular momentum; its efficiency is linked to the depth of the surface convection zone. That mechanism should also be responsible for the Sun's flat rotation profile. | During the last decade, special efforts have been devoted to improve the description of the mixing processes related to stellar rotation. The most recent works (see for example Pinsonneault et al. 1989, Zahn 1992, Maeder 1995, Talon \& Zahn 1997) describe the evolution of the internal distribution of angular momentum in a self-consistent manner under the action of meridional circulation and of shear turbulence. The mixing of chemicals is then linked directly to the rotation profile, whereas previous studies made use merely of a parametric relation between the turbulent diffusivity and the rotational velocity (cf. e.g. Schatzman et al. 1981, Zahn 1983). Such a self-consistent treatment was applied successfully by Talon et al. (1997) in the study of a 9 M$_{\odot}$ star, modeling the transport of angular momentum by the meridional circulation as a truly advective process. The only assumption in this theory is that the turbulence sustained by the shear is highly anisotropic and relies on two free parameters; the first one describes the magnitude of the horizontal shears (cf. Zahn 1992) and the second one, the erosion of the restoring force due to both the thermal and the mean molecular weight stratifications (cf. Maeder 1995, Talon \& Zahn 1997). These authors reproduce the slight abundance anomalies measured in B stars by Gies \& Lambert (1992). They also show that the widening of the main sequence, which is generally attributed to convective overshooting in massive stars, may be due to the rotational mixing present in stars having a ``typical'' velocity for the spectral type considered. Concerning low-mass stars, it has been shown that the hydrodynamical models relying on meridional circulation and shear fail to reproduce the solar rotation profile given by the helioseismic observations (Brown et al. 1989, Kosovichev et al. 1997): at the solar age, those models still have large $\Omega$ gradients which are not present in the Sun (see Chaboyer et al. 1995 and Matias \& Zahn 1997). That conclusion has been reached independently by two different groups, using different descriptions for the transport processes. On one hand, the Yale group computed the evolution of angular momentum in low mass stars with a simplified description of the action of the meridional circulation which was considered as a diffusive process rather than as an advective process. The whole evolution of momemtum and chemicals was then due to diffusion only, with a free parameter that had to be calibrated to differentiate the transport of the passive quantities with respect to that of vectorial ones. Pinsonneault et al. (1990) were then able to reproduce the surface Li abundances for low-mass cluster stars (with effective temperature lower than 6500K). However, they obtained large rotation gradients within these stars which are excluded by helioseismology (Chaboyer et al. 1995). On the other hand, Matias \& Zahn (1997) performed a complete study for the evolution of the Sun's angular momentum, where they took into account the advective nature of the meridional circulation. They also concluded that meridional circulation and shear turbulence are not efficient enough to enforce the flat rotation profile measured by helioseismology. These results indicate that another process participates in the transport of angular momentum in solar-type stars, while the so-called wind-driven meridional circulation (Zahn 1992) is successful in more massive stars. In order to study the transition between solar-type and more massive stars and to identify the mass range for which the present description for the transport of angular momentum and chemicals relying only on rotation fails, we propose to use the measures of lithium and rotational velocities in galactic cluster stars. We first review the observations of lithium abundances and rotation in the Hyades main-sequence stars, and summarize the difficulties of the various models proposed so far to explain the Li dip in F stars (\S 2). We recall the equations that describe the evolution of angular momentum due to meridian circulation and shear turbulence as well as the associated transport of chemicals (\S 3). We study the impact of rotational mixing on the lithium abundance in galactic cluster F stars, and compare this to the observations. Our models include both element segregation and rotation-induced mixing, and we treat simultaneously the transport of matter and angular momentum. The internal rotation profile thus evolves completely self-consistently under the action of meridional circulation as described by Zahn (1992) (see also Matias et al. 1997), and of shear stresses which take into account the weakening effect of the thermal diffusivity, as was first shown by Townsend (1958) (\S 4). We show that the blue side of the lithium dip is well reproduced within this framework, and that the process responsible for the shape of the solar rotation profile should become efficient only for stars on the cool side of the Li dip, where the external convection zone is thick enough. By achieving efficiently momentum transport, the global effect of this process would be to reduce the mixing due to the rotational instabilities in stars with effective temperature lower than $\sim$ 6500K. The most likely candidates for this transport process are the gravity waves generated by the external convection zone (Schatzman 1993; Zahn et al. 1997; Kumar \& Quataert 1997) and the large-scale magnetic field which could be present in the radiative interior (Charbonneau \& MacGregor 1993; Barnes et al. 1997). | Assuming rapid rotation and using a self-consistent description for the transport of angular momentum and of chemicals by meridional circulation and shear instabilities (cf. Zahn 1992, Talon \& Zahn 1997), Talon et al. (1997) successfully explained the C and N anomalies observed in some B stars. At the same time, it was shown (Matias \& Zahn 1997) that this description applied to the transport of angular momentum in the Sun is incomplete, leading to large $\Omega$ gradients which are not observed. Another transport mechanism must thus be invoked in low mass stars. At this point, 2 questions remain : firstly, the nature of that transport mechanism has to be determined unambiguously and secondly, the location of the transition between the regime which is relevant for massive stars and the one which is relevant for low mass stars has to be identified. In this paper, we addressed that second question. We presented numerical calculations of Li destruction due to rotational mixing using the {\it same} description as Talon et al. used for more massive stars and the same free parameters. We showed that this clearly reproduces the hot side of the Li dip. Let us recall that the destruction of lithium is then due solely to rotational mixing enhanced by the spin down of the outer layers. Stars hotter than 7000 K also undergo rotational mixing, but it is much milder due to the weak differential rotation. The rise of Li abundances on the right side of the dip is not explained within this framework. We propose that it is linked to the appearance of another transport mechanism for angular momentum which reduces the magnitude of the meridional circulation and shears, leading to the observed diminution of Li destruction on the red side of the Li dip. This mechanism is known to occur in the Sun where it is responsible for the flat rotation profile. | 98 | 4 | astro-ph9804211_arXiv.txt |
9804 | astro-ph9804161_arXiv.txt | The Berkeley spectrograph aboard the ORFEUS telescope made its second flight on the 14-day ORFEUS-SPAS II mission of the Space Shuttle {\it Columbia} in November/December 1996. Approximately half of the available observing time was dedicated to the Berkeley spectrograph, which was used by both Principal and Guest Investigators. The spectrograph's full bandpass is 390--1218 \AA; here we discuss its in-flight performance at far-ultraviolet (FUV) wavelengths, where most of the observations were performed. The instrument's effective area peaks at 8.9 $\pm$ 0.5 cm$^2$ near 1020 \AA, and the mean spectral resolution is 95 km s$^{-1}$ FWHM for point sources. Over most of the spectral range, the typical night-time background event rate in each spectral resolution element was about 0.003 s$^{-1}$. Simultaneous background observations of an adjacent blank field were provided through a secondary, off-axis aperture. The Berkeley spectrograph's unique combination of sensitivity and resolution provided valuable observations of approximately 105 distinct astronomical targets, ranging in distance from the earth's own moon to some of the brightest AGN. | The German spacecraft Astro-SPAS, a deployable platform designed to meet the technical performance demands of astronomical payloads and similar scientific instruments, comprised the primary payload aboard shuttle mission STS-80 ({\it Columbia}). On this, its third flight, the platform carried a trio of far-ultraviolet instruments: two independent spectrographs within the 1 meter diameter ORFEUS telescope (Grewing et al. 1991) and the IMAPS objective-grating spectrograph (Jenkins et al. 1996). All three had flown on the Astro-SPAS' 5-day maiden voyage in September of 1993, but improvements in instrument performance, and the critical need for additional observation time, motivated a reflight. A photograph of the payload is shown in Plate~1. With few exceptions, targets suitable for ORFEUS were too faint for IMAPS, so no attempt was made to coalign these instruments closely. Within the ORFEUS telescope a flip mirror was employed to direct the optical beam to one spectrograph or the other. Hence in general only one of the three instruments was operated at a time. The available observing time was shared equally between Guest Investigators selected by peer review and the Principal Investigator teams who had provided the instruments. Flight operations were directed from a control complex at the Kennedy Space Center. The general design of the Berkeley spectrograph has been discussed previously (Hurwitz \& Bowyer 1986, 1996). We changed the instrument between missions only by overcoating of two of the four diffraction gratings (including the far ultraviolet grating) with silicon carbide, introducing the multiple apertures discussed below, and modifying the detector electronics to improve the imaging at high count rates. We did not recoat the KBr photocathode on the microchannel plate detectors; the delay-line anode detector systems are discussed in Stock et al. (1993). In this work we report on the performance and calibration of the spectrograph during the ORFEUS-SPAS II mission and the instrumental effects of interest to Guest Investigators and other users of the extracted data products. | The Berkeley spectrograph aboard the ORFEUS telescope offers a unique and important combination of spectral resolution and effective area in the comparatively unexplored far-ultraviolet wavelength band. During the ORFEUS-SPAS II mission in November/December 1996, Principal and Guest Investigators utilized the spectrograph to observe some 105 astronomical targets. These data will enter the public domain in early to mid 1998. Lyman/FUSE is not far from launch, and will offer a much higher spectral resolution and sensitivity for point sources. However, near-trivial modifications would enable the Berkeley spectrograph to achieve a significantly superior performance for studies of extended emission. New replicas of the current gratings and SiC overcoating of the primary mirror would allow the spectrograph to achieve 2 \AA\ slit-limited resolution through a 45 $\times$ 420\arcsec\ slit and an effective area of $\sim$ 50 cm$^2$ across the 900 -- 1200 \AA\ band. The angular resolution along the slit length would be better than 1 \arcmin. Such an instrument would be highly desirable for studies of intracluster gas, galaxies and their halos, supernova remnants, and other extended objects. At time of writing, there are no specific plans for a third flight of the payload. | 98 | 4 | astro-ph9804161_arXiv.txt |
9804 | astro-ph9804283_arXiv.txt | In this paper we propose a new statistic capable of detecting non-Gaussianity in the CMB. The statistic is defined in Fourier space, and therefore naturally separates angular scales. It consists of taking another Fourier transform, in angle, over the Fourier modes within a given ring of scales. Like other Fourier space statistics, our statistic outdoes more conventional methods when faced with combinations of Gaussian processes (be they noise or signal) and a non-Gaussian signal which dominates only on some scales. However, unlike previous efforts along these lines, our statistic is successful in recognizing multiple non-Gaussian patterns in a single field. We discuss various applications, in which the Gaussian component may be noise or primordial signal, and the non-Gaussian component may be a cosmic string map, or some geometrical construction mimicking, say, small scale dust maps. | Current theories of structure formation may be roughly divided into two classes: active and passive perturbations. According to the inflationary paradigm, quantum fluctuations in the very early universe are produced during a period of inflation \cite{cosmcross} and grow to become classical density perturbations \cite{lidlyth,bardeen}. These perturbations evolve linearly until late times when the overdensities become galaxies. Perturbations due to inflation are called passive because they are seeded at some initial time near the Planck time and then evolve `deterministically', or linearly. They leave their imprint on the CMB at last scattering \cite{pyu,nature,HS1,HS2,richard,1.3}. Although this is not strictly necessary, in most cases the fluctuations in the CMB temperature due to inflationary perturbations form a Gaussian random field \cite{bondefst,bardstat}. There is another class of theories of structure formation, topological defects caused by phase transitions in the early universe \cite{Kib,vv}. Perturbations caused by defects are known as active perturbations \cite{mafc,coher} since they are continually being seeded by an evolving network of defects through the history of the universe. The fluctuations in the CMB in defect models have been found to be non-Gaussian \cite{phases}, even though the extent and strength of this non-Gaussianity is still far from clear. Recently a wide class of defect models have been shown to be in conflict with current data \cite{neil,abr,abr2} but some viable active models still remain \cite{mimic,durrer,abrw}. In fact, some of the most interesting \cite{abr3,shellard} of these (based on cosmic strings) require a non-zero cosmological constant of the sort that currently favoured by supernova experiments \cite{sn1}. Thus it is important to be able to distinguish Gaussian from \ngn fluctuations. Many different tests are being tried \cite{Kogut,XLuo,fermag,cumul,gorski,sergei}, adapted to different experimental settings, and types of signal. Our statistic is designed to be used for small fields where one or more distinctive shapes are obscured by an extra Gaussian component, and so are not visible in real space. It is our experience in previous work \cite{fermag} that in such situations standard statistics, based in real space, fail to recognize the non-Gaussianity of the signal. The idea is to study the statistical properties of map derivatives which are only sensitive to a given scale. In this way we may separate out different scales, some of which may have Gaussian or nearly Gaussian fluctuations, some very non-Gaussian. Although this scale filtering may be achieved using the wavelet transform \cite{cumul,pando}, in this paper we choose to use the Fourier transform as our scale filter. One problem with this approach is that the Fourier transform is a global transformation, and therefore can only recognize non-Gaussian structures globally. It may therefore offer a rather contorted description for complicated networks made up of essentially simple objects. We will however perform a transformation over the Fourier modes: a second Fourier transform, in angle, for modes in each ring in Fourier space. We will argue that by doing so we are able to recognize mostly features of individual objects, and so bypass this problem. This operation returns a set of quantities which are blind to the random orientations of individual objects. Unfortunately their random positions still affect our statistic, so there will be a limit (albeit less stringent) on the number of objects in the field before our statistic becomes confused. The plan of this paper is as follows. In Section II we look at some of the ideas involved in thinking about non-Gaussianity. In Section III we define our statistic and look at some of its properties and motivations. In Section IV we apply our statistic to practical situations in which subtle non-Gaussian signals are present. We consider non-Gaussian signals corrupted by the presence of a Gaussian process, which can be noise or primordial signal, and which dominates on all but a narrow band of scales. We consider two types of non-Gaussian signal. We consider CMB maps obtained from string simulations, implementing the algorithms in \cite{av}. We also consider geometrical constructions mimicking, say, small scale dust maps. We show how such maps look very Gaussian, but fail to confuse our statistic. | We have introduced a statistic to look for \ng in the cosmic microwave background. It is designed to pick out \ngn features which are superimposed on Gaussian fluctuations and which are therefore not visible in real space. Since some scales may be more \ngn than others we choose to separate them out. It is our experience in previous work \cite{fermag} that in such situations standard statistics, based in real space, fail to recognize the non-Gaussianity of the signal. Our statistic is naturally tailored for interferometric experiments, which make measurements in Fourier space. Indeed the algorithm exposed above could easily be turned into a data analysis package operating over visibilities. One of the problems with Fourier space statistics is that they recognize only global shapes. Therefore they become very ineffective when many individual structures are present. By taking another Fourier transform, in angle, over modes in a given ring in Fourier space, we can factor out orientations, and be only sensitive to an average shape of individual structures. Their random positions, on the other hand, will eventually make our statistic very ineffective as the number of objects becomes large. We found that although we have improved on previous work, still there is a limit on how many individual structures there may be in the field before the statistic fails to pick out their non-Gaussianity. Finally, some comments are in order regarding the applications we have considered. The Gaussian component we have considered may include noise, if it has the right power spectrum. Hence we have already considered the effects of a simplified form of noise: it will merely reduce the band of scales where the non-Gaussian signal dominates, typically providing it with an upper boundary. An exception is the case of non uniform noise in the u-v plane, present in most interferometer experiments. Non uniform noise in the Fourier domain will look non-Gaussian, for it is a Gaussian process which is not isotropic or translationally invariant. Therefore an extra element of confusion, not dealt with in this paper, will appear. Regarding the non-Gaussian component we have considered two types of signal. We have looked at some of the properties of the statistic when it is applied to a field containing one or more distinct shapes. The geometrical constructions presented are somewhat reminiscent of the jagged structures present in small scale dust maps. Experimenting on fields containing many structures we saw our statistic rapidly become blind to their non-Gaussianity. However interferometer fields are often small. They would therefore contain only a few, but more than one, of these structures. This is something our statistic can cope with, unlike previous Fourier space based statistics. We have also studied our statistic when applied to cosmic string maps. We found that some individual small fields showed \ng in the signal. If the results were averaged over many fields the effect of \ng on $|F(\beta)|$ would not show. Once again this strategy is ideal for interferometers, for which combining small fields into a large field is in fact a rather awkward operation. | 98 | 4 | astro-ph9804283_arXiv.txt |
9804 | astro-ph9804230_arXiv.txt | La d\'ecouverte en 1965 par Penzias et Wilson du fond de rayonnement cosmologique diffus \`a 3 K a d\'efinitivement fait du mod\`ele du Big Bang chaud le cadre g\'en\'eral de la cosmologie moderne. Ces trois derni\`eres d\'ecennies ont \'et\'e marqu\'ees par de nombreux progr\`es tant th\'eoriques qu'observationnels qui n'ont fait que confirmer sa validit\'e, et de pr\'eciser petit \`a petit les d\'etails de ce sc\'enario. Parmi les avanc\'ees notables de ces derni\`eres ann\'ees, citons pour m\'emoire: \begin{itemize} \item La r\'ealisation de grands catalogues de galaxies. On a maintenant des catalogues tri-dimensionnels qui contiennent des milliers d'objets. Ces catalogues permettent de faire une v\'eritable cosmographie de l'univers local. De nouveaux moyens sont apparus r\'ecemment pour compl\'eter ce type de catalogues, la d\'etermination des champs de vitesse cosmiques, permet ainsi d'acc\'eder \`a des informations cin\'ematiques tr\`es pr\'ecieuses. Et depuis la fin des ann\'ees 80, un nouveau moyen d'investigation est en train d'\'emerger, il s'agit des cartes de distorsion gravitationnelle. Elles permettent de visualiser les lignes de potentiel de la masse projet\'ee. \item le mod\`ele de Mati\`ere Noire Froide (CDM pour Cold Dark Matter) qui est apparu au d\'ebut des ann\'ees 80 et qui a servi de point de r\'ef\'erence (\`a d\'efaut de devenir un mod\`ele standard) pour tous les travaux sur le probl\`eme de la formation des grandes structures. \item Le d\'eveloppement des th\'eories inflationnaires. C'est encore un terrain tr\`es sp\'eculatif, mais c'est le lieu tr\`es excitant o\`u les concepts de la physique des hautes \'energies rencontrent des exigences observationnelles de plus en plus fiables. \item Enfin, la d\'etection en 1992 des fluctuations de temp\'erature du fond de rayonnement cosmologique par l'exp\'erience satellitaire COBE/DMR (Smoot et al. 1992) a marqu\'e un tournant pour la cosmologie: pour la premi\`ere fois on avait une preuve directe de l'origine des grandes structures de l'univers. \end{itemize} Dans ce cours je vais principalement m'int\'eresser au probl\`eme de la formation des grandes structures. Quelques ouvrages de r\'ef\'erence, \begin{itemize} \item Relativit\'e G\'en\'erale: Weinberg, {\it Gravitation and Cosmology}, 1972; Landau et Lifschitz, {\it Classical Theory of Fields}, 1975 \item Univers primordial: Kolb et Turner {\it The Early Universe}, 1990, mais il est peu pr\'ecis; revue de Brandanberger, {\it Inflation and Cosmic Strings: two Mechanisms for producing Structure in the Universe}, Int. J. Mod. Phys. A2 : 77, 1987. \item Inflation: Linde, {\it Particle Physics and Inflationary Cosmology}, 1990; Liddle et Lyth, {\it The Cold Dark Matter density perturbation}, 1993, Physics Reports, 231, 1; Lidsey et al. {\it Reconstructing the inflaton potential - an overview}, 1997, Reviews of Modern Physics, 69, 2 \item Formation des grandes structures: Peebles {\it The Large Scale Structure of the Universe}, 1980; {\it Principle of Physical Cosmology}, 1993; V. Sahni \& P. Coles, {\it Approximation Methods for Non-linear Gravitational Clustering}, 1995, Physics Reports, 262, 1 \end{itemize} | Dans ce petit tour d'horizon de la cosmologie j'ai \'et\'e loin d'\^etre exhaustif. Dans le domaine de l'\'evolution des grandes structures je n'ai pas mentionn\'e bon nombre de domaines de recherche en plein d\'eveloppement, comme \begin{itemize} \item L'\'etude des flots cosmiques \`a grande \'echelle, et leur utilisation pour mesurer les param\`etres cosmologiques; \item L'\'etude des amas de galaxies et de leur contenu aussi bien en mati\`ere noire (\'etudes dynamique, reconstruction de masse par effets de lentilles gravitationnelles..), qu'en mati\`ere baryonique (galaxies, rayonnement X..); \item L'\'etude des absorbants et des objets de grand $z$; \item ... \end{itemize} Du point de vue du physicien th\'eoricien un domaine de recherche de pr\'edilection est inconstablement la cosmologie primordiale, lieu de rencontre entre entre la physique des hautes \'energies et la cosmologie observationnelle. Cependant la formation des grandes structures recelle un certain nombre de probl\`emes ouverts digne d'int\'er\^et, \begin{itemize} \item La th\'eorie des perturbations est loin d'avoir livr\'e tous ses secrets. En particulier il serait extremement utile de comprendre les corrections en boucles. \item De nombreux aspects du r\'egime non-lin\'eaire ne sont pas compris. Par exemple on ne sait pas d\'ecrire la transition vers le r\'egime multiflots qui conduit \`a la virialisation de la mati\`ere dans les puits de potentiel. Il serait aussi tr\`es int\'eressant de pouvoir exhiber une solution explicite des \'equations dynamiques dans le r\'egime nonlin\'eaire (m\^eme si ce n'est qu'une forme assymptotique). \item L'exploitation des catalogues tridimensionnel ou bidimensionnel notamment avec la mise en \'evidence de propri\'et\'es non-Gaussiennes n'est pas encore optimale. La question se pose par exemple pour les cartes de distorsion gravitationnelle. \item La relation entre baryons et mati\`ere noire est un domaine qui n'a pratiquement pas \'et\'e explor\'e. On ne dispose pour l'instant que d'exp\'eriences num\'eriques pour essayer de comprendre ce qui se passe. \end{itemize} | 98 | 4 | astro-ph9804230_arXiv.txt |
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9804 | astro-ph9804006_arXiv.txt | The Lyman edge at 912 {\AA} is an important diagnostic region for studying quasi-stellar objects (QSOs). In particular, it reveals a great deal about the physical conditions within the atmospheres of accretion disks, a ubiquitous component of QSO theories. A robust prediction of accretion disk models is a significant polarization due to electron scattering just longward (in wavelength) of the Lyman edge because of the wavelength dependence of the Hydrogen absorption opacity. Observations of the Lyman edge regions of QSOs have shown scant evidence for the predicted features---few QSOs show the broad, partial Lyman edges expected to be common according to most theories, and none show the high polarizations expected longward of the Lyman edge. Still, polarization spectra of a small number of QSOs have shown a rising polarization (up to 20\%) at wavelengths {\it shortward} of the Lyman edge. We have now doubled our sample of intermediate-redshift QSOs observed with the {\it HST/FOS} spectropolarimeter to determine the amount of polarization on both sides of the Lyman limit. For this new sample of six objects, polarizations are low and mostly consistent with zero below the Lyman edge. Another important result of the new data is that it strengthens the conclusion that quasars are generally not polarized significantly just longward of the Lyman edge at $\sim$1000\AA\/. There is no significant statistical wavelength dependence to the polarization longward of the Lyman edge indicating that simple plane-parallel atmospheres with scattering-dominated opacity are not significant sources of UV flux in quasars. | One of the fundamental components of most theories of quasi-stellar objects (QSOs) is an accretion disk. As gas is fed into the central regions of the QSO, residual angular momentum causes the gas to naturally settle into a disk. While most theories predict the formation of such a disk, few address the observational consequences of the disk models in detail. These theoretical studies have found that the Lyman edge at 912 {\AA} is a powerful diagnostic feature for the physical characteristics of the disk. The simplest disk models (quasi-static with viscous dissipation at large optical depth) predict Lyman edges in either emission or absorption, depending on the viewing angle and the physical details of the disk atmosphere. Such Lyman edges would be broadened by rotation of the disk and by general relativistic effects as light passes close to the central black hole. In most QSOs such edges are not seen, although Koratkar, Kinney, \& Bohlin (1992) found a small number of candidate ``partial edges'' in IUE data. A second disk signature is the linear polarization, $P$, of the continuum from the disk. A purely scattering atmosphere will produce high polarization perpendicular to the disk axis. Again, this signature is not seen in any QSOs; in fact generally QSOs show {\it low} optical polarization {\it parallel} to the inferred disk axis. Laor, Netzer, \& Piran (1990) attempted to show that a disk atmosphere should have significant absorptive opacity, and thus can produce dramatically lower optical polarization (albeit still perpendicular to the disk axis and thus inconsistent with the observations). A more robust prediction according to their work, however, is a rise in polarization with decreasing wavelength from the optical into the UV. Just longward of the Lyman edge $P$ is highest, since it is at these wavelengths that scattering best competes with absorption. Just shortward of the Lyman edge, as the absorption opacity increases, $P$ should drop again. According to Laor \etal, this polarization signature should appear even when no disk signature is seen in total flux. In our first polarization study of three of the rare objects known from IUE spectra to have partial Lyman edges at the systemic redshifts, we found low polarizations longward of the edges, so we did not confirm the Laor \ea prediction. We surprisingly did find high polarization {\it shortward} of the edge in a couple of objects, contrary to the accretion disk predictions of Laor \etal, but qualitatively explicable by effects found in the more detailed calculations of Blaes \& Agol (1996); see also Agol \& Blaes (1996). The previous studies of intermediate redshift quasars (Koratkar \ea 1995 and Impey \ea 1995), included 4 objects, one of which shows only a marginal detection of polarization, while the remaining three objects show significant polarization ($>$ a few percent) shortward of the Lyman edge. A study of three high redshift objects observed from the ground failed to show any polarization changes at the edge position; most had tight limits on the polarization longward of the edge, but noisy data shortward of the Lyman edge (Antonucci \ea 1996). PG 1630+377 is the only object yet observed that can be studied in any detail (Koratkar \ea 1995; Paper I). In this object $P$ rises rapidly shortward of the edge, reaching 20\% by 1600 {\AA} (650 {\AA} rest wavelength). The Ly$\alpha$ emission line also shows a high (7.3\%) polarization at the same position angle. Antonucci \ea (1996) discuss polarization observations and other constraints on disk models in some detail. Based on the small number of QSOs observed in the UV in polarization, at the time of paper I, we could say little about whether high polarization shortward of the Lyman edge is common. Hence, we have significantly expanded the UV polarization database by observing six bright, $z > 1$ QSOs both below and above the Lyman edge. In this paper we discuss these new spectropolarimetric ultraviolet observations from the {\it Hubble Space Telescope} Faint Object Spectrograph ({\it HST/FOS}). A difference with respect to our previous study, however, is that only two of the new objects were suspected to have partial edges in total flux at the systemic redshift. The rest were simply selected because they show significant flux at short wavelengths. | Of the six QSOs identified by Koratkar, Kinney \& Bohlin (1992) as candidate targets which have partial Lyman edges consistent with edges from simple thin accretion disks, we now have spectropolarimetric observations for five QSOs. We showed in section 3.2 that one of these candidates, 0743$-$673 no longer qualifies as a partial Lyman edge object. There are only 13 high and intermediate redshift QSOs which have spectropolarimetry observations shortward of the Lyman edge region. These objects come from this paper, Koratkar \ea (1995), Impey \ea (1995), and Antonucci \ea (1996). At this point any detailed statistical tests of polarization distributions are certainly not warranted given the inhomogeneous selection criteria and data quality, and the highly model-dependent predictions. Yet, Lyman edge spectropolarimetry results can be summarized as follows: \begin{itemize} \item{}Of the 13 objects only three objects (PG 1630+377, PG 1338+416 and PG 1222+228 from paper I and Impey \ea 1995) show significant polarization at wavelengths shorter than 912\AA\/ (Lyman edge). To these three objects we can add one more marginal detection (PKS 0405$-$123 from paper I). All 13 objects in the sample show a polarization signature which is inconsistent with any simple accretion disk model. Furthermore, $\sim$30\% of the sample show a rise in polarization shortward of the Lyman edge. This observed rise in polarization is qualitatively consistent with the disk models of Blaes \& Agol (1996). A number of different interpretations of the UV signature have been given by Lee \& Blandford (1997), Shields, Wobus \& Husfeld (1997) and by us in paper I. We urge the interested reader to consult those papers for more details. \item{}There are a total of five objects (PG 0117+213 from the present sample, PG1630+377, PG 1338+416, and PKS 0405$-$123 from paper I, and 0014+813 from Antonucci \ea 1996) which show candidate partial Lyman absorption edges due to accretion disks at the systemic redshift in total flux. Of these five objects, two (PG 1630+377, PG 1338+416) have sufficient signal-to-noise at rest wavelengths of $\leq$750\AA\/, and show significant polarization shortward of the Lyman edge (at least a few percent, detected at four sigma or greater significance). If the rise in UV polarization seen in PG 1630+377 is characteristic of objects with partial Lyman edges we need to observe rest wavelengths as short as $\sim$700\AA. The effective shortest rest wavelength observed in PG 0117+213 and PKS 0405$-$123 is $\sim$800\AA\/. Thus in these objects we could have missed the rise in polarization, although we do have a marginal detection for PKS 0405$-$123. 0014+813 does not have sufficient signal-to-noise shortward of the Lyman edge. To summarize, of the objects that show candidate partial Lyman absorption edges in total flux, $\sim$40\% show polarization shortward of the Lyman edge. \item{}We have eight objects in the sample of 13 that do not show a partial Lyman edge feature in total flux. Only one object out of these, PG 1222+228, from Impey \ea (1995) shows significant polarization shortward of the Lyman edge ($P$ = 4.6$\pm$0.9\%). The rest show no detection of polarization in the 912\AA\/ spectral region. The linear polarization upper limits shortward of the Lyman edge in the remaining objects is \ltsima 4\%. \item{}PG 1630+377 is the best studied object (see Paper I), and shows UV polarization reaching 20\% at 650\AA\/ rest wavelength. Such high degree of polarization is rare in (non-blazar) QSOs. \end{itemize} Perhaps the simplest result of the current data is that it strengthens the conclusion that quasars are generally not polarized significantly just longward of the Lyman edge at $\sim$1000\AA\/. This paper, Koratkar \ea (1995), Impey \ea (1995), and Antonucci \ea (1996) together present good observations of about 20 objects in the region just longward of the Lyman edge, all with low UV polarization (\ltsima 1.5\%). Further, there is no significant statistical wavelength dependence to the polarization as predicted by the models of Laor et al (1990). From these observations we conclude that simple plane-parallel atmospheres with scattering-dominated opacity are not significant sources of UV flux in quasars. Recapitulating the previous discussions here briefly, we note that models from the 1980s generally assumed that AGN accretion disks are powered by viscous dissipation below the atmospheres, and that the atmospheric opacities are dominated by electron scattering opacity in the annuli that produce the rest optical and UV. This results in 0\% to 11.7\% polarization (Chandrasekhar 1960), depending on inclination, and in a direction perpendicular to the symmetry axis of the disks. Pioneering optical polarimetry observations showed much smaller polarizations, which are {\it parallel} to the axes when the latter could be inferred from a radio jet position angle (Stockman, Angel \& Miley 1979; Antonucci 1988). To explain the low observed polarization, subsequent models by Laor et al (1990) suggested that electron scattering was only important in the $\sim$1000-2000 \AA\/ range, with the Lyman continuum and free-free absorption opacity dominating shortwards and longwards of that interval respectively. Other more detailed calculations revealed that a lower fraction of absorption opacity was required to reduce the predicted polarization than was assumed by Laor et al; and that under rather special circumstances a large polarization {\it parallel} to the disk axis could be produced shortward of the Lyman edge (Blaes and Agol 1996, and references therein). The observed rise in UV polarization shortward of the Lyman edge has been interpreted both in the context of accretion disk models and non-disk related models. Here we do not further discuss the polarization and depolarization mechanisms discussed in detail in Paper I. An additional key complication in the accretion disk models, is that AGN variability data require that the disk atmosphere is actually illuminated from above, (e.g. Antonucci 1988, Sincell and Krolik 1996), perhaps producing a strong polarization which cannot be calculated rigorously without specification of the illumation geometry. | 98 | 4 | astro-ph9804006_arXiv.txt |
9804 | gr-qc9804041_arXiv.txt | {The mass function of primordial black holes created through the near-critical gravitational collapse is calculated in a manner fairly independent of the statistical distribution of underlying density fluctuation, assuming that it has a sharp peak on a specific scale. Comparing it with various cosmological constraints on their mass spectrum, some newly excluded range is found in the volume fraction of the region collapsing into black holes as a function of the horizon mass. } \pacs{PACS Numbers: 04.70.Bw, 04.70Dy, 98.80.Cq } | 98 | 4 | gr-qc9804041_arXiv.txt |
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9804 | astro-ph9804140_arXiv.txt | We investigate whether models based on the assumption that jets in quasars are powered by rotating black holes can explain the observed radio dichotomy of quasars. We show that in terms of the ``spin paradigm'' models, radio-loud quasars could be objects in which the black hole's rotation rate corresponds to an equilibrium between spin-up by accretion and spin-down by the Blandford-Znajek mechanism. Radio-quiet quasars could be hosting black holes with an average spin much smaller than the equilibrium one. We discuss possible accretion scenarios which can lead to such a bimodal distribution of black hole spins. | Quasars are characterized not only by an intense radiation from the central engine, but also by jets which power large scale radio structures. The ratio of radio luminosity of these structures to the optical luminosity of the central sources shows a bimodal distribution, with only $\sim 10$ \% of quasars belonging to the radio-loud category (see, e.g., Kellerman et al. 1989; Hooper et al. 1995; Falcke, Sherwood, \& Patnaik 1996; Bischof \& Becker 1997). Radio-loud quasars have never been found in spiral galaxies, while the hosts of radio-quiet quasars are either spiral or elliptical galaxies (Taylor et al. 1996; Kukula et al. 1998). On the other hand, both radio-quiet and radio-loud quasars have almost identical average IR-optical-UV spectra (Francis et al. 1993; Zheng et al. 1997), which suggests similar accretion conditions in these two samples of objects. These properties seem to support models which are based on the idea that jets are powered by rotating black holes. However, in order to explain the observed radio loudness bimodality in terms of a bimodal distribution of black hole spins, one has to assume that the population of supermassive black holes is dominated by very low-spin black holes, and one should understand why radio-loud quasars are always hosted by elliptical galaxies. This issue has been addressed by Wilson \& Colbert \shortcite{wc95}, who proposed that high-spin black holes exist only if formed by coalescence of two black holes. Such coalescences would take place mostly in a denser galactic environment (groups, clusters), which are much more populated by elliptical galaxies than the field regions. However, the basic assumption of this scenario, that black holes which do not undergo coalescence rotate very slowly -- despite the angular momentum gained from accretion discs -- must be verified. In other words, one should verify that there exists a mechanism which could keep black holes at low rotation rates despite of the accreted angular momentum. This problem was investigated by Moderski \& Sikora \shortcite{ms96a}. They showed that the Blandford-Znajek (B-Z) mechanism, which extracts rotational energy from a black hole \cite{bz77}, is not efficient enough to counteract the spinning up of a black hole by the gas accreted from standard $\alpha$-discs \cite{ss73}. Low-spin equilibrium states are possible only for very low accretion rates but the time required to established such an equilibrium is longer than the age of the Universe. This is because the power extracted from the black hole is proportional to the square of intensity of the magnetic field which threads the black hole. The field intensity is limited by the pressure in the accretion flow and, therefore, is very low for very low accretion rates. Hence, once a black hole gets a high spin, it will be rotating fast forever. Low-spin equilibrium solutions for high accretion rates are possible only for the so called $\beta$-discs \cite{ms97}. In such discs the viscous stress driving accretion is proportional to the gas pressure only (and not to the total pressure as in the case of $\alpha$ discs). Since the relation between the viscous stress and the pressure in radiation pressure dominated accretion discs is unknown, such discs are a reasonable alternative to $\alpha$-discs. For high accretion rates the gas pressure is by several orders of magnitude smaller than the radiation pressure \cite{ss73}. This implies that, for a given accretion rate, $\beta$-discs are much denser than $\alpha$-discs \cite{sc81} and, therefore, can confine much stronger black hole magnetic fields. Since the power extracted from a black hole via the B-Z mechanism scales with the square of the magnetic field intensity, equilibrium spins for $\beta$-discs are smaller than for $\alpha$-discs. This, of course, does not mean that the power extracted from a black hole with a $\beta$-disc must be smaller than the power extracted from a black hole with an $\alpha$-disc, because in the expression for the B-Z power a lower equilibrium spin is compensated by a higher intensity of the magnetic field \cite{ms97}. Thus, low equilibrium spin black holes can represent radio-quiet quasars only, if the conversion of the extracted black hole energy into jet energy is very inefficient. According to the Wilson-Colbert scenario powerful jets would then exist only in objects where a coalescence of two black holes leads to spins much higher than the equilibrium one. However, powers extracted from black holes with spins much larger than the equilibrium one are so large, that such black holes would be spun-down on time scales two orders of magnitude shorter than the typical lifetime, $\sim 10^8$ years, of radio quasars (see, e.g., Leahy, Muxlow \& Stephens 1989). Powering jets in radio-loud quasars by a black hole can last as long as $10^8$ years, only if losses of angular momentum due to the B-Z mechanism are compensated by gains of angular momentum from an accretion disc. This would give a correlation between radio and optical luminosities, in accordance with observations \cite{ser97}. Then, however, one would have to explain why the majority of super-massive black holes have spin values much lower than the equilibrium one -- the condition for the existence of radio-quiet quasars. As suggested by Moderski, Sikora \& Lasota \shortcite{msl97}, black holes in most objects could be forced to rotate slowly by multi-accretion events with random orientations of the angular momentum vector. In such a scenario, quasars which become radio-loud are only those which undergo major accretion events, induced, i.e., by a merger of two big galaxies. Following this, a black hole can easily double its mass and reach an equilibrium state. In this paper we explore this possibility and derive conditions the model must satisfy in order to explain the observed radio properties of quasars. This paper is organized as follows. In Section 2 we present equilibrium spin solutions, obtained for a variety of accretion disc models. In Section 3, we discuss multi-event accretion scenarios. In Section 4 we use our results to derive conditions which must be satisfied by quasar evolution models in order to obtain bimodal distribution of black hole spins. | Radio dichotomy of quasars was discovered many years ago (Strittmatter et al. 1980; Kellerman et al. 1989), but is still waiting for a theoretical explanation. As for now, the consensus concerns only one aspect of the problem: it is clear that jets in quasars must be formed near a supermassive black hole. This follows from the energetics of quasar jets, since no other known sources could power jets at a rate reaching $10^{46}$ ergs s$^{-1}$ for millions of years (Rawlings \& Saunders 1991; Leahy et al. 1989). Independent argument for the formation of extra galactic jets in the vicinity of supermassive black holes is provided by direct VLBI observations of nearby radio galaxies. In particular, in 3C 274 (M87) the jet is seen down to $10^{16}$ cm from the center \cite{jb95}, which corresponds to $100$ gravitational radii for the $3 \times 10^9 M_{\odot}$ \cite{har94} central black hole. There, very deep in the gravitational potential well, jets could be powered either by the innermost parts of an accretion disc (Blandford \& Payne 1982; Park \& Vishniac 1994; Contopoulos 1995; Begelman 1995) or by a rotating black hole (Blandford \& Znajek 1977; Rees et al. 1982). However, no jet production model can be successful, if it fails to explain why only a small fraction of quasars is radio loud and why radio-loudness has a bimodal distribution. In terms of the $R=F_r/F_o$ ratio, where $F_r$ and $F_o$ are the monochromatic fluxes measured at frequencies $\sim 10^{10}$ Hz and $\sim 10^{15}$ Hz, respectively, radio-quiet quasars cluster around $R \sim 0.3$ and radio-loud quasars cluster around $R \sim 300$ (Kellermann et al. 1989; Falcke, Sherwood \& Patnaik 1996). Thus, the average radio-loudness of the two quasar populations differs by a factor $10^3$ and this number, together with typical radio luminosities of radio loud quasars, $L_r \sim 10^{45}$ ergs s$^{-1}$, provides the basic quantitative conditions which should be satisfied by any unified model of quasars. These conditions, together with our results discussed in the two previous sections are used below to test spin based models of a jet activity in quasars. The predictions of such models should also satisfy such observationally established trends, as \noindent - radio-loud quasars avoid disc-galaxies and have UV-luminosities $ \ge 10^{46}$ ergs s$^{-1}$; \noindent - radio-quiet quasars are present both in spiral and elliptical galaxies \cite{tay96} and their radio properties do not depend on the galaxy morphology \cite{kuk98}; \noindent - radio properties of radio-quiet quasars suggest that they are, like in radio loud quasars, related to the jet production by a central engine. Assuming that the efficiency of conversion of jet energy into radio emission is $\sim 10$\%, the typical jet in radio-loud quasars should have $P~\sim~10^{46}$ ergs s$^{-1}$. Similar jet powers are deduced by calculating the total energy content of extended radio sources and dividing it by the age of the source \cite{lea89}, or from energetics of $\gamma$-ray production in sub-parsec jets (see, e.g., Sikora 1997). Largest powers which can be extracted from rotating black holes are given by equation (\ref{power}). For $A=1$ and $B_{\perp} = 8\pi p_{tot}$ we obtain $P_{max} \sim 3 \times 10^{44} M_8^2 p_{tot,8}$ ergs s$^{-1}$, where $p_{tot,8} = p_{tot}/10^8$dyne cm$^{-2}$ and $M_8 = M/10^8 M_{\odot}$. One can see from Figure~\ref{pmax} that for high accretion rates ($\dot m > 0.1$, say) and black hole masses $\sim M=10^9 M_{\odot}$, a pressure $\ge 10^8$ dyne cm$^{-2}$ is provided by $\alpha$-discs with $\alpha \le 0.1$, and by all $\beta$-discs. Thus, the B-Z mechanism is efficient enough to power jets in radio-loud quasars, provided the black hole magnetic field is supported by the total disc pressure. If the latter is not true and, as Ghosh and Abramowicz \shortcite{ga97} argued, the energy density of the black hole magnetic field cannot exceed the energy density of the maximum magnetic field in a disc, then the maximum pressure of the black hole's magnetic field is numerically equal to the total pressure in an $\alpha=1$ disc. In this case black hole masses $\sim 3 \times 10^9 M_{\odot}$ are required in order to get $P \sim 10^{46}$ ergs s$^{-1}$. The case of M87 seems to prove that such black holes are not necessarily exceptional \cite{har94}. However, one should note here, that the question of the diffusion of an external magnetic field into an accretion disc is still open (see, e.g., Wang 1995; Bardou \& Heyvaerts 1996). Assuming, as before, that the fraction of the jet energy converted into radiation is 10\%, and that the bolometric corrections for jet radiation at $\sim 10^{10}$Hz and for accretion disc radiation at $\sim 10^{15}$Hz are of the same order, we obtain that $P/L_d \sim 10 L_r/L_o \sim 10 F_r{\nu}_r/F_o{\nu}_o \sim 10^{-4} R$. Thus, for radio-loud quasars models should predict $P/L_d \sim 0.1$, while radio-quiet quasars should cluster around $P/L_d \sim 10^{-4}$. As is seen from Figures~2 and 3, $P/L_d \sim 0.1$ nicely corresponds to black hole equilibrium spin solutions for all but $\alpha > 0.1$ disc models. One can also check, that there are no equilibrium spin solutions which would correspond to radio-loudness of radio-quiet objects. For them $A < 0.03$ is required, provided that radio luminosity scales linearly with $P$. A population of such low spin black holes can exist only if black holes are born with very low spin and then accrete very little (Moderski, Sikora \& Lasota 1997), or if black hole evolution is determined by multi-accretion events with random angular momenta. As one can deduce from results presented in Figure~7, hundreds of accretion events per object are required in order to have more than $90$ \% of black holes with $A < 0.1$ at any given moment. This is too much to be obtained by accretion events induced by capture of dwarf galaxies, but can be achieved by accretion of molecular clouds. Molecular cloud accretion events were recently proposed by Sanders \shortcite{san98} to explain some properties of Sgr A$^*$ and other AGNs. This scenario is supported by the random orientation of central engines vs. the orientation of galactic discs, as deduced from observations of ``UV'' cones (Wilson \& Tsvetanov 1994; McLeod \& Rieke 1995) and radio axis \cite{ckp98} in Seyfert galaxies. Here we should note, that in our simplified treatment of the multi-accretion scenario (Section~3), we didn't take into account the coupling between the spin of the black hole and the orbital angular momentum of the approaching molecular clouds. Such coupling supposedly leads to random wondering of the black hole spin vector. One can now speculate that changes of orientation of the black hole spin could be interrupted and the black hole could be spun-up to very high spins following a merger process. This process could induce a massive and long lasting accretion event. If during such an event the accretion proceeds from a fixed plane and at least doubles the black hole mass, the black hole spin reaches the equilibrium spin and the object becomes a typical radio-loud quasar \cite{msl97}. Since mergers happen mostly in groups and clusters of galaxies, where the population of galaxies is dominated by ellipticals, this could explain why radio loud quasars avoid spiral galaxies. Observational arguments for such a scenario are exactly the same as those used by Wilson and Colbert \shortcite{wc95}. The only difference is, that they postulated formation of high spin black holes via coalescence of two supermassive black holes, while in our scenario high spins result from an accretion process. Note, however, that a coalescence of two black holes, if it happens, does not have to affect much our scenario. If the coalescence involves two black with very different masses, the final spin will be determined by the accretion process, otherwise both processes lead to similar spins. What are the perspectives for an observational test of the assumption that radio-quiet objects have low spins? A possibility to measure the spin of supermassive black holes is provided by the detailed studies of profiles of the X-ray fluorescent iron line produced in the surface layer of the innermost parts of accretion discs. Such lines are detected in many Seyfert galaxies, which represent the low luminosity branch of radio-quiet quasars. For at least one of such objects the line profile was claimed to be consistent with the kinematics given by the rotation of a disc around a black hole in fast rotation \cite{iwa96}. However, as demonstrated by Reynolds and Begelman \shortcite{rb97}, similar line profiles can be produced around non-rotating black holes, provided that a large part of the line emission comes from below the marginally stable orbit. Therefore, much more detailed theoretical models and sensitive observations are required to get conclusive diagnostics from this type of investigations. The remarkable discovery of relativistic jets in several Galactic X-ray sources (cf. Mirabel \& Rodriguez 1994; Hjellming \& Rupen 1995; Newell, Spencer \& Garrett 1997) suggests that the radio-dichotomy exists for Galactic compact objects as well. As was argued recently by Zhang, Cui \& Chen \shortcite{zcc97}, jet activity in these sources can also be conditioned by the value of the black hole spin. | 98 | 4 | astro-ph9804140_arXiv.txt |
9804 | astro-ph9804154_arXiv.txt | We study the interpretation of the mean surface density of stellar companions as a function of separation (or, equivalently, the two point correlation function of stars) in star-forming regions. First, we consider the form of the functions for various simple stellar distributions (binaries, global density profiles, clusters, and fractals) and the effects of survey boundaries. Following this, we study the dependencies of the separation at which a transition from the binary to the large-scale clustering regime occurs. Larson \shortcite{Larson95} found that the mean surface density of companions follows different power-law functions of separation in the two regimes. He identified the transition separation with the typical Jeans length in the molecular cloud. However, we show that this is valid only for special cases. In general, the transition separation depends on the volume density of stars, the depth of the star-forming region, the volume-filling nature of the stellar distribution, and on the parameters of the binaries. Furthermore, the transition separation evolves with time. We also note that in young star-forming regions, binaries with separations greater than the transition separation may exist, while in older unbound clusters which have expanded significantly, the transition contains a record of the stellar density when the stars formed. We then apply these results to the Taurus-Auriga, Ophiuchus, and Orion Trapezium star-forming regions. We find that while the transition separation in the Taurus-Auriga star-forming region may indicate a typical Jeans length, this is not true of the Orion Trapezium Cluster. We caution against over-interpreting the mean surface density of stellar companions; while Larson showed that Taurus-Auriga is consistent with the stars having a fractal large-scale distribution we show that Taurus-Auriga is also consistent with stars being grouped in non-hierarchical clusters. We also argue that to make a meaningful study of the stellar distribution in a star-forming region requires a relatively complete stellar survey over a large area. Such a survey does not currently exist for Ophiuchus. Finally, we show that there is no evidence for sub-clustering or fractal structure in the stars of the Orion Trapezium Cluster. This is consistent with the fact that, if such structure were present when the stars formed, it would have been erased by the current age of the cluster due to the stellar velocity dispersion. | \label{introduction} Stars generally do not form in isolation. Instead, on small scales, they frequently form as members of bound binary or higher-order multiple systems (e.g. Duquennoy \& Mayor 1991; Mayor et al. 1992; Fischer \& Marcy 1992; Ghez, Neugebauer, \& Matthews 1993; Leinert et al. 1993, Simon et al. 1995), while on larger scales they are often members of associations or clusters of stars (e.g. Gomez et al. 1993; Lada, Strom, \& Myers 1993; Zinnecker, McCaughrean, \& Wilking 1993). Studying the clustering properties of stars on different length scales may help to determine what processes are involved in their formation. Gomez et al. \shortcite{GHKH93} found that the pre-main-sequence stars in the Taurus-Auriga molecular cloud are not randomly distributed, but instead are in small associations of $\sim 15$ stellar systems within radii of $\sim 0.5-1.1$ pc. As one method of analysing the spatial distribution of stars, Gomez et al. determined the two-point angular correlation function and found that it could be represented by a single power-law over separations from $0.005$ to $5$ pc, implying that stars are clustered self-similarly. However, they also found weak evidence that two-point angular correlation function may be better represented by two different power laws with a break at $\approx 0.05$ pc. Using data from searches for binary companions to pre-main-sequence stars in the Taurus-Auriga molecular cloud, Larson \shortcite{Larson95} extended the two-point angular correlation function to smaller separations than Gomez et al. \shortcite{GHKH93} and demonstrated that, indeed, there is a break at $\approx 0.04$ pc. Rather than using the standard two-point angular correlation function, Larson used the closely-related mean surface density of companions (MSDC) (see Section \ref{MCSD}). The MSDC has the advantage that no normalisation is required, whereas the two point correlation function must be normalised by the average density in the survey area which can be difficult to determine if the stars are clustered. Larson \shortcite{Larson95} found that, for stars in the Taurus-Auriga molecular cloud, the MSDC has a power-law slope of $\approx -0.6$ on large scales, but steepens below $\approx 0.04$ pc with a slope of $\approx -2$ on small scales. The fact that a break occurs indicates that a single scale-free process is not responsible for the formation of stars on both scales. The power-law slope of $\approx -0.6$ on large scales is due to the clustering of stellar systems that Gomez et al. \shortcite{GHKH93} studied. Furthermore, Larson pointed out that a power-law slope of $-0.6$ means that the number of stars within an angular distance $\theta$ of an average star increases as $\theta^{1.4}$ and, thus, the distribution of stars on this scale can be described as a fractal point distribution with dimension 1.4. Larson identified the power-law slope of $-2$ for small angular separations with the distribution of binary separations, since stellar pairs closer than $0.04$ pc in Taurus-Auriga are typically mutually bound. However, the power-law slope of $\approx -2$ is not due to a fractal distribution. Rather, it results from the fact that the frequency distribution of binary separations is roughly uniform in log-separation \cite{DuqMay91}. Finally, Larson noted that the length scale of $\approx 0.04$ pc is essentially equal to the typical Jeans length in the Taurus-Auriga molecular cloud. Thus, Larson associated the location of the break in the MSDC with the Jeans length, speculating that companions with separations smaller than this formed due to the fragmentation of a single collapsing molecular cloud core, while on larger scales stars are grouped self-similarly due to hierarchical structure in the progenitor molecular clouds. Following Larson's analysis of the Taurus-Auriga star-forming region (SFR), Simon \shortcite{Simon97} considered the spatial distribution of stars in the Ophiuchus and Orion Trapezium regions. As with Taurus-Auriga, a break was found in the MSDC for each region. On small scales, both Ophiuchus and the Orion Trapezium could be fit by power laws with slopes of $\approx -2$. On large scales, flatter power laws were required of $-0.5\pm0.2$ for Ophiuchus and $-0.2\pm0.2$ for the Orion Trapezium. However, the break between the two regimes was found to occur at $\approx 400$ AU for the Orion Trapezium and $\approx 5000$ AU for Ophiuchus, compared to $\approx 10000$ AU (taking the mean of Simon's and Larson's results) for Taurus-Auriga. Simon concluded that all three SFRs had similar distributions of binary separations and similar fractal structure on large scales, but that the location of the break seemed to depend not only on the Jeans length, but also on the stellar density of the SFR. Finally, Nakajima et al. \shortcite{NTHN98} considered the MSDC of stars in the Orion, Ophiuchus, Chamaeleon, Vela, and Lupus star-forming regions. Again, for those regions where the survey data extends to small enough separations, they find a break in the MSDC with a power-law slope of $\approx -2$ on small scales and power-law slopes ranging from $-0.15$ to $-0.82$ on large scales. The location of the break was also found to vary from a minimum of $\approx 1000$ AU to a maximum of $\approx 30000$ AU. Nakajima et al. also considered the nearest-neighbour distributions for each of the regions and found that when the nearest-neighbour distribution could be fit well by a Poisson distribution, the MSDC had a power-law index close to zero on large-scales, while when the nearest-neighbour distribution was broader than the Poisson distribution, the MSDC had a large, negative power-law index. They interpreted this as evidence that the MSDC may indicate a star formation history in the region rather than the presence of self-similar spatial structure; if the stars have a range of ages, the older stars typically will be more dispersed than the younger stars resulting a spread in the distribution of separations of nearest neighbours and a range of stellar surface density which provides the slope of the large-scale MSDC. Motivated by these papers, we make a careful study of the interpretation of the mean surface density of companions (MSDC) of star-forming regions. Amongst other goals, we wish to determine the relationship of the break between the binary and large-scale regimes to the Jeans length and the stellar density in star-forming regions. We also want to determine how sensitive the MSDC is to detecting sub-structure in a stellar distribution and, when detected, what can be said about the form of the sub-structure (e.g. whether the sub-structure is self-similar or not) and how robust the result is. In Section \ref{MCSD} we consider the calculation of the MSDC function, handling of survey boundaries, and the results for simple stellar distributions (binaries, global density profiles, clusters, and fractals). In Section \ref{posbreak} we derive the dependencies of the break between the binary and large-scale regimes, and show that the separation at which the break occurs can only be identified with the Jeans length in special cases. We also indicate how the MSDC of SFRs is expected to evolve with time. Based on these results, we reconsider the Taurus-Auriga, Ophiuchus, and Orion Trapezium star-forming regions in Section \ref{application}. Finally, we present our conclusions in Section \ref{conclusions}. | \label{conclusions} We have studied the interpretation of the mean surface density of companions (MSDC) as a function of separation $\Sigma_{\rm com}(\theta)$ in star-forming regions. We have shown how the power-law slope of $\approx -2$ for binaries is due their flat distribution of periods in the logarithm of separation, and have considered the MSDC of various global density profiles, sub-clusters and self-similar distributions. We emphasise that simply because a power-law slope can be fit to a particular MSDC, it does not mean that the stellar distribution is self-similar or fractal. We have also demonstrated the effects of survey boundaries on the calculation of $\Sigma_{\rm com}(\theta)$. Several methods of attempting to avoid boundary effects were considered, all of which provide a full correction in the case that there is no large-scale stellar density gradient across boundaries, but none of which give a perfect correction when there are such large-scale gradients. Of these, we recommend Method 5, since it allows the maximum range of separations to be studied, does not discard any information, and is simple to use for surveys with irregular boundaries. Even in the case of a uniform stellar distribution, the improper handing of boundaries results in the $\Sigma_{\rm com}(\theta)$ having a significant slope for separations greater than $\approx 1/50$ of the survey area's dimensions (i.e.~using Method 1). Larson \shortcite{Larson95} associated the separation at which a break in $\Sigma_{\rm com}(\theta)$ between the binary regime and the large-scale regime occurs with the Jeans length in the Taurus-Auriga star-forming region (SFR). However, we show this transition separation may only be associated with the Jeans length in special cases, and that the transition separation does not necessarily give the maximum binary separation. In general, the break occurs at the separation where the mean surface density of {\em binary} companions is equal to the mean surface density of {\em non-binary} companions (the latter of which may be physically close, or simply chance projections). Thus, typically, the break occurs at smaller separations for SFRs with higher stellar surface densities (as observed by Simon \shortcite{Simon97} and Nakajima et al. \shortcite{NTHN98}). In turn, the surface density of non-binary companions depends on the parameters of the binaries, the volume density of stars, the volume-filling factor of the stellar distribution and, in general, the depth of the star-forming region. The transition separation between the binary and the large-scale regimes also evolves with time. Due to a stellar velocity dispersion, initial structure is erased and the surface density of stars in an unbound region generally decreases. This effect begins at the smallest scales, extends to larger scales with time, and results in the transition separation increasing with time. Finally, the transition between the binary and the large-scale regimes may allow a truncation of binaries at large separations to be detected, especially in old clusters that were much denser when the stars were formed and have since expanded. In such cases, this provides a record of the stellar density when the stars first formed. In summary, the transition separation may be associated with the Jeans length only if the star-forming region is young enough that initial structure has not been erased, and if the SFR is `optically thin' in the sense that projection effects due to the depth of the SFR do not affect the transition separation. The latter is true if the volume-filling factor of the SFR is low (e.g. the SFR is composed of widely separated clusters consisting of only a few stars ($\sim 10$), or if the stars have a fractal distribution with dimension $\simless 1.5$). This is the case for the Taurus-Auriga SFR, which explains the good agreement between between the transition separation and the Jeans length found by Larson \shortcite{Larson95}, but it is not the case for the Orion Trapezium Cluster. It is important when studying the large-scale spatial distributions of star-forming regions to obtain the most complete sample of stars over the largest area possible. The lack of such data for the Ophiuchus SFR makes an attempt to study its large-scale spatial distribution of little use at this time. For the Taurus-Auriga and Orion Trapezium SFRs, the current data makes a meaningful study of their large-scale stellar distribution possible. For the Taurus-Auriga SFR, Larson \shortcite{Larson95} fit the large-scale MSDC with a power-law slope that implied a fractal stellar distribution. However, this is not the only possible interpretation; the data can be equally well fit by assuming the stars are formed primarily in randomly-distributed clusters of stars. For the Orion Trapezium SFR, we find that the MSDC is consistent with the stars simply being distributed according to a surface density that decreases with radius; there is no evidence for sub-structure (either fractal or sub-clusters) in the stellar distribution. We also demonstrate how upper limits can be placed on how much sub-clustering is present, and note the the sensitivity of the MSDC to detecting sub-structure appears to be slightly less than that of the human eye. The results for the Orion Trapezium SFR are consistent with the fact that if structure were present when the stars formed, it would have been erased by the current time due to the stellar velocity dispersion. Binaries in the Taurus-Auriga and Orion Trapezium SFR are roughly consistent with an MSDC with a power-law slope of $\approx -2$. However, we point out that comparing power-law indices derived from the slope of the MSDC in the binary regime is not the best way to compare the distribution of binary separations between stellar populations since any structure or deviation from a true power-law may easily be missed. In the centre of the Orion Trapezium SFR, we find very weak evidence that there may be a deficit of binaries with separations $\simgreat 500$\,AU\@. Such a deficit may be caused by the disruption of wide binaries by single-binary star encounters. Finally, in view of our studies of the Taurus-Auriga and Orion Trapezium SFRs, we emphasise caution when interpreting the MSDC. Rather than attempting to characterise star-forming regions simply by fitting power-laws to $\Sigma_{\rm com}(\theta)$, it is more instructive to also consider the global stellar distribution (e.g. a radial surface density profile) and to compare the MSDC to those of model stellar distributions to determine the robustness of any conclusions. Alternatively, rather than just considering the MSDC (or, equivalently, the two-point correlation function), correlation functions of higher order (three and four-point correlation functions) and/or the nearest-neighbour distribution can be used to differentiate between non-hierarchical and hierarchical structure. The use of higher-order correlation functions is common in studying the large-scale structure of the universe \cite{Peebles80}. The nearest-neighbour distribution has been used by Nakajima et al. \shortcite{NTHN98} to argue that the power-law slope of an MSDC on large scales may indicate a stellar age spread rather than the presence of hierarchical structure. However, while an age spread does help explain their results, we argue that their results do not exclude the possibility that stars form in hierarchical structures. More work is required on this topic. | 98 | 4 | astro-ph9804154_arXiv.txt |
9804 | astro-ph9804012_arXiv.txt | The diffuse X-ray emission from the thin disk surrounding the Galactic mid-plane (the so-called Galactic ridge) was measured with {\it RXTE} PCA in order to determine the spatial extent, spectral nature, and origin of the emission. Spatial examination of the diffuse emission in the central $30^\circ$ of the plane in Galactic longitude reveals the presence of two components: a thin disk of full width $ \lesssim 0^\circ \!.5$ centered roughly on the Galactic mid-plane, and a broad component which can be approximated as a Gaussian distribution with FWHM of about $ 4^\circ$. Assuming an average distance of 16~kpc to the edge of the galaxy, a scale height of about $ 70$~pc and 500~pc is derived for the thin and broad disk components, respectively. Spectral examination of the emission clearly reveals the presence of a hard power law tail above 10~keV and an emission line from He-like iron, indicating both thermal and possibly non-thermal origins for the diffuse emission. The averaged spectrum from the ridge in the $3-35$~keV band can be modelled with a Raymond-Smith plasma component of temperature $\sim 2-3$~keV and a power law component of photon index $\sim 1.8$. Based on this finding, we argue that the temperature of the hot phase of the Interstellar Medium (ISM) is less than the previously reported values of $5-15$~keV. Motivated by the similarities between the characteristics of the thermal component of the Galactic ridge emission in our model and the thermal emission from supernova remnants (SNRs), we discuss the origin of the thermal emission in terms of a population of SNRs residing in the Galactic disk. We find that a SN explosion rate of less than 5 per century is adequate to power the thermal emission from the ridge. The origin of the emission in the hard X-ray band modelled by a power law remains uncertain. Possible contributions from non-thermal bremsstrahlung of cosmic ray electrons and protons, inverse Compton scattering of energetic electrons from ambient microwave, infrared, and optical photons, non-thermal emission from SNRs, and emission from discrete X-ray sources are discussed. We speculate that bremsstrahlung of accelerated electrons and protons in SNR sites can play a significant role in producing the hard tail of the spectrum. Moreover, their collisional losses can play a major role in the ionization of the ISM. | A galaxy can be well described by an ecosystem. There is an intimate relationship, much like a symbiosis, between the discrete components of the galaxy such as stars, and its interstellar medium (ISM). The ISM provides the foundation for the birth of a new generation of stars, while at the same time it is enriched by the remains of the older generations and their byproducts during their life cycle. Hence, it is natural to expect that the understanding of formation and evolution of galaxies is closely related to the understanding of their ISM. The ISM of the Milky Way has been found to have many components. Magnetic fields and cosmic ray gas compose the relativistic fluid, while the gaseous phase consists of both ionized and neutral components. The hot ionized medium observable in UV and X-ray has a temperature above $\sim 10^5$~K and is composed of hot coronal gas heated by supernova shocks. A good portion of the energy of the ISM resides in this component. The warm ionized medium (e.g. HII, planetary nebulae) is visible in ${\rm H\alpha}$, UV, and optical, and has a temperature as high as $10^4$~K. The neutral atomic gas appears to have both cold ($<100$~K; e.g. HI~clouds) and warm (100~K~$<T<$~8000~K) components. Molecular clouds compose the self-gravitating gaseous phase. Finally the ISM (with the exception of its hot phase) is filled with dust, particles typically a few tenth of microns in size and visible through their infrared radiation. Since the peak emission of each component arises in different wavebands, multi-wavelength observations from radio to $\gamma$-rays are needed to probe the ISM. Diffuse X-ray emission from our galaxy is a powerful diagnostic of the hot phase of the ISM. It is our purpose here to use this tool to probe the processes that contribute to its energetics and dynamics. The first detection of the X-ray emission from the Galactic disk was achieved by the pioneering rocket experiment of Bleach et al. (1972). They detected excess emission associated with a narrow disk component of angular size $2^\circ-7^\circ$. Since then, X-ray emission from the Galactic plane and in particular the ridge (the narrow region centered on the Galactic mid-plane extending approximately to $\pm 60^0$ in longitude and $\pm 10^0$ in latitude) has been measured in the past with several satellites (e.g. {\sl HEAO-1} [2-50~keV]: Worrall et. al. 1982; {\sl EXOSAT} [2-6~keV]: Warwick et al. 1985; {\sl Tenma} [2-11~keV]: Koyama et al. 1986; {\sl Ginga} [2-16~keV]: Yamasaki et al. 1997; {\sl ASCA} [0.5-10~keV]: Kaneda et al. 1997). The presence of the $6.7$~keV iron line in the spectrum discovered with {\sl Tenma} has motivated the idea that part of the emission below 10~keV is due to a hot optically thin plasma of temperature $5-15$~keV. Because of its high spectral resolution, {\sl ASCA} has provided the most accurate measurement of the spectrum of the emission to date. The presence of Mg, Si, and Fe K-lines in the spectrum obtained by {\sl ASCA} suggests that at least part of the emission is of thermal origin. Close examination of {\sl ASCA} data has also revealed that unresolved, discrete sources are not responsible for the bulk of the emission (Yamauchi et al. 1996; Kaneda 1997). The most recent investigation of the diffuse emission from the Scutum arm region with {\sl ASCA} (Kaneda et al. 1997) has concluded that the emission below 10~keV has both soft ($kT\sim 0.8$~keV) and hard ($kT \sim 7$~keV) thermal components. If indeed the super hot gas ($\sim 7$~keV) exists in an extended form in the ISM, it is not clear how to explain its confinement to the Galactic disk since its temperature exceeds the gravitational potential of the disk by at least an order of magnitude (Townes 1989). Unfortunately, the two-temperature model does not produce a good fit to the data above 10~keV indicating the presence of additional component(s) at higher energies. Indeed, a hard power law tail has been detected in the hard X-ray/soft $\gamma$-ray band from observations of the ridge with {\sl Ginga} and the balloon experiment {\it Welcome-1} (Yamasaki et al. 1997), and {\sl OSSE} (Skibo et al. 1997). In this paper, we present the results from {\sl RXTE} measurement of the diffuse X-ray emission from the Galactic ridge in the 3-35~keV band. Observations in the hard X-ray/soft $\gamma$-ray band have usually been complicated by the presence of numerous variable discrete sources, and the fact that the instruments generally have large fields of view and no imaging capabilities, or have imaging capability but no diffuse emission sensitivity. The combination of these factors makes the separation of emission between diffuse and compact sources a difficult task. The advantage of {\sl RXTE} over previous missions is its small field of view ($1^\circ$ FWHM) combined with its wide energy bandpass (2-60~keV for the PCA), allowing for the subtraction of the contribution of discrete sources from the diffuse emission spectrum in the hard X-ray band. In addition to reporting on the detection of a hard power law tail in the {\sl RXTE} data, we also offer an alternative interpretation for the origin of the emission below 10~keV (i.e. instead of a super hot plasma of temperature $5-15$~keV). In agreement with previous studies, our results indicate that the emission is most likely of diffuse origin as opposed to the superposition of discrete sources. However, we present a model in which the X-ray emission is the superposition of both thermal (modelled by a Raymond-Smith plasma) and possibly non-thermal (modelled by a power law) components. We discuss the origin of the thermal component in terms of a population of SNRs residing in the disk. The origin of the power law component remains uncertain. By comparing the spectrum of the diffuse emission in hard X-rays ({\sl RXTE}) and soft $\gamma$-rays ({\sl OSSE}), we find indications that the emission in the two bands are related. We discuss its origin in terms of both discrete hard X-ray sources and radiation mechanisms such as non-thermal bremsstrahlung from cosmic ray electrons and protons, and inverse Compton scattering. The plan of this paper is as follows. In \S 2, we describe the observations. In \S 3, we present the results of our spatial and spectral analysis of the data. \S 4 is devoted to the discussion of the results, and their implications for the origin of the emission. Finally in \S 5, we present our conclusions. | Our results are summarized as follows: (1)~From the {\sl RXTE} survey of the Galactic plane, we find that the diffuse emission from the Galactic ridge has two spatial components: a thin disk of width $\lesssim 0^\circ\!.5$, and a broad component with a functional form that can be approximated by a Gaussian distribution of about $4^\circ$~FWHM. Assuming an average distance of 16~kpc to the edge of galaxy, this translates to a height of 70~pc and 500~pc for the thin and broad disk, respectively. (2)~A hard power law tail is clearly detected in the spectrum and dominates above 10~keV, implying that the emission in hard X-ray possibly has non-thermal origin. On the other hand, the detection of the emission line from He-like iron in the spectrum (and also lower energy lines from Mg, Si, S in the spectrum from {\sl ASCA}) motivates the idea that part of the emission below 10~keV has thermal origin. (3)~The spectrum in the $3-35$~keV band can be well modelled with a Raymond-Smith plasma component of $2-3$~keV plus a power law component of photon index $\sim 1.8$. The spectrum above 10~keV simultaneously fitted with {\sl OSSE}'s spectrum can be fitted with a power law of photon index $2.3\pm0.2$. The change in the power law slope at lower energies hints at the possibility that the power law either flattens or gradually attenuates below some energy between 10-100~keV. (5)~The characteristics of the thermal component of the diffuse emission resembles that of the SNRs. From this interpretation, we calculate that a SN explosion rate of less than 5 per century is adequate to power the thermal emission from the ridge. (6)~The origin of the emission in the hard X-ray band modelled by a power law is uncertain. While unresolved discrete sources are expected to contribute, the bulk of the emission is expected to be of diffuse origin. One possibility is the non-thermal bremsstrahlung of electrons and protons which may have been accelerated in the SNR sites. Based on the empirical modeling of the data in a wide energy band, we speculate that the power law either flattens or gradually attenuates at lower energies. This lessens the large power injection required to explain the hard power law tail via either electron or proton bremsstrahlung. It will also make it consistent with the power injected to the Galaxy via supernova explosion. (7)~The hydrogen ionization rate implied from our model is $3.3 \times 10^{-15}n_H^{-1}\,{\rm (atom \, s)}^{-1}$ averaged over the central ridge. Generally, values of $10^{-15}-10^{-14}\, {\rm (atom \, s)}^{-1}$ are expected in the ISM. This is an indication that the collisional losses associated with the bremsstrahlung radiation of energetic particles that produce the hard power law tail are sufficient to explain the ionization rates observed in the ISM. Finally, we expect that simultaneous {\sl RXTE}/{\sl OSSE} observations of the diffuse emission from the Galactic ridge will allow the exclusion of hard discrete sources from the spectrum and will tightly constrain the hard tail slope that dominates in the hard X-ray/soft $\gamma$-ray band. This will enhance our understanding of the origin of the power law tail and the energetics of ISM. | 98 | 4 | astro-ph9804012_arXiv.txt |
9804 | astro-ph9804068_arXiv.txt | \footnotemark{} \footnotetext{Table 1 is 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.} Spectroscopic observations of a sample of 73 very luminous IRAS galaxies ($\rm log(L_{IR}/L_{\odot})\ge11.5$ for $\rm H_{0}=50km s^{-1} Mpc^{-1}$,$\rm q_{0}=0.5$) from the 2Jy redshift survey catalogue were carried out using the 2.16m telescope at the Beijing Astronomical Observatory. The observational data, including the optical images (extracted from Digital Sky Survey) and spectra for these galaxies, are presented in paper I (Wu et al., 1998). In this paper, we give the spectral and morphological classifications for these very luminous IRAS galaxies (VLIRGs). We show that about 60\% of VLIRGs exhibit AGN-like spectra (Seyfert 1s, Seyfert 2s, LINER-like galaxies). This fraction goes up to 82\% for the ultraluminous IRAS galaxies (ULIRGs) subsample ($\rm Log(L_{IR}/L_{\odot}) \geq 12.0$). 56\% of the VLIRGs show strong interaction or merging signatures; this fraction rises to 91\% for the ULIRGs. These statistical results strongly suggest that interaction triggers nuclear activities and enhances the infrared luminosity. We find that LINER and a mixture type which have optical properties of both HII galaxies and LINERs could be at the transition stage from infrared luminous HII galaxies to AGNs; their main energy production is from starbursts as well as AGNs. Both infrared luminosities and $\rm H\alpha$ equivalent widths increase dramatically as nuclear separations between VLIRGs and their nearest neighbors decrease. There is little doubt that strong starbursts happen in the nuclei of VLIRGs. Assuming class 0 as advanced merger, we construct a simple merger sequence, from morphological classes 1 to 4 (with near or far companions), to class 5 and 6 (interacting pairs and mergers) and then to class 0 (isolated galaxies). Along this sequence, VLIRGs evolve from HII galaxies to AGNs. | Very luminous infrared galaxies (VLIRGs) , especially ultraluminous infrared galaxies (ULIRGs) have bolometric luminosities comparable with quasars and dominate the top end of the galaxy luminosity function (Soifer, et al., 1987). Many theoretical and observational studies (Sanders et al., 1988; Norris et al., 1988; Struck-Marcel and Tinsley, 1978; Rieke et al., 1980; Norman and Scoville, 1988; Rieke, 1988; Condon et al., 1991; Leitherer and Heckman, 1995) show that the high infrared luminosities are associated with the phenomena of active galactic nuclei (AGNs) and starbursts, rather than the radiation from an old stellar population (Thronson, et al., 1990) and cloud-cloud collision (Harwit, et al., 1987). The main energy output mechanism for luminous infrared galaxies is still being debated. Both AGNs and starbursts have been proposed as the main energy provider. Sanders et al. (1988) found that a high proportion of AGNs seems to be present in ULIRGs, but Rigopoulou, Lawrence and Rowan-Robinson (1996) found that the starburst model fits well the radio, sub-millimeter/IR to soft X-ray observations for Sanders' 10 ULIRGs. Recent 3D infrared spectroscopic imaging observation (Kroker et al., 1996) suggests that both Seyfert 1 nucleus and circumnuclear star formation contribute significantly to the total luminosity. It is possible that AGNs and starbursts co-exist in the central region. Recently, Veilleux, Sanders and Kim (1997) reported the results of a near-infrared search for hidden broad-line regions in a sample of 25 ULIRGs with no broad-line regions in the optical band. They detected broad recombination lines in five objects and all of them are Seyfert 2 galaxies. This observation provides important clue that there could exist evolution connection from nuclear starbursts to AGNs in the central region of luminous infrared galaxies. In order to learn more about the central region of luminous infrared galaxies, and to understand the possible evolutionary process from starbursts to AGNs and how they are associated with galaxy' interaction/merging, we carried out spectroscopic observations for a large sample of 73 VLIRGs compiled from the 2Jy redshift survey catalogue (Strauss et al. 1990, 1992). The spectra and the optical images were presented in a companion paper (Wu et al. 1998, hereafter paper I). In this paper, we report the spectral classification and analysis for the sample (section 2). The morphological classification and their environmental properties are discussed in section 3. In section 4, we point out a possible spectral evolutionary process of luminous infrared galaxies using the relationship between the spectral and morphological classifications. A possible nuclear and evolutionary model for these galaxies is then constructed. Throughout this paper, we use $\rm H_{0}=50km s^{-1} Mpc^{-1}$ and $\rm q_{0}=0.5$. | We have conducted a spectroscopic survey for a complete sample of VLIRGs selected from the 2Jy redshift survey catalogue. We have studied their optical properties and the spectral and morphological classifications for these objects. Important clues are found concerning the nature of VLIRGs and the possible evolution connection between starbursts and AGNs. \begin{enumerate} \bibitem{} Spectral classification shows that over half (60\%,44/73) of VLIRGs present AGN-like spectra. As the infrared luminosity increases, this fraction increases from 56\% ($\rm 11.5 \leq Log(L_{IR}/L_{\odot}) \leq 12.0$) to 82\% ($\rm Log(L_{IR}/L_{\odot}) \geq 12.0$). If we exclude LINER-like galaxies, the fraction of AGNs is 13\%, 36\% for the two subsamples, respectively (see table 2 a-c). \bibitem{} From HII galaxies to AGNs, there is no clear boundary in the diagnostic diagrams (see Figs. 1-3). We classified galaxies with both HII and LINER spectral features as a mixture type. Evidences show that this type of VLIRGs has similar properties as LINERs. We combine the LINERs and the mixture type into a single class as LINER-like galaxies. \bibitem{} Statistics show that there exist correlations among three dust indicators: E(B-V), EW(NaID) and C6563/C4861 in the nuclei of HII galaxies, but not in the nuclei of LINER-like galaxies and Seyfert 2s. This indicates complex nuclear structures of AGNs. The line and continuum emission could come from different nuclear or circumnuclear regions. Seyfert 1's show less reddening. \bibitem{} As separations between VLIRGs and their nearest neighbors decrease, infrared luminosities and $\rm H\alpha$ equivalent widths increase significantly. There exists a character separation, about 10kpc, in ULIRGs, below which the infrared luminosity becomes $\rm \ge 2\times10^{12} L_{\odot}$. The histogram of the relative velocity shows that many of our sample galaxies have even smaller separations, which provides the favorable condition for triggering starbursts, and hence the high infrared luminosities for theses objects. \bibitem{} It is clear from the relationship between spectral and morphological types that the interaction between galaxies triggers massive starbursts as well as nuclear activities. VLIRGs have circumnuclear starburst regions around central AGN. As the spectra evolve from HII-like to AGN-like, the dominant energy output mechanism for VLIRGs changes from starbursts to AGNs. \end{enumerate} | 98 | 4 | astro-ph9804068_arXiv.txt |
9804 | astro-ph9804297_arXiv.txt | We present new determinations of bolometric corrections and effective temperature scales as a function of infrared and optical colors, using a large database of photometric observations of about 6500 Population II giants in Galactic Globular Clusters (GGCs), covering a wide range in metallicity (--2.0$<$[Fe/H]$<$0.0). \par\noindent New relations for BC$_K$ {\it vs} (V--K), (J--K) and BC$_V$ {\it vs} (B--V), (V--I), (V--J), and new calibrations for T$_{eff}$, using both an empirical relation and model atmospheres, are provided. \par\noindent Moreover, an empirical relation to derive the R parameter of the Infrared Flux Method as a function of the stellar temperature is also presented. | A global test of stellar evolutionary models requires a direct comparison between theoretical tracks and observations for stars spanning a wide range in stellar parameters, such as temperature, luminosity and metallicity. In order to achieve these goals at least two fundamental ingredients are needed: \par {\it i)~} a complete and homogeneous database of photometric observations; \par {\it ii)~} a suitable set of transformations between observables and absolute quantities. GGCs are the best empirical laboratory to obtain complete and homogeneous spectrophotometric information on Pop. II stars over a wide range of metallicities. Rewieving the published works on the transformations to the absolute plane (see Sect.3), that is bolometric corrections (BCs) and temperature scales as a function of different colors, it is easy to see that very often these calibrations are not based on a complete and homogeneous set of data, spanning a wide range of stellar parameters. Moreover, {\it ad hoc} correction factors are usually adopted to take into account for example possible systematic differences between different photometric systems and/or different assumptions for the reference solar quantities, for the adopted model atmospheres, and for different laws to extrapolate the data towards the UV/IR ranges. Such a scenario indicates that any calibration in the absolute plane can hardly be fully self--consistent as it always depends on the adopted transformations and, more crucial, the residuals among different scales are very rarely linear with the involved parameters. In order to improve the available determinations of the bolometric corrections and temperature scales, we use here our IR photometric database on GGC stars combined with available optical data from the literature to calibrate new, independent, (hopefully) self--consistent transformations, particularly useful to study the red stellar sequences in Pop II stars. In Sect.2 we present the complete database used in our analysis which includes about 6500 RGB and HB stars in a sample of 10 GGCs observed in both optical and near IR bands. In Sect.3 we derive the transformations from observed magnitudes and colors to absolute quantities, such as bolometric corrections and effective temperatures. All the results are listed in Table 3. In Sect.4 we compare the inferred scales with existent ones and we give a fully empirical calibration of the R parameter of the Infrared Flux Method (IRFM) as a function of the effective temperature. Schematic conclusions are eventually presented in Sect.5. | The main results obtained from our analysis can be summarized as follow: \begin{itemize} \item By exploiting the use of a large photometric database of Pop II stars in GGCs we derived new relations for: BC$_K$ {\it vs} (V--K) and (J--K), BC$_V$ {\it vs} (B--V), (V--I) and (V--J), to infer empirical bolometric corrections from observed colors. \item By making use of both an empirical relation and model atmospheres we calibrated two different scales to infer reliable stellar effective temperatures in the range 3000 -- 7500 K. \item We also calibrated an empirical relation between the R parameter of the IRFM and the stellar temperature. \end{itemize} All these relations are summarized in Table 3, where for a given BC$_K$ the corresponding colors and BC$_V$ in two different metallicity regimes and T$_{eff}$ can be read. These relations should be the most suitable to calibrate the red stellar sequences in the color--magnitude diagrams, like the RGB and AGB in old GGCs. | 98 | 4 | astro-ph9804297_arXiv.txt |
9804 | astro-ph9804318_arXiv.txt | We present a comprehensive investigation of the radio polarization properties within the theory of natural electromagnetic wave modes in pulsar magnetospheres. Taking into account the curvature of the field lines, aberration effects and magnetic sweep-back we use the relativistic dielectric tensor in the low-density cold plasma approximation and derive the following polarization characteristics, which are in full agreement with the observational findings. Specifically, we demonstrate that {\bf 1.} The degree of linear polarization decreases with increasing frequency. {\bf 2.} The degree of circular polarization increases with increasing frequency. {\bf 3.} At high frequencies ($\geq$ a few GHz) the degree of linear polarization is correlated to the spin down luminosity $\dot E$. {\bf 4.} At high frequencies long-period pulsars exhibit weaker linear polarization than their short-period counterparts. {\bf 5.} The difference between the refractive indices of the two natural wave modes decreases with increasing frequency which possibly results into a depolarization via superposition. | \label{intro} The polarization properties of pulsar radio emission are highly diverse. Nearly all possible polarization states have been detected, from pulsar to pulsar, from frequency to frequency, from pulse to pulse and even within one pulse profile. Nevertheless, some general statements can be made from an observational point-of-view, which must be explained by any plausible theoretical model: \begin{itemize} \item Pulsar radio emission is usually highly linearly polarized at low frequencies (in the following ``low frequencies'' defines frequencies up to about one GHz, ``high frequencies'' corresponds to everything well above one GHz). Towards higher frequencies the radiation tends to depolarize (e.g. \cite{M71}, \cite{X96}). \item The observed polarization position angle (hereafter PPA) is a very stable feature which usually does not depend strongly on frequency. Apart from the occurrence of sudden $\sim90^\circ$ jumps, the so-called orthogonal polarization modes (hereafter OPM), the PPA usually follows a characteristic S-shaped curve. This curve can well be modelled by the rotating vector model (hereafter RVM) as proposed by Radhakrishnan \& Cooke (1969). It is therefore thought to reflect the geometry of the radiating field lines. \item The sum of many (typically at least a few hundred) single pulses usually converges towards a characteristic average profile. In contrast, the individual pulses show a high variability. One well known single pulse feature is the {\it quasi-periodic microsecond-structure} (e.g. \cite{CH77}, \cite{L98}) observed over a broad frequency interval with no obvious frequency dependence of the quasi-periodicity. It is usually highly linearly and circularly polarized with a roughly constant PPA and often OPMs at the edges. This suggests that, during one micro-pulse, only one of the orthogonal polarization modes is observed. Another feature are the {\it subpulses}, sometimes drifting in pulse phase from one pulse to the next one. The PPA usually shows a swing during one subpulse, with the swing being stable relative to the subpulse. So when the subpulse drifts, the swing drifts as well, causing a depolarization in the integrated profile (e.g. \cite{MTH75}). \item In average profiles and in individual pulses OPMs can be observed (e.g. Stinebring et al. 1984, Gangadhara 1997). The magnitude of these jumps is usually $90^\circ $ but can also be less sometimes. Sudden orthogonal jumps in the PPA can be found in nearly all pulse phases and even at the same pulse phase in successive single pulses. The OPMs often coincide with changes in the handedness of the circular polarization. Their existence therefore suggests that the radiation preferentially takes two orthogonal elliptical states. \item Some pulsars show a high degree of circular polarization, preferentially near the centroid of the profile. Sometimes it also sometimes shows a sense reversal near the centroid (\cite{RR90}). Recent observations show that a few pulsars have a strong increase in their degree of circular polarization towards high frequencies together with a decrease in linear polarization, thus suggesting that a process might be active which transforms linear into circular polarization (\cite{HKK98}). \item At low frequencies there seems to be no significant relation between the average degree of polarization and any other pulsar parameter. At higher frequencies, however, correlations have been found between the degree of polarization and the spin down luminosity $\dot E$ respectively the surface acceleration potential (\cite{M81}, \cite{X95} and \cite{HKK98}). Pulsars with a high $\dot E$ (respectively a short period and a large period derivative) show a higher average polarization at high frequencies than those with a lower $\dot E$. This suggests that for these pulsars a possible depolarization process becomes important at higher frequencies than for the other pulsars. Closely connected to this is an anti-correlation between the degree of polarization at high frequencies and the period. Only rapidly rotating pulsars seem to have highly polarized radio emission (\cite{HKK98}). \item Several authors estimated the emission altitudes where the observed radiation is produced and studied if there is a frequency dependence of this altitude (Cordes 1978, Rankin 1990, Thorsett 1991, Blaskiewicz et al. 1991, Gil 1991, Phillips 1992, Gil \& Kijak 1997, von Hoensbroech \& Xilouris 1997, Kramer et al. 1997). Various independent methods were used and all authors came to the conclusion that the radio emission originates from a region close to the pulsar surface at a few percent of the light cylinder radius (at least for ``normal'' -- non millisecond -- pulsars). The frequency dependence of the emission height is thought to be small, if it exists at all. \end{itemize} All these observational facts have to be explained within a general model for the polarization of pulsar radio emission. Basically, there are two aspects which must be considered. One is the polarization characteristic of the emission process itself, and the other one is the influence of the magnetosphere on the propagating radiation. In this paper we concentrate on the possible role of propagation for the polarization because this can be considered independently of the possible radiation mechanisms. The magnetosphere of a pulsar is filled with a plasma which is streaming relativistically outwards along the open field lines. A number of articles has already been published on propagational effects in this plasma. Cocke \& Pacholczyk (1976) considered the effect of Faraday pulsation for quasi-transverse propagation. Faraday pulsation is the general case of Faraday rotation. It occurs when a polarized beam decays into two elliptical orthogonal modes which propagate with different phase velocities. The PPA of the resulting beam gets rotated and the polarization changes between linear and circular. \noindent In a series of papers Melrose \& Stoneham (1977), Melrose (1979) and Allen \& Melrose (1982) derived an approximation for the dispersion tensor and calculated the properties of the natural wave modes in the plasma (see also Lyutikov 1998). They assumed that the polarization of the propagating electromagnetic wave follows the shape of the natural wave modes up to a certain radius of limiting polarization (hereafter $R_{\rm LP}$). The radiation then escapes with the polarization properties of the natural wave modes at $R_{\rm LP}$. Barnard (1986) places $R_{\rm LP}$ at the place of cyclotron resonance. For normal pulsars this resonance occurs at a few 10\% of the light cylinder radius, but for short period pulsars ($P\leq 0.06 s$) this resonance is outside $R_{\rm LC}$. \noindent Cheng \& Ruderman (1979) considered adiabatic walking to be responsible for the variable polarization in subpulses. The polarization properties were derived for different emission mechanisms and plasma compositions. \noindent Harding \& Tademaru (1979, 1980, 1981) presented numerical calculations on the propagation of linearly polarized pulses through a shearing plasma. They show that this can account for micro-structure, circular polarization and rotation of the PPA. \noindent Barnard \& Arons (1986) and Arons \& Barnard (1986) have calculated the dispersion relations for the X- and the O-mode in an ultra-relativistic, one-dimensional plasma for distribution functions in $e^-$ and $e^+$ (the O-mode has its electric field vector in the plane of curvature of the magnetic field, the X-mode perpendicular to it). The X-mode is purely linearly polarized and propagates easily through the magnetosphere whereas the O-mode can have some circular contributions and follows the bending of the field lines. This causes a separation of the OPMs (\cite{McK97}). \noindent Calculations concerning the propagation characteristics of the various natural modes have also been presented by Beskin et al. (1993). Using the limit of infinite magnetic field strength they conclude that only the electromagnetic X-mode can propagate freely through the magnetosphere. The angular dependence of the properties of the modes and the influence of the cyclotron resonance are derived. This resonance becomes only significant, if the secondary particle production rate is high ($\ge 10^4$ per primary particle) thus leading to a much higher plasma density. The different mode properties also account for the circular polarization often observed in core components. In this paper we extended the ideas proposed by Melrose \& Stoneham(1977), Melrose (1979) and Allen \& Melrose (1982). We derive the properties of the natural wave modes throughout the magnetosphere with special respect to the angle between the propagating wave and the magnetic field. We make predictions for the qualitative dependence of the polarimetric properties with frequency and various pulsar parameters. These predictions are then compared with the observations. | We have used the approximation of the dielectric tensor for the low frequency limit to derive the general properties of the natural wave modes in the radio pulsar magnetosphere. The approximation of the dielectric tensor in the given limit implies that the plasma is sufficiently cold. Following Melrose (1979) this should be true for pulsar magnetospheres. Electromagnetic waves, which propagate through a plasma at an oblique angle to the magnetic field, adopt two orthogonal natural wave modes with individual indices of refraction. We identify these modes with the familiar orthogonal polarization modes. It was our aim to derive the polarimetric properties of these modes and calculate their dependence on different parameters under special consideration of the various angles involved. We regard only one natural mode assuming that we see only one mode at a time. This assumption is of course not always true, but the existence of highly polarized emission in single pulses indicates that at least sometimes it is true. If both modes are recorded simultaneously, it will lead to a depolarization of the observed radiation. Our results can therefore be regarded as upper limits for the observed degree of polarization. In the case of a highly linearly polarized natural mode we cannot predict the degree of linear polarization one will observe. But if the mode has only a low level of linear polarization, a low level for the observed radiation is predicted. The main uncertainty of our calculations is the radius of limiting polarization ($R_{\rm LP}$). However, as we are mainly interested in the qualitative behaviour of the polarization, we have used a fixed value for $R_{\rm LP}$ without restriction of generality. For a different value of $R_{\rm LP}$ the numbers may change but the general behaviour remains. This is even true if one moves $R_{\rm LP}$ to its possible extremes: the light cylinder radius and just above the emission region. Our results are in excellent agreement with the observations: {\bf 1.} The degree of linear polarization decreases towards high frequencies (Fig. \ref{freq1}). The decrease in linear polarization implies an increase in circular polarization (within this concept). This is in agreement with recent observations which show the existence of some pulsars which have this peculiar polarimetric property (\cite{HKK98}). {\bf 2.} Given a set of specific pulsar parameters, a variation of $\dot P$ leads to a variation in the spin down luminosity $\dot E$. Figure \ref{edot} shows that the degree of linear polarization of pulsars with a high $\dot E$ decreases at higher frequencies than for pulsars with a lower $\dot E$. Thus a correlation between the polarization at high frequencies and $\dot E$ should exist. Such a correlation has been indeed observed. {\bf 3.} Similar to the above point, we have varied the period, keeping all other pulsar parameters fixed (Fig. \ref{period}). The degree of linear polarization for long period pulsars is expected to decrease at lower frequencies than for short periods pulsars. This should give an inverse correlation between the degree of linear polarization at high frequencies and the period. Again, such an inverse correlation has been observed at a frequency of 5 GHz. No investigation could be made for the change of the polarimetric characteristics along the different field lines, which are observed during a pulse. The reason is that the qualitative change depends heavily on $R_{\rm LP}$. For future work it therefore seems to be important to find an estimate for this distance. The difference between the refractive indices of the two modes decreases with frequency. As this implies a closer propagation of the modes, we expect an increasing superposition. This would lead to a depolarization of the observed radiation at high frequencies. Such a depolarization could account for low abundance of pulsars which show an increasing degree of circular polarization with frequencies. We note again that our calculations are based on the two assumptions that the background plasma is cold and is dominated by particles with one sign of charge. The obvious advantage of these assumptions -- negligible spread of the energy distribution and no strong pair production -- is their physical and mathematical simplicity. The more it is astounding that we can qualitatively reproduce the complex properties of pulsar radio polarization. In general both assumptions are not widely accepted. Either the energy distribution of the plasma is usually taken to be extended (Arons \& Barnard 1986; Lyutikov 1998 and references therein) or the secondarily produced electron and positron pairs are expected to dominate the particle content of the pulsar magnetosphere (e.g. Sturrock 1971, Cheng \& Ruderman 1977 and Daugherty \& Harding 1982). However the expressiveness of our model deserves further considerations. | 98 | 4 | astro-ph9804318_arXiv.txt |
9804 | astro-ph9804182_arXiv.txt | We numerically investigate chemodynamical evolution of interstellar medium (ISM) in gas-rich disk-disk galaxy mergers in order to explore the origin of fundamental chemical properties of halo ISM observed in elliptical galaxies. Main results obtained in this chemodynamical study are the following three. (1) Elliptical galaxies formed by gas-rich mergers show steep negative metallicity gradients of ISM especially in the outer part of galaxies. The essential reason for this is that chemical evolution of ISM in mergers proceeds in such an inhomogeneous way that in the central part of mergers, metal-enrichment of ISM is more efficient owing to radial inflow of metal-enriched ISM during dissipative galaxy merging, whereas in the outer part, metal-enrichment is less efficient owing to a larger amount of metal-enriched ISM tidally stripped away from mergers. This result provides a clue to the origin of gaseous metallicity gradients of elliptical halo recently revealed by $ASCA$. (2) Because of inhomogeneous chemical evolution of ISM in mergers, $some$ merger remnants show mean gaseous metallicity discernibly smaller than mean stellar one. The degree of difference in mean stellar and gaseous metallicity in a merger remnant depends on chemical mixing length, galactic mass, and the effectiveness of supernova feedback. (3) Elliptical galaxies formed by multiple mergers are more likely to have metal-poor gaseous halo components and steep gaseous metallicity gradients than those formed by pair mergers. This is principally because a larger amount of less-metal enriched ISM can be tidally stripped away more efficiently from galaxies in multiple mergers. These three results demonstrate that dynamical evolution of gas-rich galaxy mergers, in particular, tidal stripping of less metal-enriched ISM during galaxy merging, greatly determines chemical evolution of ISM of galaxy mergers. These results furthermore imply that recent $ASCA$ observational results concerning mean and radial chemical properties of halo ISM in elliptical galaxies can be understood in terms of chemodynamical evolution of gas-rich galaxy mergers. | Recent observational studies by $ASCA$ ($Advanced$ $Satellite$ $for$ $Cosmology$ $and$ $Astrophysics$) have revealed a number of fundamental chemical properties of interstellar medium (ISM) of elliptical galaxies, thus have provided valuable information on the formation and evolution of elliptical galaxies (e.g., Awaki et al. 1994; Loewenstein et al. 1994; Matsushita et al. 1994; Mushotzky et al. 1994; Matsumoto et al. 1997; Matsushita et al. 1997). For example, Fe abundance of hot gaseous X-ray halo has been revealed to be appreciably smaller than that of the stellar component in the host elliptical galaxy (Awaki et al. 1994; Matsushita et al. 1994; Matsumoto et al. 1997; but see Matsushita et al. 1997). This smaller gaseous metallicity is considered to be largely inconsistent with the theoretical prediction of the conventionally used one-zone chemical evolution model (`` The iron abundance discrepancy problem'') and thus has been extensively discussed in theoretical studies (e.g., Renzini et al. 1993; Fujita, Fukumoto, \& Okoshi 1996; Arimoto et al. 1997). Furthermore, the hot $X$-ray gaseous halo in elliptical galaxies has been revealed to show strong negative metallicity gradients, which suggests that some physical mechanisms such as cooling flow, gaseous dissipation, galaxy merging, and dilution from external metal-poor gas play a vital role in the formation of the gaseous metallicity gradients (Loewenstein et al. 1994; Mushotzky et al. 1994; Matsushita et al. 1997). Although these $ASCA$ observational results on mean and radial chemical properties of ISM can provide some diagnostics for any theories of elliptical galaxy formation, only a few theoretical studies have addressed the origin of the above fundamental characteristics of chemical properties of ISM in elliptical galaxies. The purpose of this paper is to explore the origin of fundamental chemical properties of ISM of elliptical galaxies recently revealed by $ASCA$. We adopt the assumption that elliptical galaxies are formed by gas-rich disk-disk galaxy mergers and thereby investigate whether or not the merger model of elliptical galaxy formation can give a plausible explanation for the origin of recent observational results of $ASCA$ on mean and radial chemical properties of ISM in elliptical galaxies. We particularly investigate how the dynamical mixing of chemical components during galaxy merging, which has not been investigated at all in previous studies, affects mean and radial chemical properties of ISM of merger remnants. In the present study, the key physical process associated with the origin of metal poor gaseous halo and radial gaseous metallicity gradients in elliptical galaxies is demonstrated to be tidal stripping of less metal-enriched ISM during galaxy merging. Based on the present numerical results, we point out the disadvantages of the commonly used one-zone models in discussing the metallicity of ISM of elliptical galaxies and stress the importance of dynamical processes of galaxy formation in the chemical enrichment processes of ISM in galaxies. The layout of this paper is as follows. In \S 2, we summarize numerical models used in the present study. In \S 3, we demonstrate how a number of fundamental $chemical$ properties of ISM in merger remnants are affected by purely $dynamical$ processes of galaxy merging. In \S 4, we provide a number of implications on chemical properties of ISM in elliptical galaxies. The conclusions of the present study are given in \S 5. | There are a growing number of observational evidences which suggest strong radial negative gradients of hot X-ray gaseous halo in elliptical galaxies (Loewenstein et al. 1994; Mushotzky et al. 1994; Matsushita et al. 1997). Observational study of NGC 4636 (Matsushita et al. 1997) have revealed a factor of $3\sim4$ difference of ISM metallicity within $\sim 7.1 R_{\rm eff}$. Cooling flows, stellar metallicity gradients actually existing in ellipticals, the long-term chemical evolution of ISM driven by stellar mass loss or SNIa, and dilution of metal-enriched ISM by external metal-poor gas are considered to be likely explanations for the origin of gaseous metallicity gradients (e.g., Loewenstein et al. 1994; Mushotzky et al. 1994; Matsushita et al. 1997). These likely explanations are closely associated either with external metal-poor gas or with the later chemical evolution of ellipticals. The present study provides an alternative explanation for the origin: The ISM metallicity gradients can be closely associated with intrinsic and dynamical processes of dissipative galaxy merging at the epoch of elliptical formation. Inhomogeneous and radial-dependent chemical evolution of galaxies is found to play decisive roles in the formation of negative gaseous metallicity gradients of merger remnants. The present numerical results thus imply that the present-day metallicity gradients of hot ISM halo can contain a fossil record of the past dynamical evolution of ISM of elliptical galaxies at the epoch of their formation. Negative metallicity gradients derived in the present study are only true for ellipticals with a few Gyr age, primarily because we only solved a few Gyr evolution of chemical properties of ISM but did not solve the later long-term (corresponding to the Hubble time) evolution. Thus, we should investigate the following two points in our future studies in order to confirm whether or not the proposed explanation for the origin of observed metallicity gradients is actually viable for the present-day ellipticals with their ages of $\sim$ 10 Gyr in a more quantitative sense. The first is long-term chemical evolution of ISM surrounding merger remnants, which can be greatly affected by the later and continuous metal-enrichment due to SNIa and stellar mass loss of long-lived stars. As has been demonstrated by several previous studies on long-term hydrodynamical evolution of hot ISM in ellipticals (e.g., David, Forman, \& Jones 1991), the effectiveness of thermal heating by SNIa (and partly by long-lived stars) determines time-evolution radial flow patters of hot ISM (e.g., either outflow due to effective thermal heating of SNIa or inflow due to efficient cooling). Such later gaseous inflow or outflow, which can transfer metal-enriched ISM, can change radial metallicity distribution that is initially formed by gas-rich galaxy merging. Accordingly, chemodynamical effects of the later stellar mass-loss and supernovae on the radial metallicity distribution should be explored more in detail in our future studies. The second is the long-term dynamical evolution of less metal-enriched ISM that is tidally stripped and surrounds merger remnants. The present study predicts that the tidally stripped metal-enriched ISM surrounds the considerably outer part of merger remnants, where external tidal field resulting from neighbor galaxies and large-scale gravitational potential of cluster or group of galaxies can affect dynamical evolution of the stripped ISM. Later dynamical effects of external tidal force can change drastically initial radial gradients of ISM in merger remnants, thus we should also consider these in our future studies. Since radial metallicity gradients of ISM in ellipticals contain valuable information not only on chemical evolution of ISM but also on dynamical evolution of galaxies as a whole, more extensive studies are necessary for the better understanding of the origin of the gradients. Furthermore, the difference in mean metallicity between stellar and gaseous components derived in the present merger model can be compared with recent observational results by $ASCA$ (Awaki et al. 1994; Matsushita et al. 1994; Matsumoto et al. 1997) which reveals that the abundance of hot gaseous X-ray halo (Fe) is appreciably smaller than that of the stellar component in the host elliptical galaxy. These observational results seem to be consistent with results derived in some merger models of the present study, which imply that metal-poor gaseous components in some elliptical haloes can be formed by gas-rich mergers. The most recent result of Matsushita et al. (1997), however, shows that there is not so large difference between mean stellar metallicity (0.74 solar) and mean gaseous one ($\sim$ 1.0 solar) within $\sim 4.0 R_{\rm eff}$ of NGC 4636 and furthermore that the gaseous metallicity in the outer part of the halo (for the region $4.7 \leq R/R_{\rm eff} \leq 7.1$) is still smaller ($\sim 0.37$ solar) than mean stellar metallicity of NGC 4636. Although metallicity averaged out for the whole region of the gaseous halo in NGC 4636 has not been clarified yet, this new observational result provides the following implication on inhomogeneous chemical mixing derived in the present study. If the gaseous metallicity averaged out for the whole region of the gaseous halo is really larger than stellar one in NGC 4636, the present numerical result that mean gaseous metallicity can be smaller than mean stellar one in some merger remnants is not consistent with the observational result. In this case, we should consider either that inhomogeneous chemical mixing is not so efficient in real galaxy mergers as the present study predicts, or that the later long-term metal-enrichment of ISM resulting from metal-ejection from long-lived stars and supernovae (SNIa) can greatly affect the chemical evolution of ISM after galaxy merging and thus change the difference of mean stellar and gaseous metallicity in merger remnants. Alternatively, if the gaseous metallicity averaged out for the whole region of the gaseous halo is really smaller than stellar one, the present study can provide a clue to the origin of such smaller gaseous metallicity of elliptical haloes: Origin of metal-poor gaseous halo in ellipticals can be closely associated with the past dissipative galaxy merging processes. Since total number of sample X-ray gaseous halo with the mean metallicity estimated precisely is still small, it is safe for us to say, at least now, that future more extensive observational studies and more elaborated theoretical models will assess the validity of inhomogeneous chemical mixing derived in the present study. Main results obtained in the present chemodynamical study are the following three. (1) Elliptical galaxies formed by gas-rich mergers show steep negative metallicity gradients of ISM especially in the outer part of galaxies. The essential reason for this is that chemical evolution of ISM in mergers proceeds in such an inhomogeneous way that in the central part of mergers, metal-enrichment of ISM is more efficient owing to radial inflow of metal-enriched ISM during dissipative galaxy merging, whereas in the outer part, metal-enrichment is less efficient owing to a larger amount of metal-enriched ISM tidally stripped away from mergers. This result provides a clue to the origin of gaseous metallicity gradients of elliptical halo recently revealed by $ASCA$. (2) Because of inhomogeneous chemical evolution of ISM in mergers, $some$ merger remnants show mean gaseous metallicity discernibly smaller than mean stellar one. The degree of difference in mean stellar and gaseous metallicity in a merger remnant depends on chemical mixing length, galactic mass, and the effectiveness of supernova feedback. (3) Elliptical galaxies formed by multiple mergers are more likely to have metal-poor gaseous halo components and steep gaseous metallicity gradients than those formed by pair mergers. This is principally because a larger amount of less-metal enriched ISM can be tidally stripped away more efficiently from galaxies in multiple mergers. These three results demonstrate that dynamical evolution of galaxy mergers can greatly affect chemical evolution of ISM of galaxies, which cannot be attained untill both dynamical and chemical evolution of galaxies are solved in an admittedly self-consistent manner. In particular, tidal stripping of less metal-enriched ISM during dissipative galaxy merging is found to play a vital role in determining mean and radial chemical properties of ISM in elliptical galaxies. The present study accordingly implies that the origins of metal-poor gaseous halo and negative metallicity gradients of ISM in an elliptical galaxy can be closely associated with gas-rich galaxy merging at the epoch of elliptical galaxy formation. | 98 | 4 | astro-ph9804182_arXiv.txt |
9804 | astro-ph9804307_arXiv.txt | We investigate the spectroscopic characteristics of the optical components of Be/X-ray binary systems, using data collected during our seven-year monitoring campaign. We find examples of major changes in the emission line profiles associated with Type II X-ray outbursts, later developing into V/R variability cycles. We show that the time-scales for V/R variability in Be/X-ray transients extend from a few weeks to years and interpret all these changes as due to the presence of global disruptions of the axisymmetric density distribution in the extended envelopes of the Be stars in these systems. The association between X-ray outbursts and V/R variability, the occurrence of very fast changes and the very short quasi-periods of variability displayed by Be/X-ray binaries lead us to conclude that the presence of the neutron star is an important factor affecting the dynamics of the disc-like envelopes. The interaction between the compact companion and the disc would explain the correlation between H$\alpha$ strength and orbital period recently found. The characteristics of the V/R cycles are, however, mainly independent of the binary parameters. | Be/X-ray binaries constitute the major subclass of massive X-ray binaries, in which X-ray emission is due to accretion of matter from an early-type mass-losing star by a compact companion (see Apparao 1994, White et al. 1995, for reviews). Be stars are early-type non-supergiant stars, which at some time have shown emission in the Balmer lines. Both the emission lines and the characteristic strong infrared excess when compared to normal stars of the same spectral types are attributed to the presence of a cool circumstellar envelope, presumably in the shape of a disc (see Slettebak 1988). The physical reasons which give rise to the disc are unknown, but it is generally believed that the high rotational velocity of Be stars plays an important role, even though it is accepted that some other mechanism(s) must be at work. Most Be/X-ray binaries have relatively eccentric orbits and the compact companion (in general, a neutron star, but in some cases possibly a white dwarf) spends most of its time far away from the disc surrounding the Be star. Three kinds of X-ray activity are observed (Stella et al. 1986, henceforth SWR): \begin{enumerate} \renewcommand{\theenumi}{(\arabic{enumi})} \item Persistent low-luminosity ($L_{{\mathrm x}} \la 10^{36}$ erg s$^{-1}$) X-ray emission. Some sources (e.g., X Persei) have always been observed in this state. \item Periodical (Type I in SWR) X-ray outbursts ($L_{{\mathrm x}} \approx 10^{36} - 10^{37}$ erg s$^{-1}$), coinciding with the periastron passage of the neutron star. Type I outbursts have been observed in numerous sources, such as A\,0535+262 (Motch et al. 1991) and EXO\,2030+375 (Norton et al. 1994). \item Giant (Type II in SWR) X-ray outbursts ($L_{{\mathrm x}} \ga 10^{37}$ erg s$^{-1}$), which do not show any orbital modulation. Type II outbursts are normally seen in those sources that also display Type I activity (see Parmar et al. 1989, Finger et al. 1996a for examples). \end{enumerate} Be/X-ray binary systems which display outbursts are collectively termed Be/X-ray transients. Most transients (e.g. A\,0535+26) also show low-luminosity X-ray emission when they are not in outburst, but in systems containing fast-rotating neutron stars, centrifugal inhibition of accretion prevents X-ray emission (SWR) except during outbursts (e.g., 4U\,0115+634). Type I outbursts occur in series between long periods of X-ray inactivity (or low-luminosity emission), while the onset of Type II outbursts is completely unpredictable. | We have presented observational evidence showing that global disruptions are frequent in the extended circumstellar envelopes of Be/X-ray binaries. These perturbations are reflected in the asymmetric line profiles normally observed from these systems. V/R ratio variability is observed to occur with typical time-scales ranging from a few days to several years. In at least two cases (the giant outbursts of A\,0535+26 in February 1994 and 4U\,0115+634 in December 1995), a major disruption seems to have originated in coincidence with the X-ray outburst. Further evidence of the association between fast changes in the line profiles and X-ray outbursts has been seen in most Be/X-ray transients. We believe that all these observation suggest that the presence of the neutron star represents a major factor controlling the dynamics of the discs around the Be stars in X-ray binaries. This fact provides an explanation to the correlation between maximum H$\alpha$ EW and orbital period found by Reig et al. (\cite{reigb}). The frequent presence of major density perturbations in the envelopes of Be/X-ray binaries introduces a new element of complication in the modelling of these systems. Rather than assuming that the disc is static and homogeneous, new models should take into account the presence of global density waves and explore the possibility that the series of Type I outbursts are caused by the interaction of the neutron star with the regions of enhanced density which these waves generate. Continued monitoring of Be/X-ray transients and careful optical coverage of future Type II outbursts will provide the only test for these hypothesis. | 98 | 4 | astro-ph9804307_arXiv.txt |
9804 | astro-ph9804077_arXiv.txt | We have observed with the IRAM interferometer at two different epochs and simultaneously the two transitions $v=0, J=2 \rightarrow 1$ and $v=1, J=2 \rightarrow 1$ of \dSiO and \uSiO in Orion IRc2. We have made the first maps of \dSiO $v=0, J=2 \rightarrow 1$ emission from Orion. These maps and properties of the \dSiO spectra attest to maser emission. Our \uSiO maps show the stable ring of maser spots observed in previous works. Combining our own data with published works we derive that the relative motion between the two ridges of the \uSiO emission ring is less than about 0.7 AU/yr over a period of 7 years. On the other hand, the weak high velocity maser features observed around 30 \kms ~seem to move with respect to the stable ring of \uSiO main emission. Our relative \dSiO ($v=0$) and \uSiO ($v=1$) spot maps show that most \dSiO and \uSiO emission features are closely related but have not the same spatial extent. We conclude that these masers are not excited in the same gas layers in agreement with pumping models which predict that various $v$ state masers peak in different spatial regions. In addition, our maps of $v=0$ and $v=1$ emission suggest that local line overlaps due to turbulence and high gas temperature do not play a dominant role in the excitation of \uSiO and \dSiO, although excitation effects resulting from the overlap of Doppler-shifted ro-vibrational lines may still be significant. | Since the discovery nearly 25 years ago of the SiO molecule in Orion (Snyder \& Buhl 1974), SiO masers have been observed in the envelopes of hundreds of late-type stars, and in the direction of a few galactic HII regions. Orion remains a unique source of SiO emission because it contains all known SiO isotopic species and because it is the only star-forming region exhibiting a very strong maser in the $v=1, J=2 \rightarrow 1$ transition. The strong and compact $v=1$ SiO maser is associated with the luminous infrared source IRc2 and is closely connected with the extended and weaker $v=0$ maser emission mapped at 43 GHz (Chandler \& De Pree 1995) and 86 GHz (Wright et al. 1995). Recent high resolution observations showed that IRc2 is a complex object resolved into four components (Dougados et al. 1993) and that the center of the strong ($v=1$) and weak ($v=0$) SiO maser outflows coincide with the radio continuum source I (Menten \& Reid 1995, Wright et al. 1995), and is displaced from the center of the molecular hot core (lying to the east of source I). On the other hand, the center of the large-scale high velocity bipolar outflow traced by CO lies roughly $3"$ to the north of source I. Considerable efforts have been made to properly model the SiO maser phenomenon in late-type stars. Models include radiative and/or collisional excitation schemes (e.g. Kwan \& Scoville 1974, Elitzur 1980, Langer \& Watson 1984, Lockett \& Elitzur 1992 or Bujarrabal 1994). All models share two general characteristics: $(i)$ they require high volumic densities of order $10^{8}-10^{10}$ cm$^{-3}$; $(ii)$ inversion of the SiO level populations depends on the column density, and maser emission in higher vibrational states peaks at higher values of the column density. On the other hand, there are major observational facts that cannot be explained by any of the present radiative/collisional SiO pumping schemes. In particular, the "standard" pumping schemes fail to explain in stars the peculiar distribution of line intensities within a given vibrational state (e.g. Cernicharo \& Bujarrabal 1993), and fail to explain the absence or weakness of $v=2, J=2 \rightarrow 1$ emission from Orion and late-type stars (Olofsson et al. 1981 b, Bujarrabal et al. 1996). In fact, line overlaps among transitions of the isotopic species of silicon monoxide are an important addition to radiative/collisional pumping in stars (e.g. Gonz\'alez-Alfonso \& Cernicharo 1997), while the line overlap between two near infrared lines of SiO and water explains the weakness of the $v=2, J=2 \rightarrow 1$ transition (Bujarrabal et al. 1996). Depending on the relative importance of property $(ii)$ above (namely various $v$ state masers should peak in different spatial regions) with respect to line overlap effects among nearby transitions of SiO and isotopes the spatial distribution of various vibrational transitions should differ or not. Therefore, high spatial resolution and sensitive maps of SiO and isotopes should provide a test of these predictions. With this idea in mind we have compared the spatial distributions of two nearby transitions of \dSiO and \uSiO toward Orion IRc2 which contains the strongest SiO source in the sky. In Sect. 2 we present our observations and give details of data reduction. In Sect. 3 we discuss spectral variability and present our maps of \dSiO and \uSiO emission from Orion. In Sect.~4 we discuss some properties of the apparent ring of \dSiO and \uSiO masers, the relative spatial extents of both species and implications on their excitation. Some conclusions are summarized in Sect. 5. | \subsection{Stability of the ring of \uSiO masers} The regular position$-$velocity pattern of $v=1, J=2 \rightarrow 1$ emission (Plambeck et al. 1990, Wright et al. 1995) was interpreted by Plambeck et al. as a collection of maser clumps lying in an expanding and rotating disk. This pattern does not change much with time and has been observed with the IRAM interferometer in 1990 (Guilloteau et al. 1992), 1992 (Baudry et al. 1995), and 1995 and 1996 (this work). A similar position$-$velocity pattern was also observed in the $J=1 \rightarrow 0$ transition of SiO around 43 GHz by Morita \etal (1992) and Menten \& Reid (1995). We can estimate the long-term stability of the 86 GHz pattern by measuring the mean separation between the two ridges delineated by the dominant positive and negative velocity features of Orion IRc2. To this end, we have used the 6 different maps made at 86 GHz with the BIMA and IRAM interferometers (Plambeck et al. 1990, Wright et al. 1995, Guilloteau et al. 1992, Baudry et al. 1995, and this work). The general orientation and the mean separation between the two ridges of main SiO features do not seem to evolve with time. The mean separation between these two ridges is $\approx 0.^{''}165 \pm 0.^{''}01$; for the uncertainty we have assumed that the 6 independent measurements behaved as gaussian variables. Therefore, any apparent contraction or expansion of the ring, would be less than or of order $0.^{''}01 / 7$yr, namely $\leq 0.7$ AU/yr at the 480 pc distance of Orion A. On the other hand, stability of the intermediate velocity pattern ($\approx 0 \rightarrow 11$ km s$^{-1}$) is not obvious when we compare our 1990 data with the present IRAM maps. The complex shape observed in 1995 or 1996 is not quite similar to that in 1990 (see Fig. 11 in Guilloteau et al. 1992). Such differences cannot be due to relative position errors which are less than about 2 to 5 mas in the \uSiO maps. These discrepancies seem to agree with the model of Plambeck et al. (1990) which predicts that intensity changes in the intermediate velocity features could cause large position changes of the maser spots in the disk. \subsection{Nature of the \uSiO high velocity features} The high velocity features lying around $28\rightarrow 31$ km s$^{-1}$ are located close to the positive velocities of the SiO ring (see e.g. the 29.2 and 29.1 km s$^{-1}$ features in Figs. \ref{relat28map} and \ref{relat2829map}). Our observations of 1995 August show that these features are excited in an area similar, although not identical, to that observed in 1995 January by Wright et al. (1995, see their Fig. 1c) for their $30 \rightarrow 33$ km s$^{-1}$ features. However, 6.5 months later our data show that the $28 \rightarrow 31$ km s$^{-1}$ features have migrated toward the most positive velocity end of the main SiO emission ridge (see location of the 29.1 km s$^{-1}$ component in Fig.\ref{relat2829map}). The apparent migration of weak high velocity components is consistent with the model of Plambeck \etal (1990) where small changes in brightness distribution of extended features may look like rapid motion. Nevertheless, such rapid motions should be confirmed in future maps of \uSiO emission. These high velocity features could be related to the spectral changes observed for the same components; they could be weakly masing as suggested in Sect. 3.2. It is interesting to note that anomalous gas motion beyond the expansion velocity of the ring of maser clumps could perhaps explain the high velocity components. Such components are reminiscent of the weak features observed in the line wings of \uSiO emission from late type stars (Cernicharo et al. 1997, Herpin et al. 1998). SiO line wing emission in stars is related to bipolar gas outflows and to pulsations of the underlying star. \begin{figure} [t] \begin{center} \epsfxsize=8.5cm \epsfbox{figure6vit.epsf} \end{center} \caption[]{Comparison of \uSiO $v=1, J=2 \rightarrow 1$ and \dSiO $v=0, J=2 \rightarrow 1$ spot maps using the \uSiO feature at 15.6 km s$^{-1}$ as a phase reference in both maps. The epoch of the observations was March 5, 1996 for both isotopes. We have plotted the centroids of the \uSiO features (open squares and small full circles) together with the main \dSiO features (dotted circles). The diameter of each circle is proportional to the peak intensity for \dSiO. The LSR velocity labels are given every 5 and 6 channels for \uSiO and \dSiO, respectively.} \label{relat2829map} \end{figure} \begin{figure} [tbh] \begin{center} \epsfxsize=8.5cm \epsfbox{figure6rond.epsf} \end{center} \caption[]{Comparison of \uSiO $v=1, J=2 \rightarrow 1$ and \dSiO $v=0, J=2 \rightarrow 1$ spot maps using the \uSiO feature at 15.6 km s$^{-1}$ as a phase reference in both maps. The epoch of the observations was March 5, 1996 for both isotopes. We only show the centroids of \dSiO (dotted circles) and \uSiO (full circles) emissions. The diameter of each circle is proportional to the peak intensity.} \label{relat2829mapbis} \end{figure} \subsection{Nature of the \dSiO $v=0$ emission} Time variability as well as notable changes in the \dSiO line profile (Fig. \ref{variable29SiO}) are clearly in favour of maser emission. Short-term variability in the excitation of the \dSiO molecule is suggested by our maps because the mean distance observed in 1996 between the two ridges of emission, $\approx 0.13"-0.14"$, is significantly larger than the $0.11"$ measured in 1995. This fact indicates non thermal processes in the excitation of \dSiO. In addition, the relatively high flux density observed in 1995 and 1996 suggests also non thermal emission. The array cannot give the size of the individual features, but we may use the synthesized beamwidth and the observed peak flux density to estimate a minimum brightness temperature. In March 1996 the flux density peaks around 37 Jy and we derive $T_{B} \geq 2500$ K. This temperature is greatly above the kinetic temperature usually adopted in Orion, $\approx 60$ K, and is another indication for maser emission although high temperatures would be plausible in a shocked environment. However, for some of the weaker \dSiO features in the range $0\rightarrow 10$ \kms we obtain T$_{B}\geq 80-100$ K, and \dSiO could thus be part thermal and part maser. The actual spatial structure of the \dSiO $v=0$ emission is complex, and we recall that Fig. \ref{29SiOrelatmap} does not show all of the emission detectable with the array. By discarding the shorter baselines (Sect. 3.1), we have concentrated our analysis on more compact emission sources. We were not able here to map the \dSiO $v=0$ emission counterpart to the extended \uSiO $v=0$ emission seen by Wright et al. (1995). \subsection{Comparison of the \dSiO and \uSiO spot maps. Implication on excitation mechanisms} The relative distribution of \uSiO and \dSiO emission observed in 1996 is shown in Fig. \ref{relat2829map} and in Fig. \ref{relat2829mapbis} where, for clarity, we have not given the velocities. The two ridges of \dSiO $v=0, J=2 \rightarrow 1$ emission (dotted circles) and \uSiO $v=1, J=2 \rightarrow 1$ emission (full circles) are clearly visible on these figures. Our data do not show a complete spatial overlap as well as a close correlation among features of both species although in both cases the positive velocities lie to the NW of the negative features. Another clear difference between \dSiO and \uSiO is that there is no obvious pattern for the \dSiO intermediate velocity features (Fig. \ref{29SiOrelatmap} or Fig. \ref{relat2829map}); this could be related to mixed thermal and masing features as suggested in the previous Section. Fig. \ref{relat2829map} shows that velocities in the range $ \approx -6 \rightarrow -10$ km s$^{-1}$ tend to be found in the same area for both species although it is not possible to make an exact position-velocity pairing of the \dSiO and \uSiO features. In Fig. \ref{relat2829mapbis} the northern ridges of both \dSiO and \uSiO seem to be co-aligned and are not coincident. We cannot exclude, however, that uncorrected instrumental effects still affect our relative map and these observations should be repeated using frequent bandpass calibrations as used in 1996. In order to force the spatial coincidence of the main features in both species we have shifted the main \dSiO features lying around 16 km s$^{-1}$ on top of the \uSiO features in the range $15 \rightarrow 17$ km s$^{-1}$. Nevertheless, the \dSiO negative velocity emission ridge appears well outside the \uSiO negative velocity ridge. Any rotation of coordinates axis around the 16 km s$^{-1}$ features does not improve the spatial coincidence of both species. Hence, we conclude that both isotopic species are not excited in the same gas layers. This is strengthened by the analysis of our 1995 data which similarly show no spatial coincidence and a smaller distance between the two ridges of emission for \dSiO than for \uSiO . We note that a similar picture also emerges from the 43 GHz interferometric observations made by Morita et al. (1992) in Orion. Although their observations of the \uSiO $v=2, J=1 \rightarrow 0$ and $v=1, J=1 \rightarrow 0$ transitions were not made simultaneously and were less sensitive than here, the mean separation between the two ridges of emission is slightly smaller for $v=2$ ($\approx 0.12"$) than for $v=1$ ($\approx 0.14"$ in agreement with the separation measured on the 43 GHz map of Menten \& Reid 1995). The 43 GHz observations, the maps of \uSiO $v=1$ and $v=0$ emission (Wright et al. 1995) and our 86 GHz maps indicate that different vibrational states of silicon monoxide do show a close connection but do not exactly coincide. We comment below on possible explanations of the observed similarities without exact co-location of \dSiO ($v=0$) and \uSiO ($v=1$) maser sources. First, the silicon monoxide reservoir seems identical for both isotopic species since their large-scale spatial distributions are alike. This is expected if shocks generated in the expanding flow traced by the $v=1, J=2 \rightarrow 1$ masers enhance the sputtering of silicon which will then react quickly in the gas phase to form both \dSiO and \uSiO. Second, differences in the small-scale spatial distributions of \dSiO and \uSiO could simply result from differences in the excitation of both species or from different physical conditions within the silicon monoxide cloud. SiO pumping models do not require any isotopic differentiation to obtain \uSiO, \dSiO or $^{30}$SiO maser sources. In all cases the general physical conditions are grossly similar for one isotopic species or another apart from the total column densities. On the other hand, all radiative/collisional pumping models show that different $v$ state masers peak in different spatial regions. We believe that this fact, combined with different degrees of saturation in the \uSiO and \dSiO masers, is essential to explain the slightly different distribution of \dSiO ($v=0$) and \uSiO ($v=1$) maser spots. Our maps of relative \uSiO and \dSiO emission also show that some features from both isotopes and with different velocities tend to be excited in the same area. This is observed in the range $-6 \rightarrow -10$ km s$^{-1}$. Collisional pumping with high temperature ($\approx 1500$ K) and high molecular hydrogen density ($\approx 10^{9}-10^{10}$ cm$^{-3}$) provides a range of \uSiO column densities where both $v=1$ and $v=2$, $J=1 \rightarrow 0$ masers are excited (Lockett \& Elitzur 1992). Such a scheme does not apply to our apparently overlapping \dSiO ($v=0$) and \uSiO ($v=1$) 86 GHz features because their velocities are not in good agreement. However, further observations should be conducted to investigate the detailed kinematics and stability of the \dSiO emission. Line overlap effects among various transitions of silicon monoxide cannot be ignored to explain the excitation of this molecule. First, {\em local} line overlaps due to turbulence play a role as soon as the velocity dispersion reaches about 5 \kms. Limiting ourselves to the lower $J$ values, we find that 10 to 15 ro-vibrational transitions of \uSiO, \dSiO and \tSiO overlap within 5 \kms for $\Delta v=2$ and 1. If local line overlaps would dominate the excitation of low $J$ rotational levels in Orion, we would expect exact spatial coincidence of the isotopic species. Our \dSiO and \uSiO maps contradict this idea. Second, {\em non-local} line overlap effects as described by Gonz\'{a}lez-Alfonso \& Cernicharo (1997) in a non static circumstellar environment are most important. The relative distribution of \uSiO and \dSiO emission in our maps is not inconsistent with such non-local line overlaps. In addition, it is also plausible that the overlap between two near infrared lines of water and \uSiO (Olofsson et al. 1981 b) is an important excitation process of silicon monoxide in Orion. Analysing the spatial extents of different $v$ state masers is clearly important to better understand the pumping mechanisms of the SiO molecule. This kind of work should be extended to strong stellar SiO-emitters since the physical conditions in late-type stars and Orion are so different. VLBI observations are required in stars in order to make a detailed comparison of the different $v$ emission layers; such observations have been made for the first time to map the $v=2, v=1, J=1 \rightarrow 0$ lines in W Hya and VY CMa (Miyoshi et al. 1994). | 98 | 4 | astro-ph9804077_arXiv.txt |
9804 | astro-ph9804288_arXiv.txt | We have calculated several representative models of vertical structure of an accretion disk around a supermassive Kerr black hole. The interaction of radiation and matter is treated self-consistently, taking into account departures from LTE for calculating both the disk structure and the radiation field. The structural equations are described in detail, and various approximations are discussed. We have demonstrated that departures from LTE are very important for determining the disk structure, even at the midplane, as well as the emergent radiation, particularly for hot and electron--scattering--dominated disks. We have shown that at least for the disk parameters studied in this paper, NLTE effects tend to reduce the value of the Lyman jump with respect to the LTE predictions, regardless whether LTE predicts an emission or absorption jump. We have studied the effects of various values of viscosity on the model structure and predicted spectral energy distribution. The viscosity is parameterized through a parameter $\alpha_0$ which describes the vertically-averaged viscous stress, two power-law exponents $\zeta_0$ and $\zeta_1$, and the division point $m_{\rm d}$ between these two forms. The disk structure and emergent radiation is sensitive mainly to the values of $\alpha_0$, while the other parameters influence the disk structure to a much lesser extent. However, although the detailed shape of the predicted spectrum is sensitive to adopted value of $\alpha_0$, the overall appearance of the spectrum is quite similar. | Accretion disks around massive black holes have long been the most popular candidates for providing the ultraviolet and soft X-ray flux observed in Active Galactic Nuclei (AGN). Observational evidence is mainly based on the `big blue bumps' seen in the UV (e.g. Shields 1978; Malkan \& Sargent 1982). However, despite its attractiveness this model faces a number of problems when confronted with multi-wavelength observations. The current situation has been recently summarized by Koratkar (1998) from the observational point of view, and by Blaes (1998) from the theoretical viewpoint. The most pressing problems of the theoretical models are a near absence of observed Lyman discontinuity (first pointed out by Antonucci, Kinney, \& Ford 1989), the UV/EUV continuum polarization (e.g., Antonucci 1992, Blaes 1998 and references therein), the overall continuum spectral energy distribution (e.g. Laor 1990), and phased optical/UV variability (e.g., Alloin et al 1985). What is clearly needed is a self-consistent model which would explain all the observed features. Such a goal is still far away, but we feel that the first step towards it is to answer the fundamental question: In view of all current problems, is the accretion disk paradigm still a viable model? In other words, are the observed phenomena truly inconsistent with the accretion disk picture in general, or is the lack of agreement rather a result of inaccuracies or even inconsistencies in the computed models or in the current modeling techniques? Therefore, we have embarked on a systematic study of these questions. We recognize that there are many possible sources of inconsistencies in the AGN accretion disk modeling. There are essentially two types of problems. First, there are basic physical uncertainties. Among them, the most important is our ignorance of the basic physics of viscous energy dissipation in the AGN disks, so that we are left with a necessity to employ certain {\it ad hoc} parameters. Although this is certainly a viable approach when nothing else is currently available, we should bear in mind that corresponding models will lack predictive power---we may explain what we see, but we cannot predict it. In any case, when adopting a model based on some chosen set of parameters, one at least has to study carefully the influence of these parameters on the computed model. Other fundamental physical problems are uncertainties of the effects of other structures forming the AGN complex onto the accretion disk; for instance an irradiation of the disk from external sources. Again, in the absence of any detailed theory, we are usually left with a necessity to parameterize these effects. Second type of problems concerns the degree of approximation used in the actual modeling procedure. A particularly important class of such approximations deals with the description of interaction of radiation and matter in the disk. It should be emphasized that AGN disks, like atmospheres of hot stars, are typical examples of a medium where radiation is not only a {\em probe} of the physical state, but in fact a crucial {\em constituent}. In other words, radiation not only carries an information about the medium, it in fact {\em determines} its structure. Consequently, a treatment of this interaction is in a sense the very gist of the problem. We should therefore study the influence of various approximations, such as the degree of equilibrium assumed, or, more specifically, the extent to which the local thermodynamic equilibrium (LTE) applies, and the completeness of considered opacity and emissivity sources on the computed model, etc. In the previous paper (Hubeny \& Hubeny 1997 -- hereafter called Paper I), we have presented some representative self-consistent, non-LTE models of the vertical structure of AGN accretion disks. The basic aim of that study was to investigate the differences in the predicted spectrum between this and previous approaches (Sun \& Malkan 1989; Laor \& Netzer, 1989; Ross, Fabian, \& Mineshige 1992; Wehrse et al. 1993; St\"orzer and Hauschildt 1994; Coleman 1994; Shields \& Coleman 1994; Blaes \& Agol 1996; D\"orrer et al. 1996). The emphasis was to clarify the role of departures from LTE and to study the effects of simplifications of the hydrostatic equilibrium equation based on assuming a depth-independent vertical gravity acceleration. In the present paper, we will consider models of vertical structure of AGN disks in detail. In particular, we will study the dependence of computed vertical structure and corresponding emergent spectrum on the adopted value of viscosity and on the degree of sophistication of the modeling procedure. In order to better emphasize the observable consequences of self-consistent, non-LTE models, we present the predicted spectra for a few representative points on the disk. Complete spectra which are obtained by integrating the local spectra over the disk surface, taking into account general relativistic photon transfer functions (Cunningham 1975; Speith, Riffert, \& Ruder 1995; Agol 1997; Agol, Hubeny, \& Blaes 1998), will be considered in subsequent papers of this series. Our aim here is not to construct a model to be used for comparison with actual observations. Instead, we intend to study the sensitivity of computed models on the degree of approximations used in the modeling procedure. Therefore, we will first give a detailed overview of the structural equations and the modeling procedures. This is meant to provide a firm framework on which our approach is based, which in turn is useful to assessing possible systematic effects within our models. | We have calculated several representative models of vertical structure of an accretion disk around a supermassive Kerr black hole. The interaction of radiation and matter is treated self-consistently, taking into account departures from LTE for calculating both the disk structure and the radiation field. The viscosity is parameterized through the parameter $\alpha_0$ that describes the vertically averaged viscous stress, and two power--law exponents $\zeta_0$ and $\zeta_1$, and the division point $m_{\rm d}$ between these two forms. The disk structure and emergent radiation is sensitive mainly to the values of $\alpha_0$, while the other parameters influence the disk structure to a much lesser extent. However, although the detailed shape of the predicted spectrum is sensitive to adopted $\alpha_0$, the overall appearance of the spectrum is quite similar in all cases. We have shown that effects of departures from LTE are very important for determining the disk structure and emergent radiation, particularly for hot and electron--scattering dominated disks. We have shown that at least for the disk parameters studied in this paper, NLTE effects typically tend to diminish the value of the Lyman jump; in hot models they suppress the Schuster mechanism by which the LTE models produce a strong emission jump, and in cooler models they increase the flux in the Lyman continuum due to an underpopulation of the hydrogen ground state. Also, we have shown that relaxing the approximation of detailed radiative balance in the hydrogen and helium lines (i.e., computing the so-called NLTE/L models) changes the predicted line profiles significantly, but otherwise does not yield significant changes in computed vertical structure or emergent continuum flux. This result shows that for estimating the {\em continuum} radiation of AGN disks composed of hydrogen and helium, the NLTE/C models provide a satisfactory approximation. So far, we have limited our analysis to a simple H-He chemical composition. A preliminary study (Hubeny \& Hubeny 1998) indicates that the effects of numerous metal lines on the prediced spectral energy distribution of AGN disks may be quite significant. However, that study used a simplified approach in which the vertical structure was fixed by a H-He model, while the line opacity was taken into account only in the spectrum synthesis, assuming LTE source function in metal lines. Such an approach is inconsistent, since, first, the metal lines may change the disk vertical structure (the so-called metal line blanketing effects, long known form the theory of classical stellar atmospheres) and, second, the source function in metal lines may depart significantly from the LTE value. We will therefore need to construct self-consistent, fully metal-line-blanketed models of vertical structure of AGN disks, taking into account effects of literally millions of spectral lines in NLTE. A work of this project is under way, and will be reported in a future paper of this series. The results presented here do not indicate any fatal flaw of the AGN accretion disk paradigm. In contrast, they show that one of the previous critical arguments against the accretion disk paradigm, the magnitude of the Lyman jump, essentially disappears when increasing a degree of realism of the modeling procedure by relaxing previous simplifying approximations, in particular the local thermodynamic equilibrium and a simplified vertical disk structure. However, this study has concentrated on only one aspect of the problem, the spectral energy distribution in the optical, UV, and EUV region. Many questions, such as the overall spectral energy distribution of the whole disk, the effects of external irradiation, the continuum polarization, etc., remain to be explored in detail. This is exactly what we intend to do in future papers of this series. | 98 | 4 | astro-ph9804288_arXiv.txt |
9804 | astro-ph9804241_arXiv.txt | We report on an observation of the low-mass X--ray binary \U\ performed during the \B\ Science Verification Phase. An absorption feature at $\sim 37$~keV, attributable to electron cyclotron resonance, has been discovered in its pulse averaged spectrum. The inferred neutron star magnetic field strength is $3.2\cdot (1+z) \times 10^{12}$ G, where $z$ is the gravitational redshift. The feature is deep and narrow and is resolved in both the broad-band fit and in the ratio of observed counts to those seen from the Crab. The cyclotron resonance energy is in good agreement with the empirical relation between cyclotron energy and high energy cutoff, while its width is in agreement with the expected Doppler broadening of thermal electrons at the cyclotron resonance frequency. The broad-band 0.1--200 keV spectrum is well fit by a two-component model: a $0.27\pm 0.02$ keV blackbody and a power law with a photon index of $0.89\pm 0.02$. This is the first broad-band observation made after the change from spin-up to spin-down that occurred in mid 1990: it confirms the harder spectrum with respect to those observed in the 2--10 keV range. | Following the first detection in Her X--1 \cite{576}, line features in the 10--110 keV energy range were detected in some X--ray pulsars \cite{407,375}. These features are interpreted as being due to electron cyclotron transitions in the $\ga 10^{12}$~G magnetic field of the neutron star \cite{407}. Because the energy of the fundamental cyclotron harmonic, $E_{\rm cyc}$, is related to the neutron star magnetic field strength in units of $10^{12}$~G, $B_{12}$, by the relation $E_{\rm cyc} = 11.6\, B_{12}\cdot (1+z)^{-1}$~keV, where $z$ is the gravitational redshift, the observation of cyclotron resonance features (CRFs) in pulsar spectra gives a direct measurement of the magnetic field of the neutron star. From a theoretical point of view CRFs are expected to be visible as absorption features. This was shown {\em e.g.\/} by Nagel (1981) \nocite{306} for any reasonable optical depth. The feature is not strictly due to absorption, but rather to scattering of photons in the wings of the line. Up to now CRFs have been observed mainly in the spectra of young high-mass binary systems, with the notable exception of Her X--1 \cite{576}. The 7.7~s X--ray pulsar \U\ is one of the few pulsators among LMXRBs. Its timing history is characterized by a sudden change in its spin state that occurred on 1990 June \cite{1548}. Before this date the source was steadily spinning-up, while from the transition to the present time \U\ is spinning-down. Besides coherent pulsation, \U\ also shows quasi-periodic oscillations (QPOs) both in X--rays \cite{29} and in the optical \cite{1618}, with a centroid frequency of $\sim 40$ mHz. The pulse-phase averaged spectrum of \U\ has been described in terms of a two component model: an absorbed blackbody with a temperature $kT\sim 0.6$ keV and a power law \cite{658}. The spectrum also exhibits a high energy cutoff at $\sim 20$ keV, interpreted as due to a possible CRF at that energy \cite{617}, and an emission Ne line complex at $\sim 1$~keV \cite{393}. While the spin transition did not affect the high energy part of the spectrum, within the accuracy of previous measurements, it did affect the 2--10 keV part \cite{1590}: the power law photon index changed from $\sim 1.6$ \cite{617} to $\la 0.7$ \cite{393}, with a correspondingly significantly lower 2--10 keV flux, that shows an overall fading by a factor four from the HEAO1/A2 measurement \cite{1548}. | This is the first \U\ broad-band spectrum obtained after the spin transition. From our flux measurement we can see that the overall {\em bolometric\/} flux of the source decreased by a factor 4.1 since the HEAO1/A2 measurements. We find a power law index slightly higher than those found in narrow-band spectra \cite{393,1612}. This is an effect due to the broad--band fit, in which we simultaneously take into account both the low--energy power law and the high--energy cutoff. Our \U\ CRF measurement, together with its cutoff energy, fits the empirical relation found in X--ray binary pulsars between these parameters: higher the cyclotron resonance energy, higher the cutoff --- and therefore harder the spectrum \cite{61,1286}. Our CRF measurement also is in agreement with the expected electron Doppler broadening. Indeed, at the cyclotron resonance frequency $\omega_c$, electrons at rest absorb photons of energy $\hbar\omega_c$. For moving, thermal electrons the Doppler broadening $\Delta\omega_D$ is predicted to be \cite{1614} \begin{equation} \Delta\omega_D = \omega_c \left( \frac{2kT_e}{m_ec^2} \right)^{1/2} |\cos\theta| \label{eq:deltaD} \end{equation} where $kT_e$ is the electron temperature, and $m_ec^2$ is the electron rest mass. The angle $\theta$ measures the direction of the magnetic field with respect to the line of sight. Outside the range $\omega_c\pm\Delta\omega_D$ the cyclotron absorption coefficient decays exponentially, and other radiative processes become important. From Eq.~\ref{eq:deltaD} and the CRF parameters given in Table~\ref{tab:fit} we obtain a lower limit to the electron temperature of $\sim 5$ keV, in reasonable agreement with the calculations of self-emitting atmospheres of Harding et~al.\ (1984) \nocite{304}, according to which the temperature for a column or slab with optical depth $\sim 50$ g~cm$^{-2}$ is 4--8$\times 10^7~^\circ$K, corresponding to $kT_e\sim 3.5$--7 keV. Finally, we want to mention that if we assume the QPO frequency as due to the beating between the pulse frequency and the Keplerian motion at the magnetospheric radius, then the \U\ magnetic field strength is related to its luminosity by the relation $B_{12}\sim 5.5\sqrt{L_{37}}$ \cite{29}, where $L_{37}$ is the X--ray luminosity in units of $10^{37}$~erg~s$^{-1}$. Assuming a source distance $5\la \rm d_{\rm kpc} \la 13$ \cite{1618}, this corresponds to $2.4\la B_{12} \la 6.3$, in good agreement with our measurement. | 98 | 4 | astro-ph9804241_arXiv.txt |
9804 | gr-qc9804024_arXiv.txt | \noindent In the coming decade, the LIGO/VIRGO network of ground-based kilometer-scale laser interferometer gravitational wave detectors will open up a new Astronomical window on the Universe: gravitational waves in the frequency band $10$ to $10^4$ Hz. In addition, if the proposed, 5 million kilometer long, space based interferometer LISA flies, another window will be opened in the frequency band $10^{-4}$ to $1$ Hz. We review the various possible sources that might be detected in these frequency bands, and the information that might be obtainable from observed sources. Several key possible sources are inspirals and coalescences of neutron-star neutron-star and/or neutron-star black-hole binaries; inspirals, mergers, and ringdowns of black-hole black-hole binaries (both solar mass and supermassive); stellar core collapse; rapidly rotating neutron stars; the formation of supermassive black holes; and inspirals of compact objects into supermassive black holes. | This review of gravitational wave sources is divided into three sections. First, we review the detector sensitivities that have been achieved to date and discuss projected sensitivities for detectors now under construction. Second, we summarize the current observational upper limits on gravitational waves in various frequency bands. The main body of this review will consist of a survey of various anticipated sources of waves. Each anticipated source can be roughly characterized by a characteristic frequency $f$ and a characteristic value of strain amplitude $h$. However, it is important to also note that sources vary widely with respect to how uncertain is the rate of their occurrence in the Universe. The enterprise of anticipating potential gravitational wave sources is extremely uncertain. For most sources that we can conceive of, either the wave strengths are uncertain by several orders of magnitude, or the event rate is uncertain by several orders of magnitude, or the very existence of the source itself is very uncertain. While there are some important exceptions such as coalescing compact binaries, these are the exception rather than the rule. The upside of this great uncertainty is the potential for gravitational wave astronomy to bring us new and interesting information. For more details on the topics discussed here, the reader is encouraged to consult the detailed recent review articles by Thorne \cite{300yrs,thorne95,thorne97a,thorne97b}, and also the review article by Allen \cite{allen96} on stochastic gravitational waves. | With the kilometer scale interferometer network about to come online in the next few years, this is an exciting time for gravitational wave astronomy. Almost certainly we have not anticipated all the detectable sources of gravitational waves that Nature produces in the real Universe. Hopefully she will bring us some surprises. | 98 | 4 | gr-qc9804024_arXiv.txt |
9804 | astro-ph9804063_arXiv.txt | We analyze the linear, 3D response to tidal forcing of a disk that is thin and thermally stratified in the direction normal to the disk plane. We model the vertical disk structure locally as a polytrope which represents a disk of high optical depth. We solve the 3D gas-dynamic equations semi-analytically in the neighborhood of a Lindblad resonance. These solutions match asymptotically on to those valid away from resonances (previously obtained by Korycansky \& Pringle 1995) and provide solutions valid at all radii $r$. We obtain the following results. 1) A variety of waves are launched at resonance, including r modes and g modes. However, the f mode carries more than 95\% of the torque exerted at the resonance. 2) These 3D waves collectively transport exactly the amount of angular momentum predicted by the standard 2D resonant torque formula. 3) Near resonance, the f mode behaves compressibly and occupies the full vertical extent of the disk. Away from resonance, the f mode behaves incompressibly, becomes confined near the surface of the disk, and, in the absence of other dissipation mechanisms, damps via shocks. In general, the radial length scale for this process is roughly $r_{\rm L}/m$ (for resonant radius $r_{\rm L}$ and azimuthal tidal forcing wavenumber $m$), {\it independent of the disk thickness $H$}. This wave channeling process is due to the variations of physical quantities in $r$ and is not due to wave refraction. 4) However, the inwardly propagating f mode launched from an $m=2$ inner Lindblad resonance experiences relatively minor channeling (accompanied by about a factor of 5 increase in nonlinearity), all the way to the radial center of the disk. We conclude that for binary stars, tidally generated waves at Lindblad resonances in highly optically thick circumbinary disks are subject to strong nonlinear damping by the channeling mechanism, while those in circumstellar accretion disks are subject to weaker nonlinear effects. We also apply our results to waves excited by young planets for which $m \approx r/H$ and conclude that the waves are damped on the scale of a few $H$. | Gaseous disks are found in many types of binary star systems, including cataclysmic variables (CVs) and pre-main-sequence stars, and young planetary systems. The orbiting objects (stars or planets) exert tidal forces on these disks, which generally act merely to distort the disks from an axisymmetric form. However, at special locations in disks where resonances occur, the tidal forces generate waves that transport energy and angular momentum. As a result, a resonant torque is exerted by the system objects. The orbital evolution of the perturbing objects sometimes depends on the strength of these torques (Goldreich \& Tremaine 1980, hereafter GT80; Lin \& Papaloizou 1993; Lubow \& Artymowicz 1996). The waves transfer their angular momentum and energy to the disk in the regions of space where they damp, and this in turn affects the evolution of the disk. For example, gaps could be created in disks in regions where the waves damp. This paper concentrates on Lindblad resonances (LRs), which are due to horizontal forcing (along the disk plane). We assume throughout that the disk is coplanar with the orbit of the system objects. A 2D, linear theory for resonant tidal torques and associated wave propagation was developed by Goldreich \& Tremaine (1979, hereafter GT79). This theory provides an explicit formula for the Lindblad resonant torque. The 2D theory considers the disk to have only radial and azimuthal extent and ignores effects over its vertical extent (perpendicular to the orbit plane). Important progress has been made in the study of 2D nonlinear waves in disks (Shu, Yuan, \& Lissauer 1985; Yuan \& Cassen 1994). The torque in the nonlinear case was found to be within a few percent of that predicted by the 2D linear formula. The nonlinearities produce highly spiked density profiles which increase the level of dissipation present in a viscous disk. Radiative damping of linear waves can be important particularly when the disk is warm (Cassen \& Woolum 1996). Turbulent viscosity in the disk provides another means of wave damping. The 2D treatment is valid if both the vertical structure of the disk and its thermodynamic response are locally isothermal. Under such circumstances, tidal forcing will generate a 2D wave in a 3D disk. The wave front remains perpendicular to the disk plane at all heights, as the wave propagates radially in the disk. However, this 2D wave is highly singular in that it does not exist in a disk with a vertical temperature variation (Lin, Papaloizou, \& Savonije 1990a, hereafter LPS; Lubow \& Pringle 1993, hereafter LP). Furthermore, a vertically isothermal structure is not realistic for many important classes of gaseous disks, such as accretion disks in CVs, circumstellar and circumbinary disks of YSOs, and protoplanetary disks. Such disks often have optical depths much greater than unity and can be expected to have substantial vertical temperature variations, if they have an internal heat source such as turbulent dissipation. We demonstrate in this paper that 2D tidal forcing, caused by LRs, excites 3D waves in a thermally stratified disk. However, another class of resonances exist due to 3D effects. These resonances are due to vertical, tidal forcing by a coplanar perturber of a disk with nonzero thickness (Lubow 1981). The vertical resonances also generate horizontally propagating waves. Although intrinsically weaker than the LRs, the vertical resonances may be of importance in close binary star systems. Some investigations of 3D effects have been carried out using numerical simulations (Lin, Papaloizou, \& Savonije 1990a,b). Such approaches can be used to explore a limited range of physical parameter space. However, recent progress has been made in obtaining semi-analytic solutions for waves in 3D disks (LP; Korycansky \& Pringle 1995, hereafter KP). The aim of this paper is to extend that approach to understand 3D wave generation at LRs and the subsequent wave propagation. We determine the linear response of a thin (but nonzero thickness) disk. We model the disk locally as a polytrope in the vertical direction, which is valid for a disk of very high optical depth, and we ignore the effects of atmospheric layers. We aim to understand which modes of a thermally stratified disk are excited at LRs and how much torque is carried by such waves. By studying the properties of linear wave propagation, we can also understand where nonlinearity sets in that will likely lead to shocks and subsequent wave dissipation. The outline of this paper is as follows. In \S2, we review the properties of disk modes. In \S3 and \S4, we derive and solve the equations for waves generated at LRs. In \S5, we compute the total torque carried by these waves and determine which modes are excited. In \S6, we summarize properties of the f mode of the disk and present our numerical results. In \S7, we discuss the application of our results to binary and protoplanetary systems. \S8 contains a summary. | \subsection{Summary of results} We have analyzed the 3D response to tidal forcing of a gaseous disk having a vertical temperature variation. The unperturbed vertical disk structure was modeled as locally polytropic. We have considered effects at Lindblad resonances (LRs), subject to the thin disk approximation $H \ll r/m$, for which the LRs provide horizontal forcing that is independent of height. In the standard vertically isothermal case considered to date, the horizontally propagating 2D mode (sound wave) carries all the resonantly generated angular momentum, but only if the thermodynamic response of the disk is also isothermal. In the vertically polytropic case, the f mode has a role equivalent to that of the 2D mode, although r and g modes are also launched at LRs (see Figure~1). The mode launched at vertical resonances (Lubow 1981) can be shown to be the ${\rm p}_1^{\rm e}$ mode. We have found that nearly all the torque exerted at an LR is carried by the f mode (see Table~1). Near resonance, the f mode behaves in a similar manner to the 2D mode in the isothermal case, in that it occupies the full vertical extent of the disk and behaves somewhat compressibly. The f mode is almost two-dimensional near resonance in that the vertical velocities are smaller than the horizontal by a factor of order $(H/r)^{1/3}$. However, away from resonance, the behaviors of the f and 2D modes differ radically (see \S6.2). At a distance $\sim r_{\rm L}/m$ from the resonance, the f mode begins to concentrate its energy near the surface of the disk, a process which we have called wave channeling. In this regime, the f mode behaves like a surface gravity mode (see Plates~1 and~2). In most cases, the wave amplitude increases by many orders of magnitude with distance from resonance, so that it is very likely that shocks will develop which would damp the wave. An important exception to this last statement occurs for the $m=2$ ILR. The wave generated there undergoes relatively mild wave channeling and consequently a relatively mild increase in nonlinearity (by about a factor of 5). The amount of increase depends on the density distribution in the disk and could even result in a decrease in dimensionless amplitude near the disk center. Dissipation by turbulent viscosity is likely important in some cases, such as CV disks. On the other hand, the wave generated at the $m=2$ OLR does exhibit strong wave channeling. These results indicate that circumbinary disks, as found around young binaries, and protoplanetary disks perturbed by planets, are subject to strong effects of wave channeling (see \S7). \subsection{Discussion} The results in this paper differ from the previously accepted picture for wave propagation in thermally stratified disks (e.g. LPS). The standard expectation was that a wave launched at a Lindblad resonance would begin propagating horizontally, but the wavefront would be rapidly refracted upwards as a result of the decrease in sound speed with increasing height above the mid-plane. After advancing a distance comparable to the disk thickness $H$, the wavefront would be substantially tilted upwards. The wave would then propagate vertically into the atmosphere of the disk, where it would shock. This model is based on the idea that the launched wave is a pressure wave or p mode in a high-frequency acoustic limit, so that one can consider the wavefront to propagate at the local sound speed without being affected by inertial or buoyancy forces. Our results provide a different picture. We have found that the wave cannot be considered to propagate vertically, since it is in fact a vertically evanescent f mode. In this view, taken in LP and KP, the disk behaves like a waveguide in which the wave is vertically confined. However, somewhat similar to the standard picture, the wave energy does rise to the surface of the disk as the wave propagates away from the resonance. But this wave channeling process is effective over a distance of order $r_{\rm L}/m$, where $r_{\rm L}$ is the radius of the Lindblad resonance and $m$ is the azimuthal wavenumber. It does not depend on the disk thickness, provided that $H/r \ll1$. This can be understood roughly as follows. The wave channeling occurs because of the radial variation of the dimensionless intrinsic frequency of the wave, $(\omega-m\Omega)/\Omega$. This is equal to $\pm\kappa/\Omega$ at the Lindblad resonance, but increases rapidly in magnitude over a distance of order $r_{\rm L}/m$. The wave therefore proceeds rapidly along the f-mode branch of the dispersion relation. In the high-frequency limit, as described in \S6.1, the mode behaves like a surface gravity wave and is confined near the surface of the disk. (This would be true even in an incompressible disk in which refraction cannot operate.) The surface-gravity-mode dispersion relation is approximately $(\omega -m \Omega)^2 \approx g_{\rm s} k$, where the surface gravity $g_{\rm s}=\Omega^2 H$ and $k$ is the radial wavenumber. It can be seen (see \S6.1) that the mode is confined to the disk surface in a layer of vertical thickness $\delta \sim k^{-1}$. It then follows from the dispersion relation that $(\omega-m\Omega)^2/\Omega^2 \sim H/\delta $. As a result, the wave becomes confined (channeled) near the disk surface (i.e., $\delta /H \la 1/2$) over a radial distance from resonance of order $r_{\rm L}/m$. Another issue concerns the boundary conditions. In LPS, the anticipated shocks in the disk atmosphere were represented in their simulation by an upper boundary condition that included some amount of dissipation. In our analysis of a polytropic disk, we have applied a boundary condition at the disk surface that acts to reflect waves rather than absorb energy. Even in a purely isothermal disk without a definite surface, as demonstrated by LP, all the modes are confined vertically and do not propagate vertically to infinity. This is a consequence of the increase in vertical gravity with height. Our treatment of the boundary is valid in the limit of high optical depth, where the atmosphere occupies a negligible mass. In particular, the effects of atmosphere are not strongly felt by the wave until $k H_{\rm atmos} > 1$, with $H_{\rm atmos}$ being the density scale height at the base of the atmosphere. For a highly optically thick disk, $H$ is much greater than $H_{\rm atmos}$, and so $k H \gg 1$. When this condition occurs, the wave has been strongly channeled before much wave energy enters the atmosphere. We plan to extend our current analysis to include the effects of a disk atmosphere in a future paper. The asymmetry in wave properties between the inner and outer LRs is most pronounced for low-$m$ cases. For fixed $m$, the level of nonlinearity for ILR waves is lower than for OLR waves (see \S6.2.2). One reason is that the density increases as the ILR wave propagates inward, while the density decreases as the OLR wave propagates outwards (see eq. [\ref{calN2}]). The other cause of the asymmetry is that the f-mode dispersion relation is followed to higher wavenumbers in the exterior of the OLR than in the interior of the ILR, for low $m$ (see Figure~3). Away from the LR, the f mode is vertically confined to a surface layer of thickness $\sim k^{-1}$. Due to wave action conservation, the confinement acts to increase the wave amplitude, causing substantial asymmetry in the case $m=2$. In the high-$m$ case, the behavior of the wavenumber becomes similar near the inner and outer LRs, and consequently the nonlinearity becomes important at comparable distances from the respective resonances. A variant on the vertically polytropic model occurs for layered disks (Gammie 1996). In this model, some regions of protostellar or protoplanetary disks have turbulence restricted to the upper layers of the disk. The resulting vertical temperature structure is isothermal in the non-turbulent layer near mid-plane and the temperature declines with height in the outer turbulent layers. A wave launched from an LR in such a disk would probably require a longer distance from resonance to undergo wave channeling. Given the large increases in wave velocities associated with wave channeling (see Plates~1 and~2), we expect that nonlinearities will cause the waves to shock and deposit angular momentum preferentially near the surface of the disk. We have described where nonlinearities are expected to become important in \S6.2, based on a quasilinear estimate. Little work has been carried out to investigate the nonlinear behavior of f modes or surface gravity waves in a compressible fluid. It would be useful to conduct nonlinear simulations of the waves in such a disk. | 98 | 4 | astro-ph9804063_arXiv.txt |
9804 | astro-ph9804313_arXiv.txt | We have used the Deep Survey telescope of EUVE to investigate shadows in the diffuse EUV/Soft X-Ray background cast by clouds in the interstellar medium. We confirm the existence of a shadow previously reported, and provide evidence for two new shadows. We used IRAS data to identify the clouds producing these shadows and to determine their optical depth to EUV radiation. The EUV-absorbing clouds are optically thick in the EUV, and all EUV emission detected in the direction of these shadows must be produced from material in front of the clouds. We obtained new optical data to determine the distance to these clouds. We use a new differential cloud technique to obtain the pressure of the interstellar medium. These results do not depend on any zero level calibration of the data. Our results provide evidence that the pressure of the hot interstellar gas is the same in three different directions in the local interstellar medium, and is at least 8 times higher than derived for the local cloud surrounding our Sun. This provides new evidence for large thermal pressure imbalances in the local ISM, and directly contradicts the basic assumption of thermal pressure equilibrium used in almost all present models of the interstellar medium. | The discovery of the soft X-ray / EUV background (\cite{bow68}) and its anti-correlation with the Galactic H{\sc i} distribution has led to an on-going debate on the origin of that emission. Since an absorption column density of $N_{\rm H{\sc i}} = 1.2 \times 10^{20}$\,cm$^{-2}$ (which corresponds to an optical depth $\tau = 1$ at 0.25\,keV) is reached at a distance of approximately 100\,pc in the direction of the Galactic plane, the diffuse background must arise from hot gas in the local interstellar medium (ISM). Optical and UV absorption line measurements of nearby stars together with observations in the X-ray and EUV wavelength range have established a local region with a significant H{\sc i} deficiency which is partly filled with a plasma, the so-called ``Local Bubble'' (LB). No definite model exists for this region and its origin is unclear. One view is that the gas in the LB has been heated by a supernova explosion which occurred in the solar neighborhood about $10^7$\,yr ago (cf. \cite{mckost77}, \cite{coxrey87}). Other authors claim that the LB is part of an asymmetrically shaped superbubble created by stellar winds and supernova explosions (cf. \cite{frisch95}). However, the physical properties such as density, temperature, pressure, and extension of the LB, are not well determined and so far observational constraints are insufficient to establish a canonical model of the evolution and the origin of the LB. The McKee \& Ostriker model is the best known model but it is not easily testable. Only soft X-ray and EUV observations provide a direct method to study the hot gas in the ISM and test LB models. This paper presents new measurements for the thermal pressure in the local ISM, using observations of nearby clouds taken by the Deep Survey (DS) telescope on-board the {\em Extreme Ultraviolet Explorer} (EUVE). Prior to this paper, Bowyer~et~al.~(\cite{bow95}) reported on the first detection of a spatial absorption feature in the diffuse EUV background; this feature was positionally coincident with an IRAS cirrus cloud at a distance of $\la$ 40\,pc. Since the cloud casting this shadow is optically thick at EUV wavelengths, the authors concluded that all residual EUV emission observed at the position of the cloud originated from material along the line of sight in front of the cloud. By attributing the background subtracted flux in front of the cloud to the hot gas in the ISM, Bowyer~et~al.~(\cite{bow95}) derived for the length of the emitting region (40\,pc) a pressure (P/$k$) of 19,000\,cm$^{-3}$K. These authors pointed out that an imperfect zero level calibration of the data can lead to a smaller value for the pressure, and estimated a lower limit for the pressure of 7,000\,cm$^{-3}$K. The purpose of this paper is to present a re-observation of the previously reported cloud shadow in the EUV background and observations of two new shadows discovered in the DS. For all three shadow regions we provide new optical photometry data which establishes the distance of the respective clouds. We employ a differential cloud technique which does not depend on any zero level calibration, to obtain the pressure of the ISM in these directions. The outline of this paper is as follows: First, in Sect.~2 we describe our new observations and the data analysis. In Sect.~3, we employ a new analysis technique and compare our results for the three cloud shadows. We derive physical quantities of the hot interstellar gas from our observations. Finally, in Sect.~4, we discuss the implications for models of the local ISM. | \label{discus} We have presented a re-observation of a cloud shadow in the diffuse EUV background discovered by Bowyer~et~al.~(\cite{bow95}) and present two additional shadows detected in the EUVE DS. We have developed a new method to derive the diffuse astrophysical EUV background from the EUVE DS data. This differential cloud technique can be applied to any cloud shadow data without a known zero level calibration. In our case we can extrapolate the results to the origin and obtain a zero level calibration for the DS detector. This zero level calibration indicates that about 50\% of the detected photons are astronomical in origin. More importantly, this zero level calibration allows a determination of the diffuse EUV background in the direction of the three nearby cloud shadows and hence to derive the pressure in the regions between us and these clouds. The pressure obtained for these nearby regions is consistent with the value obtained from the differential measurements for the more distant regions along the lines of sight. The results of our cloud shadow observations with EUVE provide evidence for a constant pressure in three different directions in the local ISM. We derived a pressure of P/$k \approx 16500 {\rm cm}^{-3}$\,K. The canonical value obtained from the analysis of solar He~{\sc i} 584\,\AA~~radiation resonantly scattered by helium in the inflowing cloud is P/$k = 730 \pm 30 {\rm cm}^{-3}$\,K (e.g., \cite{frisch95}). Based on line measurements with EUVE, Vallerga (\cite{vallerga96}) derived a pressure for the solar cloud in the range P/$k = 1700 - 2300 {\rm cm}^{-3}$\,K which is in good agreement with results obtained by \cite{bertetal85}, P/$k \approx 2600 {\rm cm}^{-3}$\,K). A comparison of these results shows the pressure of the local ISM exceeds the pressure of the cloud surrounding our Sun by a factor of $\ge 8$. The original observation (Bowyer~et~al.~\cite{bow95}) of a pressure imbalance between the solar cloud and the surrounding local ISM could have been the result of a nearby shock wave which had not yet reached the Sun. The results reported here are obtained in three different directions. The isotropy in the results for all these directions is not consistent with this hypothesis. Flux in the soft X-ray bandpass requires a larger absorption column density to reach unit optical depth and it is difficult or impossible to separate the locally produced diffuse soft X-ray emission from contributions originating at larger distances. Soft X-ray observations, therefore, typically provide only an upper limit for the pressure of the local ISM. As discussed previously, the pressure obtained from our cloud shadow observations in the EUV is strictly interpreted as a lower limit. Although different plasma codes have been used by different authors, the pressure derived from our cloud shadow observations with EUVE is about the same as values derived from soft X-ray observations (e.g., \cite{snowden97}, \cite{freyberg97})when we use the same code (P/$k$(EUV) = 13500 cm$^{-3}$K; P/$k$(soft X-ray) = 14000 cm$^{-3}$K). The near equality of these upper and lower limits implies we have measured the true pressure. Beyond the cloud surrounding our Sun we found no significant intrinsic cold gas absorption in the local ISM. This supports the idea of a cavity primarily filled with an ionized plasma. Interstellar absorption lines in the spectra of nearby stars have indicated the existence of cloudlets of cold gas (like the local cloud around our Sun) in the solar vicinity (Lallement~\cite{lallement96}). However, the effect of these cloudlets on the diffuse EUV emission observed in the directions of the cloud shadows discussed here is negligible. We point out that the different line of sights for the six clouds discussed here show that the local ISM cannot be simply described as a sphere with an almost constant radius (\cite{snowcox90}). Our observations provide evidence for a large pressure imbalance in the local ISM compared to the local cloud surrounding our Sun. This contradicts the basic assumption of almost all available models for the local ISM of a pressure equilibrium (e.g., \cite{mckost77}, \cite{coxrey87}). It is unclear what physical mechanism can maintain such a large pressure imbalance. We emphasize that these models as well as the calculations in the here presented work assume a hot plasma in collisional ionization equilibrium. A completely different approach is provided by Breitschwerdt \& Schmutzler (\cite{breischm94}). Their model calculations are based on an adiabatically cooling non-equilibrium plasma. The advantage of this model is that it can at least qualitatively explain many of the observed features of the local ISM and it does not require pressure equilibrium in the local ISM. | 98 | 4 | astro-ph9804313_arXiv.txt |
9804 | astro-ph9804125_arXiv.txt | It has been recently realized (Rephaeli 1995) that the relativistic corrections to the spectral distortions of the cosmic microwave background (CMB) measured in the direction of clusters of galaxies containing hot gas are significant and should be detectable with the forthcoming experiments. In the present paper we calculate the correction terms that are proportional to $V_{\rm r}/c\times kT_e/m_ec^2$ and $(V/c)^2$ to the standard formulae describing the spectral distortions caused by the bulk motion of the free electrons (kinematic effect) and due to the presence of the hot gas (thermal effect) for the case of a cluster having a peculiar velocity $V$ ($V_{\rm r}$ is its radial component). The results of our analytical calculations are confirmed by Monte-Carlo simulations (Sazonov \& Sunyaev 1998). | Thomson scattering of the cosmic microwave background (CMB) radiation by hot electrons in the intergalactic gas in clusters of galaxies modifies the spectrum of the CMB (Sunyaev and Zel'dovich 1972). Zel'dovich and Sunyaev (1969), basing on the Kompaneets equation (1957), derived a simple formula describing the spectral form of the distortion, which is proportional to the parameter $y=(kT_e/m_ec^2) \tau$, where $\tau$ is the Thomson optical depth along the line of sight. The effect has now been observed from a number of clusters of galaxies (see Birkinshaw 1998 for review). Recently, interest to this effect has been reactivated in view of the perspectives of accurate measurement of the CMB distortions in a number of experiments, both ground-based and on balloons, by the MAP spacecraft and especially by the Planck Surveyor mission scheduled to be flown in the middle of the next decade. These activities were motivated by the fact that the gas temperature is so high in the clusters of galaxies (ranging between 3 and 17~keV, Tucker et al. 1998) that the scattering electrons have thermal velocities of the order of 0.1 -- 0.3 $c$, so one has to include into consideration the relativistic corrections to obtain an accurate result. Rephaeli (1995), basing on extensive previous work (Wright 1979, Fabbri 1981, Taylor \& Wright 1989, Loeb et al. 1991) has demonstrated by means of numerical calculations the relevance of the relativistic corrections for the future experiments. Stebbins (1997), Itoh et al. (1998), and Challinor \& Lasenby (1998) used a Fokker-Planck approximation of the relativistic photon kinetic equation to obtain corrections, written as series in powers of $kT_e/m_ec^2$, to the standard nonrelativistic solution. These results have proved to be in excellent agreement with those of Rephaeli, demonstrating the applicability of the diffusion approximation to the problem at hand, despite the small optical depths of the clusters of galaxies ($\tau \sim 0.01$). A gas cloud moving rapidly relative to the CMB along the observer's line of sight must significantly modify the spectrum of the CMB in addition to the thermal effect. The change in the brightness temperature caused by this ``kinematic'' effect is to first order simply proportional to the radial component of the cluster velocity $\sim (V_{\rm r}/c)\tau$ (Sunyaev \& Zel'dovich 1980). The effect should be detectable in the future, and will enable measurement of cluster peculiar velocities, with significant implications for studies of the large-scale structure of the universe. It is obvious that corrections similar to those found for the thermal effect must exist and should be taken into account for the kinematic effect, if one wants to find the correct solution for the case of a moving cluster. In this paper we calculate the next-order changes in the spectrum of the CMB related to the cluster peculiar velocity by solving the photon kinetic equation. We have obtained simple formulae giving the correction terms of the orders of $V_{\rm r}/c\times kT_e/m_ec^2$ and $(V_{\rm r}/c)^2$. Our method is similar to that used by Psaltis \& Lamb (1997) who considered the more general problem of comptonization in a moving media. The solution these authors have obtained, although applicable to many astrophysical situations, does not contain the $O(V_{\rm r}/c\times kT_e/m_ec^2)$ term, because this term is third-order in electron velocity, whereas their solution is accurate only to second order in it. We aslo confirm the existence of the term of order $(kT_e/m_ec^2)^2$ found earlier using techniques different from ours (Rephaeli 1995, Stebbins 1997, Itoh et al. 1998, and Challinor \& Lasenby 1998). Earlier we have found all the correction terms mentioned above numerically using Monte-Carlo simulations (Sazonov \& Sunyaev 1998). | 98 | 4 | astro-ph9804125_arXiv.txt |
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9804 | astro-ph9804196_arXiv.txt | We present a new calculation of the intergalactic \gray\ pair-production absorption coefficient as a function of both energy and redshift. In reexamining this problem, we make use of a new {\it empirially} based calculation (as opposed to previous {\it model} calculations) of the intergalactic infrared radiation field (IIRF). We find smaller opacities than those given previously (Stecker \& De Jager 1997). We apply our results to the new observations of the flaring \gray\ spectra of Mrk 421 and Mrk 501, both at a redshift of $\sim$0.03. Our new calculations indicate that there should be no significant curvature in the spectra of these sources for energies below 10 TeV, as indicated by recent observations. However, the intrinsic spectra of these sources should be harder by amounts of $\sim$ 0.25 to 0.45 in the spectral index (in the 1 - 10 TeV range), with an intergalactic absorption cutoff above $\sim 20$ TeV. | We have previously pointed out (Stecker, De Jager \& Salamon 1992 (hereafter, SDS92)) that very high energy \gray\ beams from blazars can be used to measure the intergalactic infrared radiation field, since pair-production interactions of \grays\ with intergalactic IR photons will attenuate the high-energy ends of blazar spectra. Determining the intergalactic IR field, in turn, allows us to model the evolution of the galaxies which produce it. As energy thresholds are lowered in both existing and planned ground-based air Cherenkov light detectors (Cresti 1996), cutoffs in the \gray\ spectra of more distant blazars are expected, owing to extinction by the IIRF. These can be used to explore the redshift dependence of the IIRF (Stecker \& Salamon 1997; Salamon \& Stecker 1998). Furthermore, by using blazars for a determination of attenuation as a function of redshift, combined with a direct observation of the IR background from the {\it DIRBE} detector on {\it COBE}, one can, in principle, measure of the Hubble constant $H_{0}$ at truly cosmological distances (Salamon, Stecker \& De Jager 1994). There are now over 50 grazars which have been detected by the {\it EGRET} team (Thompson, \etal\ 1996). These sources, optically violent variable quasars and BL Lac objects, have been detected out to a redshift greater that 2. Of all of the blazars detected by {\it EGRET}, only the low-redshift BL Lac, Mrk 421, has been seen by the Whipple telescope. The fact that the Whipple team did not detect the much brighter {\it EGRET} source, 3C279, at TeV energies (Vacanti, \etal\ 1990, Kerrick, \etal\ 1993) is consistent with the predictions of a cutoff for a source at its much higher redshift of 0.54 (see SDS92). So too is the recent observation of two other very close BL Lacs ($z < 0.05$), {\it viz.}, Mrk 501 (Quinn, \etal\ 1996) and 1ES2344+514 (Catanese, \etal\ 1997) which were too faint at GeV energies to be seen by {\it EGRET}. In this paper, we calculate the absorption coefficient of intergalactic space using a new, empirically based calculation of the spectral energy distribution (SED) of intergalactic low energy photons (Malkan \& Stecker 1998; hereafter MS98) obtained by integrating luminosity dependent infrared spectra of galaxies over their luminosity and redshift distributions. After giving our results on the \gray\ optical depth as a function of energy and redshift out to a redshift of 0.3, we apply our calculations by comparing our results with recent spectral data on Mrk 421 as reported by McEnery, \etal\ (1997) and spectral data on Mrk 501 given by Aharonian, \etal\ (1997). The results presented here supercede those of our previous calculations (Stecker \& De Jager 1997), which were based more on theoretical models (see discussion in MS98). We consider the results presented here to be considerably more reliable than any presented previously. | We have calculated the absorption coefficient of intergalactic space from pair-production interactions with low energy photons of the IIRF, both as a function of energy and redshift, using new, more reliable estimates of the SED for the IIRF which were obtained by MS98. Our results predict less absorption than we obtained previously (Stecker \& De Jager 1997), because the MS98 IIRF is lower than our previous estimate. One reason for this difference is that our previous IIRF spectrum was normalized partly to reflect the estimate of Gregorich, {\it et al.} (1995) of the IIRF at 60 \mic. According to Bertin, Dennefeld \& Moshir (1997), that estimate may have been based on the inclusion of false detections in their analysis of the IRAS data. For absorption calculations, it is important to use the most reliable estimate of the IIRF avaliable, since the absorption effect depends exponentially on the magnitude of the IIRF. While we do not claim that our new results for the absorption coeeficient as a function of energy differ dramatically from those obtained previously (MacMinn \& Primack 1996; Stecker \& De Jager 1997), we {\it do} claim that they are more reliable because they are based on the empirically derived IIRF given by MS98, whereas all previous calculations of TeV $\gamma$-ray absorption were based on theoretical modeling of the IIRF. The MS98 calculation was based on data from nearly 3000 IRAS galaxies. These data included (1) the luminosity dependent infrared SEDs of galaxies, (2) the 60$\mu$m luminosity function of galaxies and, (3) the redshift distribution of galaxies. We have applied our absorption calculations to recent flaring spectra of the nearby BL Lac objects Mrk 421 and Mrk 501. The spectral calculations given here are in good agreement with the recent observations that indicate no significant curvature in the spectra of Mrk 421 and Mrk 501. The observations are also consistent with our calculated steepening of 0.25 to 0.45 in the spectral index of these sources in the 1-10 TeV range. Our new calculations predict a significant intergalactic absorption effect which should cut off the spectra of Mrk 421 and Mrk 501 at energies greater than $\sim$20 TeV. Observations of these objects at large zenith angles, which give large effective threshold energies, may thus demonstrate the effect of intergalactic absorption. Our new calculations confirm the conclusion in SDS92 that TeV spectra of sources at redshifts higher than 0.1 should suffer significant absorption. The recent detection of another XBL at a redshift below 0.1, {\it viz.}, 1ES2344+514 (Catanese, \etal\ 1997), further supports the argument that nearby XBLs may be the only significant TeV sources presently detectable (Stecker, De Jager \& Salamon 1996). | 98 | 4 | astro-ph9804196_arXiv.txt |
9804 | astro-ph9804019_arXiv.txt | The Hubble Deep Field South (HDFS) has been recently selected and the observations are planned for October 1998. We present a high resolution (FWHM $\simeq 14$ \kms) spectrum of the quasar J2233--606 ($z_{em}\simeq2.22$) which is located 5.1 arcmin East of the HDFS. The spectrum obtained with the New Technology Telescope redward of the Lyman--$\alpha$ emission line covers the spectral range 4386--8270 \AA. This range corresponds to redshift intervals for CIV and MgII intervening systems of $z=1.83-2.25$ and $z=0.57-1.95$ respectively. The data reveal the presence of two complex intervening CIV systems at redshift $z=1.869$ and $z=1.943$ and two complex associated ($z_{abs} \approx z_{em}$) systems. Other two CIV systems at $z=1.7865$ and $z=2.077$, suggested by the presence of strong Lyman--$\alpha$ lines in low resolution ground based and Hubble Space Telescope (HST) STIS observations (Sealey et al. 1998) have been identified. The system at $z=1.943$ is also responsible for the Lyman limit absorption seen in the HST/STIS spectrum. The main goal of the present work is to provide astronomers interested in the Hubble Deep Field South program with information related to absorbing structures at high redshift, which are distributed along the nearby QSO line of sight. For this purpose, the reduced spectrum, obtained from three hours of integration time, has been released to the astronomical community. | The Hubble Deep Field (HDF) program (Williams et al. 1996), which was carried out with the Hubble Space Telescope (HST) in December 1995 in the four filters $UBVI$ to the deepest ever reached limiting magnitude, together with followup observations in other spectral bands, can be considered as one of the major astronomical events of the nineties. More than 40 articles on this program have been published in refereed journals between April 1996 and December 1997, confirming the tremendous impact of the observations. Articles related to the high redshift Universe ($z\gsim1$) treat such subjects as the galaxy redshift distribution (Gwyn \& Hartwick 1996, Seidel et al. 1996, Lanzetta et al. 1996, Lowenthal et al. 1997), the Global Star Formation History (Madau et al. 1996, Connoly et al. 1997, Guzman et al. 1997) and the clustering properties of galaxies (Colley et al. 1996, Villumsen et al. 1997). After two years, the Hubble Deep Field South (HDFS) has been planned for Cycle 7 of the HST (Williams et al. 1997, see also the HDF--South Web site at the URL \\ http://www.stsci.edu/ftp/science/hdf/hdfsouth/hdfs.html), and the observations will be performed in October 1998. Around the selected Southern field only one high redshift quasar J2233--606 ($z_{em}\simeq2.22$) has been identified (Boyle 1997), and it is located 5.1 arcmin EAST of the HDFS. Two HST orbits have been dedicated to this object to obtain a low resolution STIS spectrum (3700 seconds with G230L and 2200 seconds with G430L). The spectrum in the range $\lambda\lambda=1600-5700$ \AA~reveals, among other features, the presence of a Lyman limit break at $z\sim1.9$, a few metal lines in the red part of the \lya~emission line and a \lya~forest with a few absorption lines on the top of the \lya~emission lines. No Damped \lya~profile is seen along the line of sight. Ground--based low resolution observations combined with the STIS spectrum have been analyzed by Sealey et al. (1998). They report a first tentative identification of metal lines and provide a HI column density of the Lyman limit system at $z=1.943$, of $N(\rm HI) = (3.1\pm1.0)\times10^{17}$ \cm2. Here we report high spectral resolution (FWHM $\simeq 14$ \kms) observations obtained with the New Technology Telescope (NTT) of the nearby QSO J2233--606 ($z_{em}\simeq2.22$) in the spectral range $\lambda\lambda=4386-8270$ \AA. Even though the 3 hours of integration time have not produced an extremely high quality spectrum, the reduced data have been released to the astronomical community (see the URL http://www.eso.org/ $\tilde{}$ ssavagli). It is hoped that these data will provide a valuable guide when planning HDFS followup observations and will assist scientific discussions on the relation between the absorption systems along the QSO line of sight and the emitting objects which will be identified in the HDFS. Other high resolution observations from the ground primarily covering the Lyman--$\alpha$ forest region have been scheduled for July 1998 at the high resolution spectrograph UCLES of the Anglo Australian Telescope (AAT). | Unlike in the case of the HDF North program, ``follow--up'' observations of the HDFS region from the ground and space telescopes began already long before the beginning of the program. The observations from the radio to X--rays, from low and high resolution spectroscopy to narrow and broad band imaging, programmed since November 1997 to the Spring of 1999, will be combined to provide astronomers with the deepest and most complete view of the Universe for redshifts $z>1$, where most of the stars have probably already formed. Moreover, unlike the Northern field, the HDFS is close to a high redshift quasar. The QSO line of sight ($\alpha = 22^h 33' 37.67''$, $\delta = -60^\circ 33' 28.95"$, J2000 Equinox) is located $5'7''$ away from the Hubble Deep Field South ($\alpha = 22^h 32' 56.22''$, $\delta = -60^\circ 33' 02.69"$, J2000 Equinox). This corresponds to a real separation that is shown in Fig.~11 in the redshift range $z=0.5-3$. The separation at the QSO redshift ($z_{em}\simeq2.22$) is 2.8 and 8.9 $h_{50}^{-1}$ Mpc ($h_{50}$ is the Hubble constant expressed in units of 50 \kms~Mpc$^{-1}$) for a flat Universe in the physical and comoving space respectively. Deep imaging of the region near the QSO line of sight is fundamental to the understanding the quasar environment, the origin of high redshift quasars in connection to that of galaxy clusters and the relation between \lya~absorbers and high redshift galaxies. The typical wavelength of density fluctuations whose collapse generate galaxy clusters is around 20 $h_{50}^{-1}$ Mpc. On the other hand, once virialized, clusters have a typical size (i.e., virial radius) of about $2-5~h_{50}^{-1}$ Mpc, depending on their mass and on the background cosmological model (e.g., Kitayama \& Suto 1996, and references therein). At redshift $z\sim 2.2$ we generally expect to find protoclusters which are not yet virialized. Instead, they should correspond to detectable overdensities involving scales of about 10 $h_{50}^{-1}$ Mpc comoving. Therefore, that would roughly correspond to the comoving separation between the HDFS and the main absorption clusters seen in the QSO spectrum, whereas at the same redshift, the comoving HDFS size is around 4.5 $h_{50}^{-1}$ Mpc. The size of the HFDS and the separation from the quasar might be suitable for the identification of clustered structures associated with the QSO absorption systems. The idea of the existence of clustered structures at high redshifts has recently been confirmed by observations. Steidel et al. (1998) found evidence for a protocluster of galaxies at $z\simeq 3.1$ in the field of two high redshift QSOs, one of which is at the same redshift as the structure. The discovered 15 Lyman break galaxies are distributed on a plane which is at least $20\times15~h_{50}^{-2}$ Mpc$^2$ comoving. Numerical simulations have shown that non--linear structures can be present on large scales already by $z=3$. Using N--body simulations and semi--analytical methods, Governato et al.~(1998) were able to produce in the generic Cold Dark Matter scenario, structures similar to the observed one. These high redshift associations of Lyman break galaxies are strongly biased with respect to the already overdense local dark matter distribution and are the seeds of present day galaxy clusters. Other protoclusters of galaxies candidates at very high redshift ($z\sim2$) have been found in a limited number of cases (Dressler et al. 1993, Francis et al. 1996, Pascarelle 1996, Hutchings 1995). At low redshift ($z<1$) extensive imaging and spectroscopic observations of galaxy clusters in quasar environments and other AGNs indicated that quasar activity is strongly correlated with the properties of the environment of the host galaxy (Yee \& Green, 1987; Yee\& Ellingson, 1993). The study of the HDFS QSO should allow to extend this concept to much higher redshifts. In particular, it should reveal if the highest density peaks that give origin to quasars at high redshift are also the seeds of structure formation on much larger scales. | 98 | 4 | astro-ph9804019_arXiv.txt |
9804 | hep-ph9804230_arXiv.txt | For neutrinos with a magnetic moment, we show that the collisions in a hot and dense plasma act as an efficient mechanism for the conversion of $\nu_L$ into $\nu_R$. The production rate for right-handed neutrinos is computed in terms of a resummed photon propagator which consistently incorporates the background effects. Assuming that the entire energy in a supernova collapse is not carried away by the $\nu_R$, our results can be used to place an upper limit on the neutrino magnetic moment $\mu_\nu < (0.1-0.4)\times 10^{-11}\mu_B$. | 98 | 4 | hep-ph9804230_arXiv.txt |
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9804 | astro-ph9804042_arXiv.txt | s{ The measurements of CMB anisotropy have opened up a window for probing the global topology of the universe on length scales comparable to and beyond the Hubble radius. We have developed a new method for calculating the CMB anisotropy in models with nontrivial topology and apply it to open universe models with compact spatial topology. We conduct a Bayesian probability analysis for a selection of models which confronts the theoretical pixel-pixel temperature correlation function with the {\sc cobe--dmr} data. Our results demonstrate that strong constraints on compactness arise: if the universe is small compared to the `horizon' size, correlations appear in the maps that are irreconcilable with the observations.} \noindent The remarkable degree of isotropy of the cosmic microwave background (CMB) points to homogeneous and isotropic Friedmann-Robertson-Walker (FRW) models for the universe. This argument is a purely local one and does not refer to the global topological structure of the universe. In fact, in the absence of spatially inhomogeneous perturbations, a FRW model predicts an isotropic CMB regardless of the global topological structure. However, the observed large scale structure in the universe and CMB anisotropy allude to the existence of small spatially inhomogeneous primordial perturbations. The global topology of the universe does affect the observable properties of the CMB anisotropy. In compact universe models, the finite spatial size usually implies a suppression of the power in large scale perturbations and consequently the CMB anisotropy is suppressed on angular scales above a characteristic angle related to size of the universe. Another signature is the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the compact manifold. Much recent astrophysical data suggest the cosmological density parameter, $\Omega_0$, is subcritical.~\cite{opencase} In the absence of a cosmological constant, this would imply a hyperbolic spatial geometry for the universe (commonly referred to as the `open' universe in cosmological literature). The topologically trivial (simply connected) hyperbolic 3-space, $\hm$, is non-compact and has infinite size. There are numerous theoretical motivations, however, to favor a spatially compact universe.~\cite{cct_motive} To reconcile this with a flat or hyperbolic geometry, consideration of models with non-trivial topology is required. A compact cosmological model is constructed by identifying points on the standard infinite flat or hyperbolic FRW space under the action of a suitable discrete subgroup, $\Gamma$, of the full isometry group, $G$, of the FRW space. The FRW spatial hypersurface is the {\em universal cover}, tiled by copies of the compact space, ${\cal M}$. Any point ${{\bf x}}$ of the compact space has an image ${{\bf x}}_i = \gamma_i {{\bf x}}$ in each tile on the universal cover, where $\gamma_i \in \Gamma$. \begin{figure}[b] \plottwo{vplot_sim2.eps}{m_all_plot_sim2.eps} \caption{The figure consists of two columns of CMB sky-maps showing a pair of $140^\circ$ diameter hemispherical caps each, centered on the South (SGP) and North (NGP) Galactic Poles, respectively. The map labeled DATA, shows the {\sc cobe--dmr} 53+90+31 GHz A+B data after Wiener filtering assuming a standard CDM model, normalized to {\sc cobe}. The rest of the five maps are one random realization of the CMB anisotropy in two examples of compact hyperbolic (CH) spaces for several values of $\Omega_0$ based on our theoretical calculations of $C(\hat q,\hat q^\prime)$ convolved with the {\sc cobe--dmr} beam. Both surface and integrated (ISW) Sachs-Wolfe effects have been included in $C(\hat q,\hat q^\prime)$. No noise was added. The power was normalized to best match the {\sc cobe} data. The theoretical sky was optimally filtered using the {\sc cobe} experimental noise to facilitate comparison with data. The maps labeled L(arge)CH refer to the CH model $v3543(2,3)$. The right column shows the CMB maps for the S(mall)CH model $m004(-5,1)$. (The model number associated with the topology corresponds to that of the census of CH spaces in the Geometry center, Univ. of Minnesota; SCH is one of the smallest and LCH is one of the largest spaces in the census). LCH with $\Omega_0=0.8$ is compatible with the data with a suitable choice of orientation while all the others are ruled out (See Table~\ref{table1}). For all six maps, the average, dipole and quadrupole of the $\vert b \vert > 20^\circ$ sky were also removed and a $20^\circ$ Galactic latitude cut was used, with extra cuts to remove known regions of Galactic emission proposed by the {\sc cobe} team accounting for the ragged edges. The contours are linearly spaced at $15~\mu {\rm K}$ steps. The maps have been smoothed by a $1.66^\circ$ Gaussian filter.\hfill\mbox{}} \label{fig1} \end{figure} For Gaussian perturbations, the angular correlation function, $C(\hat q,\hat q^\prime)$, of the CMB temperature fluctuations in two directions $\hat q$ and $\hat q^\prime$ in the sky completely encodes the CMB anisotropy predictions of a model. The dominant contribution to the anisotropy in the CMB temperature measured with wide-angle beam ($\theta_{\sc fwhm} \gta 2^\circ~ \Omega_0^{1/2}$) comes from the cosmological metric perturbations through the Sachs-Wolfe effect. The angular correlation function of the CMB anisotropy, $C(\hat q,\hat q^\prime)$, depends on the spatial two point correlation function, $\xi_\Phi \equiv \langle\Phi({\bf x},\taurec)\Phi({\bf x^\prime},\taurec)\rangle $ of the gravitational potential, $\Phi$, on the hypersurface of last scattering.\footnote{ Other effects which contribute to the CMB anisotropy at smaller angular scales can also be approximated in terms of spatial correlation of quantities defined on the hypersurface of last scattering.~\cite{us}} To calculate the spatial correlation function on a compact hyperbolic (CH) manifold, described by the corresponding $\Gamma$, we have developed a general technique -- the {\em method of images}, which evades the difficult problem \footnote{The correlation function is usually computed using a mode function expansion. However, obtaining closed form expressions for eigenfunctions of the Laplacian may not be possible beyond the simplest topologies and even numerical estimation is known to be difficult in CH spaces.} of solving for eigenfunctions of the Laplacian on these manifolds.~\cite{us_texas} Using the method of images, the spatial correlation function, $\xi^c_\Phi$, between two points ${\bf x}$ and ${\bf x^\prime}$ on a compact space of volume $V_{\cal M}$ can be expressed as~\cite{us} \begin{equation} \xic = \lim_{r_*\to\infty} \sum_{r_j < r_*} \xiurj - \frac{4\pi}{V_{\cal M}} \int_0^{r_*} dr ~\sinh^2r ~\xiur\,,~~~~~ r_j = d({\bf x}, \gamma_j {\bf x}^\prime)\, . \lbl{moi_final} \end{equation} a regularized sum over the correlation function, $\xi_\Phi^u$, on the universal cover evaluated between ${\bf x}$ and images $\gamma_i{\bf x^\prime}$ of ${\bf x^\prime}$. Numerically it suffices to evaluate the above expression up to $r_*$ a few times the curvature radius, $d_c$, to reach a convergent result. We then integrate $\xi^c_\Phi$ along photon trajectories to get $C(\hat q,\hat q^\prime)$ which includes both surface and integrated (ISW) Sachs-Wolfe effects. The CMB photons can be viewed as propagating to the observer from a $2$-sphere of radius, $R_H$, -- the sphere of last scattering (SLS). In contrast to the topologically trivial models, widely separated pixels in compact spaces can still have strong correlations if in the sum over images, eq.~(\ref{moi_final}), one of the images of ${\bf x^\prime}$ happens to be close to ${\bf x}$. The dependence of the spatial correlations on the anisotropic distribution of images leads to a characteristic statistical anisotropy in the CMB in compact models. These effects are pronounced when the compact space fits well within the SLS, but persist at an observable level even when $V_{\cal M}$ is comparable to (or somewhat bigger than) $V_{\sc sls}$, the volume of SLS. If SLS does not fit completely inside a single tile -- a copy of the compact space, the CMB temperature values will be identical along pairs of circles if temperature fluctuations are dominated by the surface terms at the SLS.~\cite{circles} This pattern of matched circles is one specific manifestation of the angular patterns in $C(\hat q,\hat q^\prime)$. Figure~\ref{fig1} compares theoretical realizations of the CMB anisotropy in a selection of CH models with the {\sc cobe--dmr} data. In this work, the six {\sc cobe--dmr} four-year maps~\cite{dmr4} are first compressed into a (A+B)(31+53+90 GHz) weighted-sum map, with the customized Galactic cut advocated by the {\sc dmr} team. There is no effective loss of information when we do further data compression by using $5.2^\circ\times5.2^\circ$ pixels.~\cite{bdmr294} The theory and data maps have been postprocessed so as to facilitate a fair visual comparison. The incompatibility of models with small $V_{\cal M}/V_{\sc sls}$ (SCH-$\Omega_0=0.3,0.6$) is visually obvious: the best fit amplitudes are high which is reflected in the steeper hot and cold features. Although, SCH-$\Omega_0=0.9$ and LCH-$\Omega_0=0.6$ do not appear grossly inconsistent, the intrinsic anisotropic correlation pattern is at odds with the data. We have carried out a full Bayesian analysis of the probability of the CH models given the {\sc cobe--dmr} 4yr data. In Table~\ref{table1} we present the {\em relative likelihood} of the selected models to that of the infinite, $\hm$, model with the same $\Omega_0$. (The {\sc cobe} data alone does not strongly differentiate between the infinite hyperbolic models with different $\Omega_0$.) The anisotropy of the theoretical $C(\hat q,\hat q^\prime)$ causes the likelihood of compact models to vary significantly with the orientation of the space with respect to the sky, depending on how closely the features in the single data realization available match (or mismatch) the pattern in $C(\hat q,\hat q^\prime)$. Some optimal orientations may also have the ``ugly'' correlation features hidden in the Galactic cut. We analyzed $24$ different orientations for each of our models and found that only the model with $V_{\cal M} > V_{\sc sls}$ (LCH-$\Omega_0=0.8$) cannot be excluded (at one orientation this model is even preferable to standard CDM; this raises a question of the statistical significance of any detection of intrinsic anisotropy of a space when only a single realization of data is available). Similar conclusions were reached by some of the authors (JRB, DP and I. Sokolov~\cite{us_torus}) for flat toroidal models. Comparison of the full angular correlation with {\sc cobe} data led to a much stronger limit on the compactness of the universe than limits from other methods.~\cite{tor_refs} The main result of the analysis was that $V_{\sc sls}/V_{\!\cal M} < 0.4$ at $95\%~CL$ for the equal-sided $3$-torus. For non compact $1$-torus, the constraint on the most compact dimension is not quite as strong. In summary, our results demonstrate that the {\sc cobe} data can put strong constraints on the compact models of the universe. If the universe is small compared to the `horizon' size, correlations appear in the maps that are irreconcilable with the large angle {\sc cobe--dmr} data. \vspace*{-2mm} \begin{table}[htb] \caption{The Log-likelihoods of the compact hyperbolic models relative to the infinite models with same $\Omega_0$ are listed below. The likelihoods are calculated by comparison with {\sc cobe--dmr} data. The three columns of Log-likelihood ratios correspond to the best, second best and worst values that we have obtained amongst $24$ different rotations of the compact space relative to the sky. The number in brackets gives a convenient, albeit crude, translation to gaussian likelihood. Only the last model can be reconciled with the {\sc cobe--dmr} data.\hfill\mbox{}\label{logprob_tab}} \vspace{0.3cm} \begin{center} \begin{tabular}{|c|c|c|c|c|c|} \hline Topology &$\Omega_0$&$V_{\sc sls}/V_{\!\cal M}$&\mco{3}{|c|}{Relative Log. Likelihood (Gaussian approx.)}\\ \cline{4-6} & & &\mco{3}{|c|}{Orientation}\\ & & &`best'&`second best'&`worst'\\ \hline $m004(-5,1)$ &0.3&153.4 & -35.5 (8.4$\sigma$)& -35.7 (8.4$\sigma$) & -57.9 (10.8$\sigma$)\\ &0.6&19.3 & -22.9 (6.8$\sigma$)& -23.3 (6.8$\sigma$) & -49.4 ( 9.9$\sigma$)\\ $V_{\cal M}/d_c^3=0.98$ &0.9&1.2 & -4.4 (3.0$\sigma$)& -8.5 (4.1$\sigma$) & -37.4 ( 8.6$\sigma$)\\ \hline $v3543(2,3)$ & 0.6 &2.9 & -3.6 (2.7$\sigma$) & -5.6 (3.3$\sigma$)& -31.0 (7.9$\sigma$) \\ $V_{\cal M}/d_c^3=6.45$ & 0.8 &0.6 & 2.5 (2.2$\sigma$) & -0.8 (1.3$\sigma$) & -12.6 (5.0$\sigma$) \\ \hline \end{tabular} \end{center} \lbl{table1} \end{table} \vspace*{-6mm} | 98 | 4 | astro-ph9804042_arXiv.txt |
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9804 | astro-ph9804274_arXiv.txt | We search for angular correlation of gamma-ray bursts with cataloged quasars, BL Lac objects, and AGN using a large sample of relatively well-localized bursts detected by WATCH on board GRANAT and EURECA, IPN, and BATSE (327 bursts total). A statistically significant (99.99\% confidence) correlation between GRB and $M_B<-21$ AGN in the redshift range $0.1<z<0.32$ is found. The correlation with AGN is detected, with a lower significance, in three independent GRB datasets. The correlation amplitude implies that, depending on the AGN catalog completeness, 10\% to 100\% of bursts with peak fluxes in the range $3-30\times10^{-6}$~erg~s$^{-1}$~cm$^{-2}$ in the 100--500 keV band are physically related to AGN. The established distance scale corresponds to the energy release of order $10^{52}$~ergs per burst. | Gamma-ray bursts are distributed isotropically on the sky; their peak flux distribution shows the lack of faint bursts compared to the expectation for a homogeneous distribution in Euclidean space (see e.g.\ Fishman \& Meegan 1995 for review). These observed properties of GRB are reproduced by two popular models: 1) GRB arise in an extended Galactic halo with a core radius of $\sim 100$~kpc and 2) GRB are located at cosmological distances. The Galactic halo models are challenged by the observed isotropy of the burst positions. The isotropy constraints much improved recently, and as a result, most variants of the halo model are no longer viable (Briggs et al.\ 1996). An additional argument against the Galactic halo model comes from {\em Einstein} data. The {\em Einstein} IPC sensitivity is sufficient to detect bursts from the halos of nearby galaxies, but they were not found in the data (Hamilton et al.\ 1996). In the cosmological model, the burst isotropy and departures of $\log N - \log S$ from the Euclidean $S^{-3/2}$ law are explained naturally. The minimum redshift $z=0.835$ of the optical transient associated with GRB~970528 (Metzger \etal\ 1997) is a decisive evidence in favor of the cosmological model. Optical spectroscopy of the GRB counterparts is complicated by the faintness of associated optical transients and their relatively featureless spectrum (e.g., no strong emission lines were observed in the GRB~970528 optical transient). Therefore, indirect estimates of the GRB distance scale are still useful. The shape of the GRB peak flux distribution implies that sources of the dimmest BATSE bursts are at $z\sim1$, assuming no source evolution (Emslie \& Horack 1994; Fishman \& Meegan 1995 and references therein). For some plausible evolution of GRB volume density, the above distance estimate can vary by a factor of $\gtrsim 2$ (Horack et al.\ 1995). Using a different approach, Quashnock (1996) derived $z>0.25$ for the dimmest BATSE bursts by cross-correlating the third BATSE catalog with the known large scale structure at low redshifts. Some earlier studies already searched for a direct relationship between GRB and other astrophysical objects at cosmological distances. Kolatt \& Piran (1996) and Marani et al.\ (1997) found that GRB from the third BATSE catalog are correlated with Abell clusters. However, using more precise GRB localizations, Burenin et al.\ (1997), Hurley et al.\ (1997), and Gorosabel \& Castro-Tirado (1997) have not found any correlation with Abell clusters. Furthermore, the amplitude of correlation of well-localized GRB and $z<0.1$ Abell clusters is lower than expected in the case of the same spatial distribution of GRB and optically luminous matter (Burenin et al.\ 1997). If this result is interpreted as an indication that bright and well-localized GRB are at greater distances than nearby Abell clusters, the dimmest BATSE bursts should be at $z>0.3$, in agreement with Quashnock's (1996) results. At still higher redshifts, $z\sim0.1-1$, a natural choice is to search for correlation of GRB with quasars and other flavors of active galactic nuclei, which comprise the majority of known objects at these high redshifts. To search for correlation of AGN and GRB is also attractive because some theoretical models relate bursts to physical processes in active nuclei (e.g.\ Lejter 1980, Carter 1992). Some attempts to search for such a correlation have been undertaken earlier. A marginal evidence for excess of QSO in the small Interplanetary Network (IPN) GRB error boxes has been found (Vrba \etal\ 1995). However, Webber \etal\ (1995) and Gorosabel \etal\ (1995) found that the number of QSO and AGN in GRB error boxes is consistent with random. These earlier analyses were based on a smaller number of bursts with good localizations than available at present. Citing a recent work, Schartel \etal\ (1997) have found a correlation of well-localized GRB from the third BATSE catalog with radio-quiet QSO. The strongest correlation was detected at the $>99.7\%$ confidence level for intrinsically bright QSO with $z<1$. The goal of this work is to search for a correlation of GRB with QSO and AGN using all available data for good, with $\lesssim 1^\circ$ uncertainty, bursts localizations. We use $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0=0$. \begin{figure*}[htb] \begin{minipage}[t]{0.47\textwidth} \epsfysize=1.1\textwidth \centerline{\epsffile{fig1.ps}} \vskip -5pt \small {\bf Fig.~1}--- Number density of QSO from VCV96 as a function of redshift. Smooth solid line represents the true number density of these objects (accounting for their cosmological evolution, Franceschini \etal\ 1994). \end{minipage} \hfill \begin{minipage}[t]{0.47\textwidth} \epsfysize=1.1\textwidth \centerline{\epsffile{fig2.ps}} \vskip -5pt \small {\bf Fig.~2}--- Number density of AGN from VCV96 as a function of their redshift. Dotted histogram represents number density of all AGN, the solid one corresponds to bright $M_B<-21$ AGN. \end{minipage} \end{figure*} | This is not the first attempt to find optical counterparts of GRB by analyzing the optical content of small-area localizations, so the comparison of our results with some previous works is in order. Vrba et al.\ (1995) have performed extensive optical photometry of 8 small-area ($<70$~arcmin$^2$) IPN localizations searching for objects with unusual colors, variability, and proper motions. Only blue objects, which Vrba et al.\ interpreted as QSO, showed a marginal excess at a rate approximately one per localization. This is in agreement with our results. Webber et al.\ (1995) and Gorosabel et al.\ (1995) used several tens GRB localizations from IPN and WATCH, respectively. The optical content was taken from existing catalogs, similar to our study. No excess of either AGN or any other class of objects was found. These results do not directly contradict to ours because Webber et al.\ (1995) and Gorosabel et al.\ (1995) used a smaller number of bursts, did not account for the incompleteness of optical catalogs (\S2.2), and did not apply redshift constraints in search for correlation. Since the completeness of the VCV96 catalog (and probably other all-sky catalogs) is below, or on the level of, several percent (\S4), analyses based on several tens of bursts are inconclusive. The most similar to ours is the work of Schartel et al.\ (1997) who cross-correlated BATSE bursts with the AGN and QSO from VCV96. They found a marginal correlation with QSO, and no correlation with AGN, with no redshift constraints applied to AGN. We essentially reproduce these results (Fig.~4 and Table~1). A significant correlation with AGN found in our analysis is not found by Schartel et al.\ because they used a smaller number of bursts, with poorer localizations (only BATSE data), and used the entire VCV96 catalog which is very incomplete for AGN even at low redshift (\S2.2). Finally, we mention the two GRB with small-area localizations, in which AGN were found. Drinkwater et al.\ (1997) report that the possible X-ray counterpart of GRB~920501 is associated with a Seyfert~1 galaxy at $z=0.315$. The first X-ray localization of a gamma-ray burst by BeppoSAX (3$^\prime$ radius) contains a $z=1.038$ QSO (Piro et al.\ 1998). On the other hand, optical transients associated with other BeppoSAX bursts, GRB~970508, 970228, and 971214 probably are not AGN. Our analysis supersedes most of earlier searches for GRB-AGN association because we used a large GRB dataset, carefully accounted for incompleteness of the optical catalogs, and introduced sensible object redshift constraints. These advantages made it possible to find the strongest ever evidence for association of GRB with a known class of extragalactic objects. However, there are several problems with our analysis. First, we used relatively large area gamma-ray burst localizations, and therefore had to perform a statistical correlation rather than an object-by-object identification. Second, we used a sparse catalog of AGN and QSO, in which many objects were found in a number of small-area high-sensitivity searches. In fact, this introduces a possibility that the correlation we detect is not with AGN but rather with objects around which the AGN we searched, e.g.\ normal galaxies or other QSO (data from Arp 1980 and Monk et al.\ 1988), targets of {\em Einstein}\/ pointings (EMSS AGN, Stocke et al.\ 1991). However, the sky distribution of GRB with $0.1<z<0.32$ AGN inside the localization area, does not generally follow the regions of deep AGN surveys (Fig.~7). Also, a clear redshift dependence of the correlation is not easily explainable in such a scenario. Another problem is that the amplitude of the detected correlation implies that the fraction of bursts related to AGN is somewhat higher than the estimated completeness of the optical catalog (\S4). We can offer two possible explanations. First is that GRB prefer luminous AGN, for which the catalog completeness is higher. Second is that GRB prefer a certain type of AGN (Seyfert 1 or 2, X-ray loud or quiet, radio loud or quiet, etc.) which is more commonly present in the VCV96 catalog than the ``average'' AGN. Despite these problems, we believe that the case for an association of bright gamma-ray bursts with AGN at moderate redshift is compelling. This case can be further proved, or disproved, by an extensive optical, preferably spectroscopic, survey of small-area localizations of bright GRB, similar to the work of Vrba et al.\ (1995) but using a larger number of bursts. Another approach would be to use a more complete catalog of AGN covering a significant fraction of the sky; unfortunately, this seems impractical until the Sloan Digitized Sky Survey is completed. | 98 | 4 | astro-ph9804274_arXiv.txt |
9804 | astro-ph9804104_arXiv.txt | Recently, \citet{MParkJGott97a} claimed that there is a statistically significant, strong, negative correlation between the image separation $\Delta\theta$ and source redshift $\zs$ for gravitational lenses. This is somewhat puzzling if one believes in a flat ($k=0$) universe, since in this case the typical image separation is expected to be independent of the source redshift, while one expects a negative correlation in a $k=-1$ universe and a positive one in a $k=+1$ universe. \citeauthor{MParkJGott97a} explored several effects which could cause the observed correlation, but no combination of these can explain the observations with a realistic scenario. Here, I explore this test further in three ways. First, I show that in an inhomogeneous universe a negative correlation is expected regardless of the value of $k$. Second, I test whether the \thz\ can be used as a test to determine $\lnull$ and $\onull$, rather than just the sign of $k$. Third, I compare the results of the test from the \citeauthor{MParkJGott97a} sample to those using other samples of gravitational lenses, which can illuminate (unknown) selection effects and probe the usefulness of the \thz\ as a cosmological test. | Historically, there has been little interest in the \thz\ compared to other cosmological tests based on gravitational lensing statistics, perhaps because the inflationary paradigm \citep[\eg][]{AGuth81a}, which began about the same time as the discovery of the first gravitational lens \citep*{DWalshCW79a}, has become so influential. Since a flat ($k=0$) universe is a robust prediction of inflation, many researchers assume this and consider only flat universes (or, at most, $k=-1$ cosmological models with $\lnull=0$). Due to the fact that for the popular singular isothermal sphere model for a single-galaxy lens the average image separation $\Delta\theta$, integrated over the lens redshift $\zd$ from $\zd=0$ to $\zd=\zs$, is {\em completely independent\/} of the source redshift $\zs$ in a flat universe, there is little point in pursuing the \thz\ if one is interested primarily in flat cosmological models. If one is not committed to a flat universe, then of course one should not assume $k=0$, but even if one believes that the universe must be flat, it is still important to test this belief observationally. The situation is somewhat worsened by the fact that most `standard' cosmological tests such as the \mbox{$m$-$z$} (magnitude-redshift or `standard candle') and $\theta$-$z$ (angular size-redshift or `standard rod') relations, `conventional' gravitational lensing statistics, age of the universe) are relatively insensitive to the radius of curvature of the universe ($R_{0} \sim (|\Omega_{0}+\lambda_{0}-1|)^{-\frac{1}{2}}$), being degenerate in combinations of $\lnull$ and $\onull$ in directions roughly perpendicular to lines of constant $R_{0}$ in the $\lnull$-$\onull$ plane. A notable exception are constraints derived from CMB anisotropies \citep*[\eg][]{DScottSW95a,WHuSS97a}. | \citet{MParkJGott97a} pointed out that the image separations in gravitational lens systems show a strong significant negative correlation with the source redshift, while in a flat universe one would expect no correlation (while a negative correlation would be expected in a universe with negative curvature and a positive one in a universe of positive curvature). None of the possibilities they examined were strong enough to explain the effect. A possibility not examined by them, namely an inhomogeneous universe, produces a negative correlation regardless of the sign of the curvature, but it too is not strong enough to account for the effect. As a general test for the values of $\lnull$ and $\onull$ the test is of no use, all cosmological models being assigned roughly the same probability, but {\em which\/} value they are assigned depends on the sample used. The strong dependence of the result on the sample used seems to indicate that the result of \citet{MParkJGott97a} is due not to some physical cause but rather to unidentified selection effects in the sample of gravitational lenses taken from the literature. The large number of JVAS and CLASS lenses gives us an independent comparison sample, thus demonstrating the need for discovering a large number of lenses in a well-defined sample. As \citet{MParkJGott97a} point out, since many conclusions based on `conventional' gravitational lensing statistics are based on essentially the same lenses as in their literature sample, if this sample is for some unknown reason atypical, then conclusions drawn from statistical analyses of it must be examined with care. It will thus be interesting to see what conclusions can be drawn from a statistical analysis of the JVAS/CLASS sample after the observational tasks have been completed. (We expect to find more lenses, but have no qualms about using the present incomplete sample in this analysis since there is no reason to believe that a larger sample would show a different \thz.) | 98 | 4 | astro-ph9804104_arXiv.txt |
9804 | astro-ph9804038_arXiv.txt | We use numerical simulations of structure formation in a Cold Dark Matter model to predict the absorption lines in the soft X-rays produced by heavy elements in the shock-heated intergalactic medium at low redshift. The simulation incorporates a model for heavy element production in galaxies and the subsequent dispersion of the metals to the intergalactic medium. We analyze in particular absorption lines produced by oxygen, and calculate the ionization stage taking into account the observed X-ray background at the present time. We find that oxygen is fully ionized by the X-ray background in low-density voids, and is mostly in the form of $\OVII$ and $\OVIII$ in the sheets and filamentary regions. Strong absorption lines of $\OVII$ and $\OVIII$ with equivalent widths $W\sim 100 \kms$ are produced in filamentary regions of overdensities $\sim 100$ and temperatures $\sim 10^6$ K, located in the outskirts of groups and clusters of galaxies. The $\OVII$ line at $E=574$ eV is generally the strongest one in these systems. Our model predicts that any X-ray source (such as a quasar) should typically show about one $\OVII$ absorption line with $W > 100 \kms$ in the interval from $z=0$ to $z=0.3$. These lines could be detected with the upcoming generation of X-ray telescopes, and their origin in intervening systems could be confirmed by the association with groups of galaxies and X-ray emitting halos near the line-of-sight at the same redshift. The hot intergalactic medium may be one of the main reservoirs of baryons in the present universe, and the heavy element X-ray absorption lines offer a promising possibility of detecting this new component in the near future. | Hierarchical theories of the formation of large-scale structure are based on the hypothesis that the gravitational collapse of some type of cold dark matter (i.e., any collisionless matter with a small enough velocity dispersion to allow collapse down to scales much smaller than the observed structure), triggered by initial density fluctuations, is responsible for the formation of galaxies and for the later assembly of groups, clusters, and superclusters of galaxies. This generic scenario has been very successful in explaining a large variety of observations, even though the origin of the initial density fluctuations remains unknown, and simple models are therefore assumed for the primordial power spectrum (generally based on adiabatic, Gaussian, scale-invariant fluctuations). One of the predictions that have been made from these models, with the use of numerical hydrodynamic simulations, is that despite the wide range of scales over which the baryonic matter is able to collapse (as dark matter halos successively merge from the onset of non-linearity until the present time), a relatively large fraction of baryons should still remain as ``intergalactic matter''. This intergalactic medium should be forming a network of shock-heated gas in the form of filamentary and sheet-like structures connected to galaxy clusters and groups, as well as colder gas left out in voids, as is revealed in numerical simulations (e.g., Ostriker \& Cen 1996). A similar network of photoionized and shock-heated gas should also be present at high redshift (although at lower temperatures than at the present time, due to the lower velocities of collapse at high redshift), and probably gives rise to the hydrogen \lya forest (e.g., Cen \etal 1994, Hernquist \etal 1996, Miralda-Escud\'e \etal 1996, Zhang \etal 1997, 1998). Here, a new prediction based on the same theory of hierarchical clustering and the presence of intergalactic gas shall be studied. It has already been shown observationally that heavy elements are present in the high-redshift absorption systems (Songaila \& Cowie 1996 and references therein). It is very likely that many more elements were spread to the intergalactic medium (hereafter, IGM) at later times, either through galactic winds energized by supernova explosions or active galactic nuclei (e.g., Dekel \& Silk 1986), or simply by the process of gravitational merging, which can also lead to some gas being ejected from halos back into the IGM (Gnedin \& Ostriker 1997). As pointed out by Shapiro \& Bahcall (1980) and Aldcroft \etal (1994), heavy elements in the intergalactic medium should cause absorption lines (as well as continuum edges of absorption) on background X-ray sources In fact, for highly ionized, hot gas at $T\sim 10^6\kelvin$, absorption lines from heavy elements are probably the only method of detection, since hydrogen is highly ionized and its absorption lines are very weak. The soft X-ray emission from low-density gas is also very weak and would generally be seen superposed with emission from other structures along the line-of-sight, as well as the Galactic emission. We shall denote these absorption lines by ``X-ray forest'', in analogy to the \lya forest caused by hydrogen. The X-ray forest is much more difficult to observe than the \lya forest, due to the lower sensitivity and resolution in the X-ray band. The possibility of observing this X-ray forest has also been discussed by Aldcroft \etal (1994) and Canizares \& Fang (1998). Similar absorption lines may also be detectable in the ultraviolet when the temperature is not very high; these are generally caused by lithium-like ions like $\OVI$. Mulchaey \etal (1996) proposed this as a method to detect halos of hot has in poor groups of spiral galaxies, where the temperature of the halo gas may be too low to have been detected in emission by ROSAT. There is in fact evidence for a population of $\OVI$ absorbers among the numerous absorption systems seen in quasar spectra that may arise in hot, collisionally ionized gas (e.g., Burles \& Tytler 1996). In this paper we shall predict some properties of the X-ray forest from a hydrodynamic simulation, and discuss the prospects for detecting it in future X-ray missions. | We have presented in this paper a new generic prediction of the large-scale structure theories based on hierarchical gravitational collapse of density fluctuations: X-ray quasars should show absorption lines in their spectra from heavy elements in the intervening intergalactic gas. A substantial fraction of the baryons are predicted to be in the form of low-density intergalactic gas at the present time (Ostriker \& Cen 1996; Cen \& Ostriker 1998; Miralda-Escud\'e \etal 1996; Zhang \etal 1998). Several mechanisms are known to enrich this gas with heavy elements; given the observations of the metallicity in the centers of rich clusters, it seems inevitable that this enrichment has also taken place in the lower-density gas. In fact, the simulations presented here {\it underpredict} the metal abundance observed in clusters; had we increased the metal injection rate to fit the present cluster metallicities, the column densities of the absorbers we predict would be proportionally increased. Absorption lines in the soft X-rays, as well as far-UV lines produced by lithium-like ions (Mulchaey \etal 1996), are the only observational means we know of to detect low-density, hot intergalactic gas in the present universe. At temperatures $\sim 10^6$ K, the neutral hydrogen fraction is too low to produce significant \lya absorption, so we must rely upon the highly ionized heavy elements. The search for intergalactic gas is important to complete an inventory of the baryons in the present universe. The total baryon density observed in galaxies and X-ray emitting gas in clusters comes close to the baryon density predicted by nucleosynthesis (e.g., Persic \& Salucci 1997; Fukugita, Hogan, \& Peebles 1998), but given the uncertainties it is possible for up to $\sim$ 80\% of the baryons to be undetected, in the form of ionized intergalactic gas. Results from cosmological simulations also suggest that the amount of baryons in the \lya forest at $z\sim 3$ is higher than the sum of all the baryons observed at present (Rauch \etal 1997; Weinberg \etal 1998; Zhang \etal 1998). The detection and subsequent study of the X-ray forest would lead to a large number of potential applications. The number of absorption lines found will be measuring the product of the baryon density times the metallicity in the IGM. The ratio of strengths of the $\OVII$ and $\OVIII$ lines will probe the distribution of the gas temperature and density, and several other lines may be observed which could provide extensive tests for the predictions of large-scale structure models. The X-ray absorption lines could be correlated with structures seen in emission near the line-of-sight, such as galaxies or X-ray emitting clusters and groups. A new era in the study of the intergalactic medium could be opened with the discovery of the X-ray forest by the new X-ray observatories. | 98 | 4 | astro-ph9804038_arXiv.txt |
9804 | hep-ph9804205_arXiv.txt | s{One of the fundamental problems of modern cosmology is to explain the origin of all the matter and radiation in the Universe today. The inflationary model predicts that the oscillations of the scalar field at the end of inflation will convert the coherent energy density of the inflaton into a large number of particles, responsible for the present entropy of the Universe. The transition from the inflationary era to the radiation era was originally called reheating, and we now understand that it may consist of three different stages: preheating, in which the homogeneous inflaton field decays coherently into bosonic waves (scalars and/or vectors) with large occupation numbers; backreaction and rescattering, in which different energy bands get mixed; and finally decoherence and thermalization, in which those waves break up into particles that thermalize and acquire a black body spectrum at a certain temperature. These three stages are non-perturbative, non-linear and out of equilibrium, and we are just beginning to understand them. In this talk I will concentrate on the preheating part, putting emphasis on the differences between preheating in chaotic and in hybrid inflation. } | At the end of inflation all the energy density is in the homogeneous zero mode of the inflaton field. The Universe is in a vacuum-like state with zero temperature and vanishing particle and entropy densities. The problem of reheating is how to convert all this coherent energy into a state of thermalized relativistic particles. The original analysis~\cite{book} assumed the perturbative decay of the inflaton into bosons and fermions, as if the inflaton were already an ensemble of decoherent {\em particles}. Reheating ended when the total decay rate of the inflaton was of the order of the expansion rate of the Universe, $\Gamma\sim H$, while the total energy of the inflaton field decayed exponentially fast into other particles. As a consequence, the final reheating temperature only depended upon $\Gamma$. We understand today that, for certain parameter ranges, there is a new decay channel that is non-perturbative,\cite{KLS1} due to the coherent oscillations of the inflaton field, which induces stimulated emission of bosonic~\footnote{Fermions can also be parametrically amplified, but their occupation numbers are constrained, $n_k\leq1$, by Pauli's exclusion principle.} particles into energy bands with large occupation numbers. The modes in these bands can be understood as Bose condensates, and they behave like classical waves. The backreaction of these modes on the homogeneous inflaton field and the rescattering among themselves produce a state that is far from thermal equilibrium and may induce very interesting phenomena, such as non-thermal phase transitions~\cite{nonthermal} with production of topological defects, a stochastic background of gravitational waves,\cite{GW} production of heavy particles in a state far from equilibrium, which may help GUT baryogenesis~\cite{baryo} or constitute today the dark matter in our Universe.\cite{DM} These classical waves eventually reach a state of turbulence where, hopefully, decoherence will occur and thermalization will follow, although these stages are not yet fully understood, either analytically or numerically. The period in which particles are produced via parametric resonance is called {\em preheating}.\cite{KLS1} The idea is relatively simple, the oscillations of the inflaton field induce mixing of positive and negative frequencies in the quantum state of the field it couples to. In the language of quantum fields in curved space, creation and annihilation operators mix via Bogoliubov transformations,\footnote{For an introduction to particle production in strong external fields, see Grib et al.\cite{GMM}} $a_k = \alpha_k \bar a_k + \beta^\ast \bar a_{-k}^\dagger$, and with every oscillation of the inflaton field, new particles are produced, $n_k \equiv \langle\bar 0|a_k^\dagger a_k|\bar 0\rangle = |\beta_k|^2$. In the case of chaotic inflation, with a massive inflaton $\phi$ coupled to a massless scalar field $\chi$, the evolution equation for the Fourier modes, $\ddot X_k + \omega_k^2 X_k=0$, with $X_k=a^{3/2}(t)\chi_k$ and $\omega_k^2 = k^2/a^2(t) + g^2\phi^2(t)$, can be cast in the form of a Mathieu equation, with coefficients $A=k^2/4a^2m^2+2q$ and $q=g^2\Phi^2/4m^2$, where $\Phi$ is the amplitude and $m$ is the frequency of inflaton oscillations, $\phi(t)=\Phi(t)\sin mt$. For certain values of the parameters $(A,q)$ there are exact solutions that grow exponentially with time, and each mode $k$ belongs to an instability band of the Mathieu equation.\footnote{For a recent comprehensive review on preheating after chaotic inflation see Kofman~\cite{Kofman} and references therein.} These instabilities can be interpreted as coherent ``particle'' production with large occupation numbers. One way of understanding this phenomenon is to consider the energy of these modes as that of a harmonic oscillator, $E_k = |\dot X_k|^2/2 + \omega_k^2 |X_k|^2/2 = \hbar \omega_k (n_k + 1/2)$. The occupation number of level $k$ can grow exponentially fast, $n_k\sim\exp(2\mu_k mt)\gg1$, and these modes soon behave like classical waves. It is analogous to the well-known mechanism of generation of density perturbations during inflation.\cite{book} The parameter $q$ during preheating determines the strength of the resonance. It is possible that the model parameters are such that parametric resonance does {\em not} occur, and then the usual perturbative approach would follow, with decay rate $\Gamma$. In fact, as the Universe expands, the growth of the scale factor and the decrease of the amplitude of inflaton oscillations shifts the values of $(A,q)$ along the stability/instability chart of the Mathieu equation, going from broad resonance, for $q\gg1$, to narrow resonance, $q\ll1$, and finally to the perturbative decay of the inflaton. Parametric resonance will stop whenever the inflaton decay is dominated by the perturbative decay, $q m < \Gamma$, or when the instability modes are redshifted away from the (last) narrow resonance band, $q^2 m < H$. | 98 | 4 | hep-ph9804205_arXiv.txt |
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9804 | hep-ph9804425_arXiv.txt | During preheating after inflation, parametric resonance rapidly generates very large fluctuations of scalar fields. In models where the inflaton field $\phi$ oscillates in a double-well potential and interacts with another scalar field $X$, fluctuations of $X$ can keep the $\phi\to-\phi$ symmetry temporarily restored. If the coupling of $\phi$ to $X$ is much stronger than the inflaton self-coupling, the subsequent symmetry breaking is a first-order phase transition. We demonstrate the existence of this nonthermal phase transition with lattice simulations of the full nonlinear dynamics of the interacting fields. In particular, we observe nucleation of an expanding bubble. | 98 | 4 | hep-ph9804425_arXiv.txt |
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9804 | astro-ph9804326_arXiv.txt | { We study the galaxy number density (smoothed on a $5h^{-1}$~Mpc scale) around 18 low-redshift Lyman-alpha absorbers previously observed with {\sl HST}. The absorbers lie in the foregrounds of Mrk 335, Mrk 421, Mrk 501, I Zw 1, and 3C 273, all within regions where there are now complete redshift surveys to $m_{\rm Zw}=15.5$. We construct a smoothed galaxy number density field from the redshift survey data and determine the distribution of densities at the Lyman-alpha absorber locations. We also find the distribution of galaxy number density for a variety of test samples: all galaxy locations within the Center for Astrophysics Redshift Survey (CfA2), CfA2 galaxy locations along randomly selected lines of sight, and randomly chosen redshifts along random lines of sight. The Lyman-alpha absorbers are present in dense regions of the survey, but occur far more frequently in underdense regions than do typical luminous galaxies. The distribution of smoothed galaxy density around the Lyman-alpha absorbers is inconsistent at the 4$\sigma$ level with the density distribution around survey galaxies. It is highly consistent with a density distribution at randomly chosen redshifts along random lines of sight. This supports earlier evidence that the nearby, low column density ($\log N_{H\sc i} \lesssim 14$) Lyman-alpha forest systems are spatially distributed at random; they are not well correlated with the local large-scale structure. } | 98 | 4 | astro-ph9804326_arXiv.txt |
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9804 | astro-ph9804110_arXiv.txt | We show that the luminosity dependence of the red clump stars on age and metallicity can cause a difference of up to $\lesssim$0.6 mag in the mean absolute I magnitude of the red clump between different stellar populations. We show that this effect may resolve the apparent $\approx$0.4 mag discrepancy between red clump-derived distance moduli to the Magellanic Clouds and those from, e.g., Cepheid variables. Taking into account the population effects on red clump luminosity, we determine a distance modulus to the LMC of 18.36 $\pm$ 0.17 mag, and to the SMC of 18.82 $\pm$ 0.20 mag. Our alternate red clump LMC distance is consistent with the value (m-M)$_{LMC}$ = 18.50 $\pm$ 0.10 adopted by the HST Cepheid Key Project. We briefly examine model predictions of red clump luminosity, and find that variations in helium abundance and core mass could bring the Clouds closer by some 0.10--0.15 mag, but not by the $\approx$0.4 mag that would result from setting the mean absolute I-magnitude of the Cloud red clumps equal to the that of the Solar neighborhood red clump. | The distance to the Magellanic Clouds is a problem of great astrophysical interest due to their role in the determination of extragalactic distances (c.f. \cite{mad98}). Despite the large amount of effort that has gone into the determination of these distances, they remain a matter of some controversy (see \cite{wes97}, \cite{wes90} for thorough discussions). The success of the HST Cepheid Key Project, whose goal is the determination of H$_0$ to an accuracy of 10\%, depends critically on knowledge of the distance to the Clouds, especially the LMC. Based on the Cepheid variable Period-Luminosity relation (\cite{mad98}) and the light echoes of SN1987A (\cite{pan97}), the Key Project has adopted a distance modulus of (m-M)$_{LMC}$ = 18.50 $\pm$ 0.10 (50 $\pm$ 2 kpc) (\cite{raw97}). These determinations are in good agreement with recent derivations from RR Lyrae variables, based on both ground-based, statistical parallax methods (\cite{fea97b}), and on the {\sl Hipparcos}-calibrated distance scale (\cite{gra97}, \cite{rei97}). However, the LMC distance controversy is far from settled. A substantial fraction of recent techniques have yielded smaller values, typically in the range of (m-M)$_{LMC}$ $\approx$ 18.30 $\pm$ 0.10. These include further analysis of the SN1987A light echoes (\cite{gou98}) and recalibration of the RR Lyrae magnitude-metallicity relation (\cite{lay96}). If the distance to the LMC is uncertain, that of the SMC is even more so (\cite{wes97}, \cite{uda98}). Because of its smaller Cepheid population and its large line-of-sight depth, distance determinations to the SMC are in generally poorly constrained. We can only say with confidence that it lies some 0.3--0.6 mag beyond the LMC. The best Cepheid distance to the SMC is 18.94 $\pm$ 0.04 (\cite{lan94}). Recently, the large photometric sample of the OGLE microlensing survey and its resulting high-quality color-magnitude diagrams of the Clouds have permitted the development of the red clump of intermediate-age helium-burning stars as a ``standard candle'' for single-step distance determinations (\cite{pac98}). The key to this method has been the availability of accurate {\sl Hipparcos} parallaxes to calibrate the absolute magnitude of solar neighborhood red clump stars. The {\sl Hipparcos} color-magnitude diagrams (e.g., \cite{jim98}) show a solar-neighborhood red clump that has a very small dispersion in mean absolute I-band magnitude. Because the red clump is the dominant post-main-sequence evolutionary phase for most stars, it makes a tempting target for the application of ``standard candle'' techniques for distance determinations. The red clump method was developed very thoroughly and applied to fields in the Galactic bulge (\cite{pac98}) and M31 (\cite{sta98}). These studies used the mean absolute magnitude of solar-neighborhood red clump stars, M$_I^0$ = $-$0.23 $\pm$ 0.03, obtained from a volume-limited sample of 228 red clump stars observed with {\sl Hipparcos}. The red clump method is based on a well-populated, well-calibrated phase of stellar evolution; in fact, it is possibly a more reliable distance indicator than many other methods that have been employed. In a recent paper, \cite{uda98} extended the red clump method to the Magellanic Clouds. After taking careful account of the reddening distributions along Cloud lines of sight, they find mean red clump magnitudes of I$_0$(LMC) = 17.85 $\pm$ 0.03, and I$_0$(SMC) = 18.33 $\pm$ 0.03. Using the solar neighborhood value M$_I^0$ = $-$0.23 $\pm$ 0.03, \cite{uda98} find distance moduli of (m-M)$_{LMC}$ = 18.08 $\pm$ 0.15 and (m-M)$_{SMC}$ = 18.56 $\pm$ 0.09. These values are $\approx$0.4 mag below the ``long'' distance scale preferred by the HST Cepheid Key Project, and only marginally consistent with the ``short'' scale. \cite{sta98b} applied the same technique to an independent large photometric survey of the LMC and obtained the virtually identical result (m-M)$_{LMC}$ = 18.07 $\pm$ 0.12. | Our alternate red clump distance modulus (m-M)$_{LMC}$ = 18.36 $\pm$ 0.17 is $\approx$0.3 mag longer than that obtained under the assumption that the LMC red clump mimics the Galaxy's. This is consistent with the ``long'' distance modulus of 18.50 $\pm$ 0.10 adopted by the Cepheid Key Project, and, contrary to \cite{uda98} and \cite{sta98b}, in agreement with the most recent calibrations of the Cepheid period-luminosity relation, which give (m-M)$_{LMC}$ = 18.44 $\pm$ 0.35 or 18.57 $\pm$ 0.11 (\cite{mad98}). Our value of 18.36 $\pm$ 0.17 is also in agreement with the shorter distance from RR Lyrae stars of (m-M)$_{LMC}$ = 18.28 $\pm$ 0.13 (\cite{lay96}). Taking account of the stellar population of the SMC pushes that galaxy slightly farther away as well, from (m-M)$_{SMC}$ = 18.56 $\pm$ 0.09 derived by \cite{uda98}, to (m-M)$_{SMC}$ = 18.82 $\pm$ 0.20. The larger value is in good agreement with that from the Cepheid period-luminosity relation, (m-M)$_{SMC}$ = 18.94 $\pm$ 0.04 (\cite{lan94}). The detailed stellar physics of mass-loss, and the relations between Z, Y, and M$_c$ introduce significant uncertainty into the determination of red clump absolute magnitudes, even in the I band. It is quite likely that these effects may work to bring the Magellanic Clouds to a distance more consistent with the ``short'' (RR Lyrae) distance scale than the ``long'' (Cepheid) scale. However, the red clump method does not require Magellanic Cloud distances as much as 15\% smaller than commonly accepted. Figure 1 shows the menagerie of recent LMC distance determinations with errorbars, and it can be seen that the red clump method gives results consistent with most of the other determinations. We conclude that the red clump is indeed an extremely useful distance indicator, as described by (\cite{pac98}, \cite{sta98}, and \cite{uda98}). However, like most stellar ``standard candles'', its properties vary with the composition and age of the host galaxy, and the assumption that all red clumps are identical to the {\sl Hipparcos} red clump is probably incorrect. Among populations dominated by stars older than $\approx$6 Gyr, the standard candle approximation should be valid to a high degree for M$_I^0$. For younger populations, M$_I^0$ is probably brighter than M$_I^0$(local), and the correction may amount to as much as several tenths of a magnitude. | 98 | 4 | astro-ph9804110_arXiv.txt |
9804 | astro-ph9804260_arXiv.txt | If the first stars formed soon after decoupling of baryons from the thermal cosmic background radiation (the CBR) the radiation may have been last scattered in a cloudy plasma. We discuss the resulting small-scale anisotropy of the CBR in the limit where the plasma clouds are small compared to the mean distance between clouds along a line of sight. This complements the perturbative analysis valid for mildly nonlinear departures from homogeneity at last scattering. We conclude that reasonable choices for the cloud parameters imply CBR anisotropy consistent with the present experimental limits, in agreement with the perturbative approach. This means the remarkable isotropy of the CBR need not contradict the early small-scale structure formation predicted in some cosmogonies. | \label{sec-intro} Cosmogonies in which the baryons were concentrated in clouds at the epoch of last scattering of the thermal cosmic background radiation (the CBR), as in isocurvature models (\cite{pe94}, \cite{pe97a}), are in qualitative agreement with the appearance in quasar absorption line spectra of a well-advanced state of structure formation at redshift $z\sim 5$, but the large amplitude of the density fluctuations may produce significant small-scale anisotropy in the CBR. The analysis of the CBR anisotropy in second order perturbation theory was introduced by \cite{os86}. Vishniac's (1987) more detailed investigation has been confirmed by \cite{hu94}, \cite{do95}, and \cite{hu97}. The analysis and numerical evaluation are extended in \cite{pe95} and Persi et al. (1995). The application of perturbation theory may be uncertain if early structure formation produced highly nonlinear departures from homogeneity at the epoch of last scattering, however. To investigate this we have developed a nonperturbative model for scattering in a cloudy distribution of plasma. The expression for the CBR anisotropy $\delta T/T$ in the strongly cloudy limit bears a close resemblance to the Ostriker-Vishniac relation, and the numerical results for $\delta T/T$ accordingly are similar. Our analysis, which extends previous discussions by \cite{ho89}, \cite{pe90}, \cite{ag96} and \cite{gr98}, assumes the CBR was last scattered by free electrons in clouds with density contrast well above unity, so the mean free distance $t_f$ between intersections of gas clouds along a line of sight is much larger than the typical cloud size $\dcl$. The simplification offered by this limit is that the details of the matter distribution and motion on scales from $\dcl$ to $t_f$ are not important, because a line of sight on average samples only one cloud over the distance $t_f$. Thus we can model the clouds as an inhomogeneous random Poisson process. In the process the joint distribution of cloud motions and scattering optical depths as a function of position along the line of sight is determined by the plasma density and velocity fields averaged through a window of width $t_f$. If $t_f$ is larger than the scale of non-linear density fluctuations then the CBR anisotropy $\delta T/T$ is the sum in quadrature of a perturbative contribution and a shot noise term. The next section shows the relation between the shot noise term in $\delta T/T$ and the Ostriker-Vishniac (1986) effect. The model of the clouds as an inhomogeneous random Poisson process is presented in \S 3, and a simplified treatment of the effect of correlated cloud motions is discussed in \S 4. In \S 5 we present numerical examples of the expected CBR anisotropy. | \label{discussion} The simplifying assumption for this analysis is that structure formation is so well advanced at the epoch of last scattering of the CBR by free electrons that the mean distance $t_f$ between intersections of clouds along a line of sight is large compared to the scale of nonlinear mass fluctuations. This allows us to model the clouds as an inhomogeneous random Poisson process determined by the mass density and peculiar velocity fields smoothed through a window of width $t_f$, and it leads to the shot noise contribution to the small-scale CBR anisotropy in equation~(\ref{eqn-unsaturated}). This approach is motivated by the isocurvature CDM model for structure formation (\cite{pe97a}), in which structure formation could commence at decoupling at redshift $z\sim 1000$. A second important motivation has been to complement the usual perturbative analysis of the effect of the nonlinear growth of small-scale structure. The similarity of results from the perturbative (\cite{psco95}) and nonperturbative approaches leads us to believe we have reliable methods for estimating the effect of early nonlinear structure formation on the CBR anisotropy. The observations of young galaxies and the intergalactic medium at $z\sim 3$ indicate a situation intermediate between the perturbative and nonperturbative cases. The damped Lyman-$\alpha$ systems contain a significant baryon fraction, and the mean distance between intersections of these clouds is large (at $z=3$ it is longer than the Hubble length). There also is a significant baryon fraction in the Lyman-$\alpha$ forest, and these clouds have a relatively short mean free distance. It is not unreasonable to speculate that the situation at much larger redshifts similarly calls for a combination of the two approaches to the analysis of the angular distribution of the CBR. The CBR anisotropy produced by the Sunyaev-Zel'dovich (1970) effect of the hot electrons in clusters of galaxies, which certainly is dominated by the shot noise term, offers an important constraint on the epoch of collection of the intracluster plasma. The evidence from the analysis of Persi et al. (1995) is that this constraint does not yet rule out the early structure formation picture. And our conclusion from the numerical examples in \S 5 is that within presently known observational constraints structure formation could have commenced when the universe was optically thick to scattering of the CBR. | 98 | 4 | astro-ph9804260_arXiv.txt |
9804 | astro-ph9804056_arXiv.txt | We present the results of soft X-ray observations of the intermediate-age open cluster IC\,4651 performed with the ROSAT PSPC. We detected 25 sources. Two are identified with a giant binary and a blue straggler respectively, both belonging to the cluster, and two with probable main-sequence members. Two other cases have ambiguous identification. Of the five binaries known in the cluster, the one detected in X rays is the only one whose period is short enough to maintain fast rotation and therefore strong stellar activity at a high age. The detected blue straggler is probably binary, suggesting that binarity is the key to producing a high level of X-ray emission. It is the third blue straggler detected in X rays. The remaining sources need to be identified through optical follow-up. | Stellar activity depends crucially on the star's rotation rate (Pallavicini et al. 1981), which decreases with age because of magnetic braking (e.g. Skumanich 1972). Thus, open clusters are crucial to distinguish between truly evolutionary effects on stellar activity and effects primarily due to the rotation rate itself. The first X-ray observations of open clusters were carried out with the {\it Einstein} satellite (Stern et al. 1981, Caillault \& Helfand 1985, Micela et al. 1988, Schmitt et al. 1990, Micela et al. 1990), but these observations have been carried out in a more systematic way with the ROSAT satellite (Stern et al. 1992; Stauffer et al. 1994; Patten \& Simon 1993; Randich \& Schmitt 1995; Randich et al. 1995, 1996a,b). Various clusters of different ages have been studied in order to understand the evolution of stellar activity with age (see Randich 1997 for a review). Before ROSAT, the attention has been concentrated mainly on young clusters (30 to 700 Myr). ROSAT performed the first observations of the old open cluster M\,67 ($\sim$5\,Gyr, Belloni et al. 1993, Belloni et al. in preparation) and NGC\,188 ($\sim$9\,Gyr, Belloni et al., in preparation), leading to the detection of a number of sources. Indication of chromospheric activity has been found for most of the optical candidates (Pasquini \& Belloni 1994). In contrast to younger clusters, clusters older than $\sim$\,1 Gyr are not expected to contain rapidly-rotating single late-type stars, and therefore strong X-ray sources. However, there are stars older than $\sim$ 1 Gyr that show rapid rotation: these are members of close binary systems, where tidal interaction prevents the stars from losing angular momentum; well-known examples are the RS CVn binaries. The observations of these clusters have led to the detection of such binaries, and also to a number of peculiar and interesting objects, such as white dwarfs, catalclysmic variables, blue stragglers and wide/eccentric binaries. In the framework of a project to cover the X-ray observational gap between old and young clusters, we observed the intermediate age open clusters NGC\,752 ($\sim$2\,Gyr, Belloni \& Verbunt 1996), NGC\,6940 ($\sim$1\,Gyr, Belloni \& Tagliaferri 1997) and IC\,4651 ($\sim$2.5\,Gyr) with ROSAT PSPC. Moreover, ROSAT HRI observations of the clusters NGC 3680 ($\sim$2\,Gyr) and NGC\,2527 ($\sim 1$\,Gyr) were made in 1997 and are currently being analyzed. In the PSPC observation of NGC\,752, 49 X-ray sources have been detected; seven of them are identified with optical cluster members, four of which are short period binaries, one is a rapid rotator and one is a blue straggler (Belloni \& Verbunt 1996). In the PSPC observation of NGC\,6940 18 sources were detected, four of which are identified with members of the cluster with a fifth source a suspected member. In NGC\,6940, a high fraction of the detected members are binaries: three out of four of the identified members are among the only six binaries known in the cluster. These observations give also evidence for the presence of a saturation level, at which the whole surface of the star is chromospherically active (see Belloni 1998 for a review). Here we present the results obtained for IC\,4651. This open cluster has an estimated age of $\sim 2.5$\,Gyr and a distance of $\sim 800 - 900$\,pc (Eggen 1971, Anthony-Twarog et al. 1988). It has a relatively constant reddening, estimated to be E(B-V)$ = 0.15$ by Eggen (1971) and E(B-V)$ = 0.09$ by Anthony-Twarog \& Twarog (1987), and an angular size smaller than $\sim 15'$. The paper is organized as follows: in Sect. 2 we present the PSPC observation and our data analysis, in Sect. 3 we present and discuss the results and in Sect. 4 we compare them with those of other open clusters and discuss the implications. | The age of IC\,4651 is about twice that of NGC\,6940 and similar to that of NGC\,752. As in the case of NGC\,752 (Belloni \& Verbunt 1996), we expect to detect coronal sources only if they are in binary systems; moreover, due to saturation effects one would expect to detect more easily giant stars than main sequence stars (see discussions in Stauffer et al. 1994; Belloni \& Verbunt 1996; Belloni \& Tagliaferri 1997). As can be seen from Fig. 1, in IC\,4651 there are more than a dozen giants. Five of them are found to be binaries by Mermilliod et al. (1995). However, only one of them has a period short enough to sustain a high level of stellar activity, and this is the binary we detect as an active X-ray source. In this scenario, one expects the stars to co-rotate and therefore the orbit to be circular. This binary is slightly eccentric (e=0.1), but its characteristics are compatible with co-rotation (see Verbunt \& Phinney 1995). This is also the brightest source of the six that we suppose to be at the cluster distance, in agreement with the presence of a saturation level, which is higher for evolved stars (see Belloni \& Verbunt 1996). In NGC\,6940 we had not detected main sequence stars, but only giants. Of these, two are not classified as binaries; one in particular has been extensively studied (Mermilliod \& Mayor 1989) and is very unlikely to be a binary (see Belloni \& Tagliaferri 1997). Thus, their detection in the X-ray band is puzzling. On the contrary, in IC~4651 of all known giants only one, a binary, is detected. Two other stars lying on the cluster main sequence are also detected. Given the age of the cluster, we expect them to be binaries. At the bottom end of the cluster main sequence, somewhat offset from it, lie the two possible counterparts of source \#18. They could be either field stars or weak members of the cluster. In the error box of source \#11, which is only 13 arcsec, there are at least seven stars, which look like a sub-cluster. The three brightest of these stars are plotted as crosses in Fig. 1. If these stars are all members of the cluster, then one is a blue straggler (i.e. it lies in a region of the color-magnitude diagram occupied by blue stragglers), while the other two are stars in the process of evolving toward the red giant branch. The possibility cannot be ruled out that more than one of these stars contribute to the detected X-ray emission. Finally, source \#14 is identified with a blue straggler, found by Anthony-Twarog \& Twarog (1987) to be a binary. This is the third blue straggler detected in the X-ray band, the other two being in M67 and NGC\,752 respectively. Of all blue stragglers known in the old and intermediate-age open clusters that we studied, only three are detected in X rays. All three have measurements that indicate binarity (see Belloni 1998), and once again this is probably the key to X-ray emission. These are not the only binaries know among the blue stragglers in these clusters, showing that binarity does not imply strong stellar activity per se. One of the two stars also has to be of late spectral type in order to have an enhanced dynamo activity, like in Algol systems. For the sources marked with `OUT' in Table 1 we have no color information; however, from their magnitudes, most of them should be stars. They lie outside the inner 10 arcmin diameter region of the cluster and are probably non-members. We have already planned follow-up optical observations to determine the physical nature of all detected X-ray sources in the IC\,4651 field. | 98 | 4 | astro-ph9804056_arXiv.txt |
9804 | astro-ph9804146_arXiv.txt | % OH(1720 MHz) masers unaccompanied by 1665/7 MHz line masers have recently been proposed as indicators of the interaction of supernova remnants (SNRs) and molecular clouds. We present a model for the masing region in which water produced in a C-type shock wave driven into the molecular cloud is dissociated as a result of the X-ray flux from the SNR. We note that the magnetic field strengths inferred from Zeeman splitting of the 1720 MHz line measure the internal pressure of the supernova remnant. In addition, we discuss the interaction of Sgr A East, a SNR candidate, with the 50 km/s cloud at the Galactic Centre and present near-infrared observations of \hmol\ emission towards the regions where OH(1720 MHz) maser emission is concentrated. The magnetic field strength obtained from earlier Zeeman measurements is consistent with rough pressure equilibrium between the postshock gas and the X-ray gas filling Sgr A East detected by ASCA. Further, the intensity of the v=1--0 S(1) line of \hmol\ is consistent with the shock strength expected to be driven into the molecular gas by this pressure. The relative intensities of the \hmol\ lines in Sgr A East imply mainly collisional excitation. | OH masers have generally been used as a diagnostic for HII regions and evolved stars. However, a recent study by Frail, Goss \& Slysh (1994) revealed that the 1720 MHz transition of OH maser emission, when unaccompanied by the 1665 and 1667 MHz OH lines, can be an effective indicator of shock waves interacting with molecular clouds, particularly for supernova remnants (SNRs). Frail et al. (1994) detected 26 distinct OH(1720 MHz) maser spots along the interface between the SNR W28 and an adjacent molecular cloud. Evidence of the association between the molecular cloud and W28 comes from the distribution of the molecular material following the eastern edge of the supernova shell (Wootten 1977). There are more than a dozen Galactic sources and three extragalactic sources in which 1720 MHz OH masers have been found interior to SNR's with adjacent molecular clouds (Yusef-Zadeh, Uchida \& Roberts 1995; Frail et al. 1996; Yusef-Zadeh et al. 1996; Green et al. 1997; Seaquist, Frayer \& Frail 1997). These masers are close both in position and velocity to the interfaces between the remnants and the clouds. In addition, in several of these sources, CO emission lines reach a maximum in both brightness and linewidth at the interface between remnant and cloud (Wootten 1977). These are strong observational indications that the shocks are caused by the remnants expanding into their respective adjacent molecular clouds. Further support is provided by theoretical studies of the pumping of the OH maser lines (Elitzur 1976; Pavlakis \& Kylafis 1996a,b). The 1665 and 1667 MHz masers are pumped by far-infrared radiation and are therefore associated with HII regions and evolved stars. The OH(1720 MHz) maser is collisionally pumped in molecular gas at temperatures and densities between 15-200 K and \(10^4-10^6\) cm\(^{-3}\), respectively. Thus in the absence of the 1665/7 MHz transitions, the OH 1720 MHz line presumably traces cooling, shocked gas. However, shock chemistry predicts that OH is \emph{not} abundant in the postshock gas as it is rapidly converted to \water\ within the shock front. OH masers adjacent to compact HII regions are produced by photodissociation of \water\ (Elitzur \& de Jong 1978; Hartquist \& Sternberg 1991; Hartquist et al. 1995), but in that case there is a strong dissociating FUV flux from the star. The dissociating flux is largely absorbed and reradiated in the FIR by grains, providing a sufficient IR background to also pump the 1665/7 MHz transitions (e.g. Pavlakis \& Kylafis 1996b). This cannot be the case for the unaccompanied 1720 MHz masers associated with SNR-molecular cloud interactions. Here we propose that it is the weak X-ray flux from the SNR interior that is ultimately responsible for the dissociation of \water\ in the shocked molecular gas. We note that if the OH maser arises in the postshock gas, then Zeeman measurements determine the magnetic field strength in the postshock gas, and thus measure the pressure within the SNR more directly than other methods. The observations of the IC 443, W28 and W44 are consistent with this scenario. Finally, we apply these ideas to the interaction of the Galactic center nonthermal source Sgr A East, which is either a SNR or a multiple-SNR driven bubble, with a molecular cloud at the Galactic center. | We have proposed a model for the OH 1720 MHz masers unaccompanied by main-line transitions that are associated with SNRs interacting with molecular clouds (Frail et al. 1994, 1996; Green et al. 1997). A C-type shock wave driven into the adjacent cloud produces water that is subsequently dissociated by the secondary FUV flux produced by the interaction of X-rays from the SNR incident on the molecular cloud. The dissociation only becomes significant once the shocked molecular gas has cooled to about \( 180 \u K \), naturally producing OH at the temperature and abundance required for collisional pumping of the 1720 MHz transition. Some consistency checks suggest that this model also applies to the OH 1720 MHz masers seen towards Sgr A East. Firstly, within the framework of the model the magnetic field strengths of roughly 3 mG inferred from Zeeman splitting of the OH masers (Yusef-Zadeh et al. 1996) are measurements of the magnetic pressure in the postshock gas, which dominates the gas pressure behind the shock front. Equating the magnetic pressure, \( 4\ee -7 \u erg \percc \), to \( \rho v_s^2 \), where \( \rho \) is the preshock density and \( v_s \) is the shock speed, and adopting a preshock density of \( \nh=2 \ee 4 \percc \) (Mezger et al. 1989), we infer \( v_s \approx 25\)--\(30 \kms \). A C-type shock at this speed produces an intensity in the 1-0 S(1) line of \( 10^{-4} \)--\(10^{-3} \u erg \ut s -1 \ut cm -2 \) (Draine et al. 1983; Kaufman \& Neufeld 1996). The \emph{measured} intensity, after correcting for extinction assuming that \( A_K \approx 2.5 \), is \( 1.5\ee -4 \u erg \ut s -1 \ut cm -2 \). Second, the inferred postshock pressure is comparable to the pressure of the X-ray emitting gas filling Sgr A East detected with ASCA (Koyama et al. 1996), which has \( n_e \approx 6 \percc \) and \( T \approx 10 \u keV \), and a pressure of \( 2\ee -7 \u erg \percc \). Finally, we note that the X-ray luminosity from Sgr A East is \( 10^{36} \u erg \persec \) (Koyama et al. 1996), providing the required dissociating flux for the production of OH behind the shock front. Rho (1995) suggests that centrally-peaked thermal X-ray emission from SNRs is characteristic of SNR-molecular cloud interactions. Green et al. (1997) note that all of the SNRs detected in the OH (1720 MHz) line that have been observed in X-rays are members of this class. Our model appears to strengthen this link by relying on the X-ray flux from the SNR to produce OH, but in principle a cosmic-ray flux enhanced by a factor of 100 over the solar neighbourhood value has a similar effect on the heating and chemistry of molecular gas (see, e.g. Maloney et al. 1996). The high-energy cosmic-ray flux in the W44, W28 and IC 443 SNRs inferred from EGRET observations is roughly two orders of magnitude larger than the solar neighbourhood value (Esposito et al. 1996), so a naive extrapolation down to the cosmic-ray energies important for ionisation suggests that the cosmic-ray and X-ray contributions to the secondary FUV flux may be similar. In any case, this picture supports the suggestion by Frail et al. (1994) that the presence of the OH 1720 MHz line, and the absence of the 1665/1667 MHz lines provide a clear diagnostic of shocked molecular gas. | 98 | 4 | astro-ph9804146_arXiv.txt |
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