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0710 | 0710.1077_arXiv.txt | In this work we give special attention to the bimetric theory of gravitation with massive gravitons proposed by Visser in 1998. In his theory, a prior background metric is necessary to take in account the massive term. Although in the great part of the astrophysical studies the Minkowski metric is the best choice to the background metric, it is not possible to consider this metric in cosmology. In order to keep the Minkowski metric as background in this case, we suggest an interpretation of the energy-momentum conservation in Visser's theory, which is in accordance with the equivalence principle and recovers naturally the special relativity in the absence of gravitational sources. Although we do not present a general proof of our hypothesis we show its validity in the simple case of a plane and dust-dominated universe, in which the `massive term' appears like an extra contribution for the energy density. | \label{intro} Could the graviton have a non-zero rest mass? The observations have shown that this is a possibility. One of the most accurate bounding on the mass of the graviton comes from the observations of the planetary motion in the solar system. Variations on the third Kepler law comparing the orbits of Earth and Mars can lead us to $m_g < 7.8 \times 10^{-55}g$ \cite{Talmadge88}. Another bound comes from the analysis of galaxy clusters that lead to $m_g < 2 \times 10^{-62}g$ \cite{Goldhaber74} which is considerably more restrictive but less robust due to the uncertainties in the content of the universe in large scales. Studying rotation curves of galactic disks, \cite{JC} has found that we should have a massive graviton of $m_g \ll 10^{-59}g$ in order to obtain a galactic disk with a scale length of $b\backsim10$ kpc. The above tests are obtained from static fields based on deviations of the newtonian gravity. In the weak field limit has been proposed \cite{Finn2002} to constraint $m_g$ using data on the orbital decay of binary pulsars. From the binary pulsar PSR B1913+16 (Hulse-Taylor pulsar) and PSR B1534+12 it is found the limit $m_g < 1.4 \times 10^{-52}g$, which is weaker than the bounds in static field. It is worth recalling that the mass term introduced via a Pauli-Fierz (PF) term in the linearized approximation produces a theory whose predictions do not reduce to those of general relativity for $m_g \rightarrow 0$. This is the so called van Dam Veltmann Zakharov discontinuity \cite{Veltman1970}. Moreover the Minkowski space as background metric is unstable for the PF theory \cite{Gruzinov2005}. However, there is no reason to prefer the PF term over any other non-PF quadratic terms. It is important to emphasize that these mass terms do not have clear extrapolation to strong fields. A way to do that was proposed by Visser \cite{visser98}. To generalize the theory to strong fields, Visser makes use of two metrics, the dynamical metric ($g_{\mu\nu}$) and a non-dynamical background metric ($\left( g_0\right) _{\mu\nu}$) that are connected by the mass term. Although adding a prior geometry is not in accordance with the usual foundations underlying Einstein gravity, it keeps intact the principles of equivalence (at least in its weak form) and general covariance in the Visser's work. Some interesting physical features emerge from the theory such as extra states of polarizations of the gravitational waves \cite{wayne2004}. In the present article, we explore some aspects which are not treated by Visser in his original paper. In the great part of the astrophysical studies the Minkowski metric is the most appropriate choice to the background metric. However, in the study of cosmology, it is not possible to consider this kind of metric, and we need some prior considerations regarding a background metric. Once this problem emerges from the coupling of the two metrics and the energy's conservation condition, we analyze an alternative interpretation of this condition. We also show that this interpretation is in accordance with the equivalence principle and recovers naturally the special relativity in the absence of gravitational sources. Arguments in favor of a Minkowskian background metric in Visser's theory are also considered. This paper is organized as follows: in section \ref{sec:1} we show how to introduce a mass for the graviton through a non-PF term. We present the strong field extrapolation as given by Visser in section \ref{sec:2}. In section \ref{sec:3} we show that the theory is not in accordance with a Minkowski background metric in the study of cosmology. In section \ref{sec:4} we re-interpret the stress-energy conservation in order to keep Minkowski as background in any case. In particular, we show that our re-interpretation is in accordance with the equivalence principle. In section \ref{sec:5} we show why Minkowski is the most natural choice to the background metric. We briefly study some cosmological consequences of our interpretation of the energy-momentum conservation in section \ref{sec:6}. And finally, we present our conclusions in the last section. | \label{sec:7} Our interpretation of the energy-momentum conservation in the Visser's massive gravity is in accordance to the equivalence principle and recover naturally the results of special relativity in the absence of gravitational sources. The point of view considered in this paper allow us to consider Minkowski as background metric in Visser's theory in all astrophysical cases including cosmology. This new interpretation may lead to interesting cosmological results once we can construct a cosmological model in a theory with massive gravitons with a Minkowski background. Additional contributions to the cosmological fluids will appear due to the modifications in the interaction potential, which, maybe, would be a way of treat the dark-energy problem. The analyses of the theory in the absence of gravitational sources lead us to exclude the de Sitter space-time as a vacuum solution of the massive gravity, once a constant $\Lambda$ term is rigorously zero in a flat background. Another interesting feature is that our interpretation of the energy conservation in strong fields is independent of the form of the tensor which interact with the perfect-fluid tensor, so this can be used to other models with additional energy-momentum contribution. | 7 | 10 | 0710.1077 |
0710 | 0710.0247_arXiv.txt | We present high spatial resolution spectroscopic measurements of dynamic fibrils (DFs) in the Ca~{\small{II}}~8662~{\AA} line. These data show clear Doppler shifts in the identified DFs, which demonstrates that at least a subset of DFs are actual mass motions in the chromosphere. A statistical analysis of 26 DFs reveals a strong and statistically significant correlation between the maximal velocity and the deceleration. The range of the velocities and the decelerations are substantially lower, about a factor two, in our spectroscopic observations compared to the earlier results based on proper motion in narrow band images. There are fundamental differences in the different observational methods; when DFs are observed spectroscopically the measured Doppler shifts are a result of the atmospheric velocity, weighted with the response function to velocity over an extended height. When the proper motion of DFs is observed in narrow band images, the movement of the top of the DF is observed. This point is sharply defined because of the high contrast between the DF and the surroundings. The observational differences between the two methods are examined by several numerical experiments using both numerical simulations and a time series of narrow band H$\alpha$ images. With basis in the simulations we conclude that the lower maximal velocity is explained by the low formation height of the Ca~IR~line. We conclude that the present observations support the earlier result that DFs are driven by magneto-acoustic shocks exited by convective flows and p-modes. | The dynamical nature of the chromosphere is obvious when the Sun is imaged in the line center of strong chromospheric spectral lines, most commonly in the H$\alpha$ line \citep[e.g.,][]{2006Noort}. One of the dominating features is a vast number of thin (0.2-1 Mm) omnipresent jet like structures \citep{1968Beckers}. On the quiet solar limb they are commonly known as spicules, while on the quiet disk they are often called mottles, and finally in active regions they are known as active region fibrils or dynamic fibrils (DFs). The nomenclature can be confusing, but there are strong indications that these structures are physically closely related \citep{1994Tsiropoula,1995Suematsu,2001Christo,2007Rouppe}. Recently, a combination of high--resolution observations and advanced numerical modeling have shown that DFs are most likely driven by shocks that form when photospheric oscillations leak into the chromosphere along inclined flux tubes \citep{1990Suematsu,2004dePontieu,2006Hansteen,2007dePontieu}. The inclination of the magnetic field lowers the acoustic cutoff frequency sufficiently to allow p--modes with the dominant low frequencies to propagate along flux tubes \citep{1973Mich,1977Bel}. These insights into the formation of DFs have become possible because of recent developments in observational techniques, such as bigger telescopes combined with real time wavefront corrections by adaptive optics (AO) systems \citep[e.g.][]{2000Rimmele,SSTAO}, and post-processing methods \citep[e.g.][]{speckle,momfbd} which have made observations of these jet structures much more reliable. These developments have spurred several authors to focus on the detailed understanding of DFs \citep[e.g.,][]{2003dePontieu,2004dePontieu, 2005dePontieu,2007dePontieu,2004Kostas,2006Hansteen,2006deWijn,2007Julius, 2007Lars}. One of the important results of this work is that the DFs are driven by and can channel photospheric oscillations into the chromosphere and the corona. \begin{figure*}[!ht] \includegraphics[width=\textwidth]{f1.eps} \caption{ Spectrograms of the H$\alpha$ line (panel a) and the Ca~{\small{II}}~8662~{\AA} line (panel b). Notice the highly dynamical line center of the Ca line. The corresponding MFBD processed slitjaw image (panel c) and the narrow band H$\alpha$ DOT image (panel d). The position of the spectrograph's slit is marked with a line in the DOT image. } \label{plotone} \end{figure*} In two papers \citet{2006Hansteen} and \citet{2007dePontieu} used high spatial and high cadence observations of the H$\alpha$ line center together with realistic simulations to investigate the nature of DFs. One of their conclusions was that the DFs follow parabolic paths along their axis with decelerations lower than the solar gravitational deceleration. Previous observations did not allow an accurate determination of the nature of the trajectory because of lower quality data and line--of--sight (LOS) effects that were difficult to estimate \citep{1988Nishikawa,1995Suematsu}. \citet{2006Hansteen} and \citet{2007dePontieu} further report regional differences between DFs observed in two different plage areas. The regional differences are explained by different inclination angles of the magnetic fields in the two regions. In this way the magnetic topology of the solar atmosphere works as a filter, where only waves with certain periods can leak through. The simulations, spanning from the upper convection zone to the corona, reproduce the observed correlations between the maximum velocities and decelerations in DFs leading to the conclusion that DFs are formed by chromospheric shocks driven by global p-modes and convective flows. In this paper we add to the understanding of the DFs, by analyzing high spatial and high cadence spectrograms of the Ca {\small{II}} 8662 {\AA} line, put into context by simultaneous H$\alpha$ spectrograms and narrow-band images. Furthermore, the observational results are compared with the numerical simulations of \citet{2006Hansteen}. In \S~\ref{Obs} we describe the observing program and instrumentation. The data reduction method is described in \S~\ref{Data}. In \S~\ref{Obsres} we show the results of the observations. The main observational errors are discussed in \S~\ref{error}. To get a better understanding of the observations we present several numerical experiments in \S~\ref{Sim}. Finally, we summarize the results in \S~\ref{Con}. | \label{Con} In this work we have presented co-spatial and co-temporal narrow band H$\alpha$ images and spectroscopic measurements of the Ca~{\small{II}}~8662~{\AA} line. These observations have been used to identify 26 DFs, and measure their Doppler shifts. A reduced--$\chi^2$ analysis shows that the time evolution of the Doppler shifts are well approximated by a linear fit, if the measurement errors are about $2$~km\,s${}^{-1}$. Using this approximation we derive values for the decelerations and maximal velocities for each DF. Scatter plots of the deceleration and maximal velocity show a strong positive correlation between the two. We also observe weak correlations between the deceleration and lifetime and the maximal velocity and lifetime. These results are supporting the shock-wave theory as explanation model for the DFs \citep{2006Hansteen, 2007dePontieu,2007Lars}. Furthermore, the Doppler shifts show that at least a subset of DFs are caused by mass moving up and down in the atmosphere. The values of the maximum velocity and decelerations are all somewhat lower than earlier reported values. Using numerical experiments we have explained the differences in the two observational sets with the intrinsic differences in observational methods. Earlier observations have used the high contrast seen between the top of the DF and the background for measuring the proper motion of the DF. This high contrast is caused by the intensity increase due to contributions from the transition region. In the present observations the DF motion is measured using Doppler shifts, which are affected by the atmospheric conditions over the formation height. The formation height of the Ca~{\small{II}}~8662~{\AA}~line is much lower than the transition region. Since the shock amplitude is increasing from the formation height for the Ca IR line to the transition region we necessarily measure lower velocities using spectroscopy. The difference in maximal velocities derived from the two methods in our simulations is about a factor two, which is about the same as observed. The one dimensional nature of the slit spectrograph somewhat affects our results, but experiments with narrow band images show that this is not altering the results significantly. | 7 | 10 | 0710.0247 |
0710 | 0710.5712_arXiv.txt | The cosmological concordance model contains two separate constituents which interact only gravitationally with themselves and everything else, the dark matter and the dark energy. In the standard dark energy models, the dark matter makes up some 20\% of the total energy budget today, while the dark energy is responsible for about 75\%. Here we show that these numbers are only robust for specific dark energy models and that in general we cannot measure the abundance of the dark constituents separately without making strong assumptions. | We cosmologists are very proud that our field has finally reached the status of ``precision science'' over the last decade. The quality of the current observations of the cosmic microwave background (CMB), the galaxy distribution and the luminosity distance to type Ia supernovae (SN-Ia) is indeed impressive, and has allowed the construction of a concordance model in which the universe contains the known, baryonic, matter (5\% of the energy density today), radiation (negligible energy density today), dark matter (20\%) and dark energy (75\%). The need for dark matter became apparent long ago in order to explain the motion of galaxies in clusters \cite{darkmat} and the observed galaxy rotation curves. In cosmology, it is often modelled as a pressureless fluid with negligible interactions. Much more recently, less than ten years ago, new SN-Ia data \cite{sn1a} convinced the majority of cosmologists that dark energy was needed as well. Until today the nature of the dark energy is a deep mystery. Although many models have been proposed, there are none that can explain its current abundance in a natural way. The alternative to the model building is a more phenomenological approach, where one measures the physical properties of the dark energy. To this end, one introduces a completely general fluid and tries to determine its characteristics from observations. | We have seen that cosmology cannot measure separately the properties of the dark matter and of a general dark energy component. In order to do that, we either need to impose additional assumptions, for example that the dark energy is a scalar field, or else we need a non-gravitational measurement of the dark matter properties, specifically of its contribution to the total energy density of the universe. One possibility is a detection of supersymmetry at LHC, which may in turn determine the abundance and mass of the lightest stable SUSY particle, one of the best candidates for the dark matter. As a corollary, if the abundance determined in this way is not the one expected within the $\Lambda$CDM cosmological concordance model, one possible explanation is an evolving dark energy. We can also consider the degeneracy as a test of the generality of the different approaches to measure the dark energy equation of state. Since no analysis so far seems to have found it, we can only wonder what else has been overlooked. Finally, although we show here that one never can prove experimentally from cosmological data alone that the dark energy is a cosmological constant, it is remarkable that a model containing just cold dark matter and $\Lambda$ fits the data so well. From a model selection point of view $\Lambda$CDM is still the preferred model because of its simplicity. \ack It is a pleasure to thank Luca Amendola, Ruth Durrer, Domenico Sapone and Anze Slosar for interesting discussions. MK acknowledges funding by the Swiss NSF. | 7 | 10 | 0710.5712 |
0710 | 0710.0137_arXiv.txt | We present multiple epochs of H$\alpha$ spectroscopy for 47 members of the open cluster NGC 3766 to investigate the long term variability of its Be stars. Sixteen of the stars in this sample are Be stars, including one new discovery. Of these, we observe an unprecedented 11 Be stars that undergo disk appearances and/or near disappearances in our H$\alpha$ spectra, making this the most variable population of Be stars known to date. NGC 3766 is therefore an excellent location to study the formation mechanism of Be star disks. From blue optical spectra of 38 cluster members and existing Str\"omgren photometry of the cluster, we also measure rotational velocities, effective temperatures, and polar surface gravities to investigate the physical and evolutionary factors that may contribute to the Be phenomenon. Our analysis also provides improvements to the reddening and distance of NGC 3766, and we find $E(B-V) = 0.22 \pm 0.03$ and $(V-M_{\rm V})_0 = 11.6 \pm 0.2$, respectively. The Be stars are not associated with a particular stage of main-sequence evolution, but they are a population of rapidly rotating stars with a velocity distribution generally consistent with rotation at $70-80$\% of the critical velocity, although systematic effects probably underestimate the true rotational velocities so that the rotation is much closer to critical. Our measurements of the changing disk sizes are consistent with the idea that transitory, nonradial pulsations contribute to the formation of these highly variable disks. | \setcounter{footnote}{3} NGC 3766 is a rich, young open cluster in the Carina spiral arm that is well known for its high content of Be stars \citep{slettebak1985}, and many previous studies of this cluster have focused on the characteristics of these stars to identify their evolutionary status. The cluster has been the target of numerous photometric studies \citep{ahmed1962, yilmaz1976, shobbrook1985, shobbrook1987, moitinho1997, piatti1998, tadross2001, mcswain2005b}. But despite these intensive investigations, the cluster's age and distance remain somewhat uncertain; measurements of its age range from 14.5 to 25 Myr (WEBDA\footnote{The WEBDA database is maintained by E.\ Paunzen and is available online at http://www.univie.ac.at/webda/navigation.html.}; \citealt{lynga1987, moitinho1997, tadross2001}), and its distance is between 1.5 and 2.2 kpc. The reddening $E$($B$-$V$) is between 0.16 and 0.22 (see the discussion of \citealt{moitinho1997}). Spectroscopic investigations of NGC 3766 have targeted a limited sample of cluster members, focusing primarily on the Be star and supergiant populations (\citealt{harris1976}; \citealt{mermilliod1982} and references therein; \citealt{slettebak1985, levesque2005}). Even the eclipsing double-lined spectroscopic binary BF Centauri (= HD 100915), a member of NGC 3766, has been largely neglected by modern spectroscopic observations (\citealt{clausen2007} and references therein). For most cluster members, no detailed information about their physical characteristics such as temperature, gravity, rotation, and metallicity are known. In this work, we present red and blue optical spectra for both normal B-type and Be stars in the cluster. Like many prior studies of NGC 3766, our primary goal is to investigate the Be star population; but unlike other works, we achieve a more complete understanding of this subset of B stars by comparing these emission-line objects to their non-emission counterparts. Therefore we present measurements of the effective temperature, $T_{\rm eff}$, surface gravity, $\log g$, and in most cases the projected rotational velocity, $V \sin i$, for 26 normal B stars and 16 Be stars in NGC 3766. We use these results to improve the known reddening and distance to the cluster. From multiple epochs of H$\alpha$ spectroscopy, we also investigate the variability of the circumstellar disks and estimate the disk mass loss/gain rates for 11 Be stars. Finally, we use the observed disk masses and angular momenta to show that nonradial pulsations are a possible origin for the disks, and they probably fill during short-lived bursts of mass flow from the stellar surface. | Our spectroscopic analysis of NGC 3766 has revealed that Be stars may be much more common than we originally thought. In our photometric study of NGC 3766 \citep{mcswain2005b}, we found up to 13 Be stars (5 definite, 8 uncertain) out of an expected 191 B-type stars, not counting the one Be star that saturated our photometry. The new total of 16 Be stars is 23\% greater. Among these 16 Be stars, 2--5 of them appear to have almost no disk at any given time, and an additional 2--4 have extremely subtle emission in their H$\alpha$ line profile that could easily be mistaken for other phenomena (such as NRP manifesting themselves as bumps moving across the line or SB2 line blending). Therefore 25--50\% of the Be stars may go undetected in a single spectroscopic observation, and photometric snapshots are even less likely to discern such weak emitters. We note four stars (Nos.\ 27, 45, 49, and 77) that were found to be possible or likely Be stars in the photometric study by \citet{shobbrook1985, shobbrook1987}, but they never showed emission during our observations and thus remain unconfirmed. The existence of transitory, weak disks (especially Nos.\ 130 and 196) could mean that many more Be stars are waiting to be discovered. For our total sample of 48 Southern open clusters in our photometric survey, we found a low Be fraction of $2-7$\% \citep{mcswain2005b}. Considering the very weak disks that are observed in NGC 3766 and the exceptionally high variability among the cluster's Be population, the total fraction of Be stars could be much greater. We are currently performing a similar spectroscopic study of several other clusters from our survey, and we will address those results in a future paper. While the Be stars of NGC 3766 are not distinguishable from normal B-type stars by their evolutionary states, they do form a population of rapidly rotating stars. With two exceptions, their measured velocities are consistent with a uniform population of rapid rotators having $V = 0.7-0.8 \; V_{\rm crit}$. Gravitational darkening and weak emission in the \ion{He}{1} lines may mean that these velocities are underestimated by as much as 33\% \citep{townsend2004}, so the true $V_{\rm rot}$ is probably at least $0.84 \; V_{\rm crit}$. From the measured changes in the disks' masses and angular momenta, NRP are a capable source for the mass flow into the equatorial plane. The pulsations may be a transitory phenomenon, however, and the variable nature of the Be stars probably reflects dramatic changes in the surface activity. | 7 | 10 | 0710.0137 |
0710 | 0710.2903_arXiv.txt | We use the kinetic theory of nucleation to explore the properties of dust nucleation in sub-saturated vapors. Due to radiation losses, the sub-critical clusters have a smaller temperature compared to their vapor. This alters the dynamical balance between attachment and detachment of monomers, allowing for stable nucleation of grains in vapors that are sub-saturated for their temperature. We find this effect particularly important at low densities and in the absence of a strong background radiation field. We find new conditions for stable nucleation in the $n-T$ phase diagram. The nucleation in the non-LTE regions is likely to be at much slower rate than in the super-saturated vapors. We evaluate the nucleation rate, warning the reader that it does depend on poorly substantiated properties of the macro-molecules assumed in the computation. On the other hand, the conditions for nucleation depend only on the properties of the large stable grains and are more robust. We finally point out that this mechanism may be relevant in the early universe as an initial dust pollution mechanism, since once the interstellar medium is polluted with dust, mantle growth is likely to be dominant over non-LTE nucleation in the diffuse medium. | Dust particles are one of the fundamental components of the interstellar medium (ISM) and an ever-present worry for observers due to their opacity at optical and UV wavelengths (Cardelli, Clayton \& Mathis 1989). The ISM of the Milky Way is polluted by a mixture of grains made of a variety of materials, likely dominated by carbonaceous grains, silicates, and small PAHs particles (Mathis, Rumpl \& Nordsieck 1977; Weingartner \& Draine 2001). The dust properties are supposed to be the result of dust formation in the outflows of evolved stars (e.g. Salpeter1977; Stein \& Soifer 1983; Mathis 1990; Whittet 1992; Draine 2003, and references therein) and subsequent evolution, and eventual dissolution, in the ISM, mainly as the effect of shock waves that destroy the grains through sputtering (Draine 1989; McKee 1989; Edmunds 2001). Alternative dust production sites are supernova explosions (Kozasa, Hasegawa \& Nomoto 1989, 1991; Todini \& Ferrara 2001; Nozawa et al. 2003; Schneider, Ferrara \& Salvaterra 2004; Bianchi \& Schneider 2007) and quasar outflows (Elvis, Marengo \& Karovska, 2002). The theory of dust nucleation in astrophysics is heavily influenced by the theory of the nucleation of phase transitions in super-saturated vapors (Becker \& Doring 1935; Feder et al. 1966; Abraham 1974). The theory had mild success in reproducing nucleation rates, but is still controversial in many aspects, especially because it extrapolates the properties of macroscopic bodies to clusters of few molecules and because it extends the thermodynamic approach to systems with a handful of particles. To add to these problems, astrophysical dust nucleation requires chemical reactions, since grains of materials that do not have a vapor state do nucleate (think, for example, to the nucleation of olivines from silicon oxides and metals; Draine 1979; Gail \& Sedlmayr 1986). An alternative approach is the so-called kinetic theory, which describes nucleation as the result of attachment and detachment of monomers from a seed cluster of $n$ particles (atoms, molecules or radicals; Nowakowski \& Ruckenstein 1991ab). Both the thermodynamic and the kinetic nucleation theory have been developed in conditions of {\it true equilibrium}, i.e., when the two phases have the same temperature. In the astrophysical scenario, however, the temperature of the dust grains can be sensibly lower than the temperature of the gas in which they are embedded due to efficient radiation cooling (e.g., Draine 1981). This would seem to be irrelevant to nucleation theory, since a vapor needs to be already nucleated in order to have grains that can be colder than the gas phase. Even a sub-saturated vapor, however, has a large number of unstable clusters that form by random association of monomers (and rapidly evaporate). In this paper we study the effect of cooling of these proto-clusters in a sub-saturated vapor, and the effect this has on the balance between attachment and detachment of monomers. Using the kinetic theory of nucleation, we find that even largely sub-saturated vapors can nucleate, provided they are not immersed in a strong radiation field. We compute, albeit under some controversial assumptions, the non-LTE nucleation rate. We show that, even though it is not as large as in super-saturated vapors, it can produce dust grains at a rate that can reproduce the average dust grain density in the Milky Way over a timescale of several million years. In addition we show that non-LTE effects can increase the nucleation rate in super-saturated vapors. Non-LTE nucleation could therefore provide a slow channel for dust formation, in which dust is built over a relative long time in a slowly evolving region. Such an example could be the outflow from AGN nuclei (Elvis et al. 2002). Such evolution is different from the one envisaged in the classical dust factories -- AGB star atmospheres and supernov\ae\ -- where dust nucleation is rapid but short lived since the favorable conditions are rapidly lost. This paper is organized as follows: in \S~2 we briefly review the classical kinetic theory of nucleation; in \S~3 we compute the dust grain temperature and in \S~4 we compute the new conditions for nucleation. In \S~5 we consider the nucleation rate and discuss our results in \S~6. | We have considered the effect of grain (cluster, droplet) cooling in the nucleation of liquid and solid phases in vapors. We find that the effect can be dramatic on the nucleation rate and on the nucleation phase diagram, allowing for nucleation in large regions of the parameter space that are classically considered to be non-nucleating. As is in general true for nucleation, there are several limits and approximations that we should bear in mind when considering the theory from the quantitative point of view. As exemplified by the comparison of the data with the theory in Fig.~\ref{fig:j}, several orders of magnitude can separate the nucleation rate prediction from the observations. The controversial points are: \begin{itemize} \item {\it Sticking coefficients} --- Most of the figures and computations in this work assume $k_s=1$. This is not always true (Batista et al. 2005). In addition to its dependence on temperature, the sticking coefficient may depend on the size of the cluster. A big cluster could more easily absorb the extra kinetic energy of the incoming monomer, compared to a small cluster (K00), and therefore $k_s$ may be significantly smaller than unity for very small clusters. \item {\it Capillary approximation} --- The capillary approximation, i.e., the assumption that the surface tension does not depend on the cluster size, is very controversial, and a change in the surface energy for very small clusters could result in big changes on the nucleation rate. For example, the discrepancy in Fig.~\ref{fig:j} could be solved by assuming a larger surface tension for the very small water droplet. In addition, macromolecules do not even have a properly defined surface, and the whole concept does not apply. Finally, the exponential factor in Eq.~\ref{eq:sigma} depends on the assumption that the clusters are spherical. A different ratio of the surface to the volume would modify this term. This is likely for small graphite clusters, since graphite tends to aggregate in a planar form. \item {\it Detachment rate for small clusters} --- In this paper, and in most nucleation theory, the detachment rate is computed by propagating to very small clusters the detachment rate of macroscopic bodies. It is very likely that, in the very small limit, the detachment is governed by completely different processes. Let us analyze the two limits. For a macroscopic body, the number of monomers is so large that monomers with a statistically higher energy can detach since their energy is larger than the binding energy. In the opposite limit of a dimer, the detachment has to be due to an external action: either a collision with a fast monomer or with a photon or with another cluster. In this limit destructive collisions have to be considered. \item {\it Cooling of macromolecules} --- This is a new problem that arises when the cooling of the grains is considered. We have assumed that the grains cool as modified black bodies down to the smallest sizes. When the number of monomers in the cluster becomes small, the cooling will not be through a continuum spectrum, but through lines and bands. A more refined treatment of cooling is necessary to compute accurate nucleation rates. \item It should finally be kept in mind that we allowed for the complete cooling of the grains, neglecting the effects of a background radiation field in setting a lower limit to the temperature. In addition, we neglected the fact that the temperature of small grains is largely stochastic and that at high temperature some collisions between grains and gas particles can result in sputtering rather than accretion. \end{itemize} Despite all these caveats, the main result of this paper holds: the region where a vapor spontaneously nucleate is not limited to the classic region, when nucleation takes place in thermal equilibrium and $T_g=T_s$. Allowing for the clusters to cool inhibits the detachment of monomers from the clusters and allow for nucleation at higher temperatures and lower densities than in the classical scenario. The nucleation rates tend to be orders of magnitude smaller than those derived in thermal equilibrium, but provide a non-negligible correction to the LTE rate for mildly super-saturated vapors with $S\gsim1$, especially at low temperature. In the astrophysical scenario small nucleation rates are not a big worry. The ISM of our Galaxy contains approximately 1 per cent of its mass in dust grains (Mathis et al. 1977). This corresponds to approximately one dust particle every cubic meter (assuming a power-law grain size distribution as in Mathis et al. 1977). However, in the present day Universe the ISM is already polluted with dust and in the presence of seed grains it is likely that the process of mantle growth dominates over non-LTE nucleation in the diffuse medium. The process of non-LTE nucleation may therefore be important at high redshift, when the ISM is first polluted with metals by supernova explosions. It is unclear if supernov\ae\ do generate dust by themselves, and even more whether the generated dust can survive the sputtering in the forward-reverse shock systems (Bianchi \& Schneider 2007; Nath, Laskar \& Shull 2007). In the case that supernov\ae\ do mainly pollute the ISM with metals but with no or a negligible quantity of dust, non-LTE nucleation could become the dominant process of dust nucleation in the early universe. Detailed estimates of nucleation in the various scenarios require a more detailed understanding of the properties of the very small nuclei and are beyond the scope of this paper. | 7 | 10 | 0710.2903 |
0710 | 0710.2418_arXiv.txt | { We report here on a new VHE source, HESS~J1908+063, disovered during the extended H.E.S.S. survey of the Galactic plane and which coincides with the recently reported MILAGRO unidentified source MGRO~J1908+06. The position, extension and spectrum measurements of the HESS source are presented and compared to those of MGRO~J1908+06. Possible counterparts at other wavelenghts are discussed. For the first time one of the low-lattitude MILAGRO sources is confirmed. } \begin{document} | H.E.S.S. observations of the inner Galactic plane in the [$270^{\circ}$, $30^{\circ}$] longitude range have revealed more than two dozens of new VHE sources, consisting of shell-type SNRs, pulsar wind nebulae, X-ray binary systems, a putative young star cluster, etc, and yet unidentified objects (see e.g. \cite{HESSScanII} and \cite{HESSSurveyICRC07} in these proceedings for a summary). The extended H.E.S.S. survey in the [$30^{\circ}$-$60^{\circ}$] longitude range performed between 2005 and 2007 overlaps with regions covered by the MILAGRO sky survey at longitudes greater than $30^{\circ}$. The latter experiment has recently reported \cite{MILAGRO} three low-latitude sources including, MGRO~J1908+06, detected after seven years of operation (2358 days of data) at 8.3$\sigma$ (pre-trials) confidence level. MGRO~J1908+06, of which the extension remains unknown but bounded to a maximum diameter of 2.6$^{\circ}$, is located near the galactic longitude $\sim 40^{\circ}$ and hence is covered by the H.E.S.S. galactic plane survey. A new H.E.S.S. source, HESS~J1908+063, which coincides with MGRO~J1908+06, is presented here. Its position, size and spectrum are measured and compared to the MILAGRO source. Possible counterparts at other wavelengths are discussed in the light of the H.E.S.S. measurements. | 7 | 10 | 0710.2418 |
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0710 | 0710.2126_arXiv.txt | $\mu$ Orionis was identified by spectroscopic studies as a quadruple star system. Seventeen high precision differential astrometry measurements of $\mu$ Ori have been collected by the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES). These show both the motion of the long period binary orbit and short period perturbations superimposed on that caused by each of the components in the long period system being themselves binaries. The new measurements enable the orientations of the long period binary and short period subsystems to be determined. Recent theoretical work predicts the distribution of relative inclinations between inner and outer orbits of hierarchical systems to peak near 40 and 140 degrees. The degree of coplanarity of this complex system is determined, and the angle between the planes of the A-B and Aa-Ab orbits is found to be $136.7 \pm 8.3$ degrees, near the predicted distribution peak at 140 degrees; this result is discussed in the context of the handful of systems with established mutual inclinations. The system distance and masses for each component are obtained from a combined fit of the PHASES astrometry and archival radial velocity observations. The component masses have relative precisions of $5\%$ (component Aa), $15\%$ (Ab), and $1.4\%$ (each of Ba and Bb). The median size of the minor axes of the uncertainty ellipses for the new measurements is 20 micro-arcseconds ($\microas$). Updated orbits for $\delta$ Equulei, $\kappa$ Pegasi, and V819 Herculis are also presented. | $\mu$ Orionis (61 Ori, HR 2124, HIP 28614, HD 40932) is a quadruple star system that has been extensively studied by radial velocity (RV) and differential astrometry. It is located just North of Betelgeuse, Orion's right shoulder (left on the sky); $\mu$ Ori is a bright star that is visible to the unaided eye even in moderately light-polluted skies. \cite{Frost1906} discovered it to be a short period (four and a half day) single-lined spectroscopic binary; this was component Aa, whose short-period, low mass companion Ab has never been detected directly. \cite{Aitken1914} discovered it also had a more distant component (B) forming a sub-arcsecond visual binary. Much later, \cite{Fekel1980} found B was itself a short-period (4.78 days) double-lined spectroscopic binary, making the system quadruple; these stars are designated Ba and Bb. Most recently, \cite{Fekel2002} (hereafter F2002) reported the astrometric orbit of the A-B motion, double-lined RV orbits for A-B and the Ba-Bb subsystem, and a single-lined RV orbit for the Aa-Ab subsystem. F2002 estimate the spectral types as A5V (Aa, an Am star), F5V (Ba), and F5V (Bb), though they note these are classifications are less certain than usual due to the complexity of the system. For a more complete discussion of the history of $\mu$ Ori, see F2002. Until now, astrometric observations have only been able to characterize the long period A-B motion, lacking the precision necessary to measure the astrometric perturbations to this orbit caused by the Aa-Ab and Ba-Bb subsystems. The method described by \cite{LaneMute2004a} for ground-based differential astrometry at the $\sim 20$ $\microas$ level for sub-arcsecond (``speckle'') binaries has been used to study $\mu$ Ori during the 2004-2007 observing seasons. These measurements represent an improvement in precision of over two orders of magnitude over previous work on this system. The goal of the current investigation is to report the center-of-light (photocenter) astrometric orbits of the Aa-Ab and Ba-Bb subsystems. This enables measurement of the coplanarities of the A-B, Aa-Ab, and Ba-Bb orbits. The masses and luminosity ratio of Aa and Ab are measured for the first time.Also presented are updated orbits for the PHASES targets $\delta$ Equ, $\kappa$ Peg, and V819 Her. Astrometric measurements were made at the Palomar Testbed Interferometer \citep[PTI;][]{col99} as part of the Palomar High-precision Astrometric Search for Exoplanet Systems (PHASES) program \citep{Mute06Limits}. PTI is located on Palomar Mountain near San Diego, CA. It was developed by the Jet Propulsion Laboratory, California Institute of Technology for NASA, as a testbed for interferometric techniques applicable to the Keck Interferometer and other missions such as the Space Interferometry Mission (SIM). It operates in the J ($1.2 \mu{\rm m}$), H ($1.6 \mu{\rm m}$), and K ($2.2 \mu{\rm m}$) bands, and combines starlight from two out of three available 40-cm apertures. The apertures form a triangle with one 110 and two 87 meter baselines. | The center-of-light astrometric motions of the Aa-Ab and Ba-Bb subsystems in $\mu$ Ori have been constrained by PHASES observations. While four degenerate orbital solutions exist, two of these can be excluded with high reliability based on mass-luminosity arguments, and the fact that Ab is not observed in the spectra. Ba and Bb are stars of a class (mid-F dwarfs) whose properties have been well established by studying other binaries. Their association with Aa and Ab, which are members of more poorly studied classes (Am and late K dwarfs) allows a better understanding of those objects in a system which can be assumed to be coevolved. The orbital solution finds masses and luminosities for all four components, the basic properties necessary in studying their natures. Complex dynamics must occur in $\mu$ Ori. The Ba-Bb orbital plane is nearly perpendicular to that of the A-B motion, and certainly undergoes Kozai-type inclination-eccentricity oscillations. It is possible that the mutual inclination of the A-B pair and Aa-Ab subsystem is a result of KCTF effects over the system's evolution. Finally, it is noted that the orbits in the $\mu$ Ori system are quite non-coplanar. This is in striking contrast with the planets of the solar system, but follows the trend seen in triple star systems. With the solar system being the only one whose coplanarity has been evaluated, it is difficult to draw conclusions about the configurations of planetary systems in general. It is important that future investigations evaluate the coplanarities of extrasolar planetary systems to establish a distribution. Whether that distribution will be the same or different than that of their stellar counterparts may point to similarities or differences in star and planet formation, and provide a key constraint on modeling multiple star and planet formation. | 7 | 10 | 0710.2126 |
0710 | 0710.0854_arXiv.txt | {Many thermally emitting, isolated neutron stars have magnetic fields that are larger than $10^{13}$~G. A realistic cooling model that includes the presence of high magnetic fields should be reconsidered.} {We investigate the effects of an anisotropic temperature distribution and Joule heating on the cooling of magnetized neutron stars.} {The 2D heat transfer equation with anisotropic thermal conductivity tensor and including all relevant neutrino emission processes is solved for realistic models of the neutron star interior and crust.} {The presence of the magnetic field affects significantly the thermal surface distribution and the cooling history during both, the early neutrino cooling era and the late photon cooling era.} { There is a large effect of Joule heating on the thermal evolution of strongly magnetized neutron stars. Both magnetic fields and Joule heating play an important role in keeping magnetars warm for a long time. Moreover, this effect is important for intermediate field neutron stars and should be considered in radio--quiet isolated neutron stars or high magnetic field radio--pulsars.} | Observation of thermal emission from neutron stars (NSs) can provide not only information on the physical properties such as the magnetic field, temperature, and chemical composition of the regions where this radiation is produced but also information on the properties of matter at higher densities deeper inside the star \citep{Yakovlev2004,Page2006}. To derive this information, we need to calculate the structure and evolution of the star, and compare the theoretical model with the observational data. Most previous studies assumed a spherically symmetric temperature distribution. However, there is increasing evidence that this is not the case for most nearby neutron stars whose thermal emission is visible in the X-ray band of the electromagnetic spectrum \citep{Zavlin2007,Haberl2007}. The anisotropic temperature distribution may be produced not only in the low density regions where the spectrum is formed and preliminary investigations had focused their attention, but also in intermediate density regions, such as the solid crust, where a complicated magnetic field geometry could cause a coupled magneto-thermal evolution. In some extreme cases, this anisotropy may even be present in the poorly known interior, where neutrino processes are responsible for the energy removal. The observational fact that most thermally emitting isolated NSs have magnetic fields larger than $10^{13}$ G \citep{Haberl2007}, which is sometimes confirmed by spin down measurements, leads to the conclusion that a realistic cooling model must not avoid the inclusion of the effects produced by the presence of high magnetic fields. The transport processes that occur in the interior are affected by these strong magnetic fields and their effects are expected to have observable consequences, in particular for highly magnetized NSs or magnetars. Moreover, the large surface magnetic field strengths inferred from the observations probably indicate that the interior field could reach even larger values, as theoretically predicted by some models \citep{TD1993}. The presence of a magnetic field affects the transport properties of all plasma components, especially the electrons. In general, the motion of free electrons perpendicular to the magnetic field is quantized in Landau levels, and the thermal and electrical conductivities exhibit quantum oscillations. In the limit of a strongly quantizing field, in which almost all electrons populate the lowest level, such as in the envelope of a NS, a quantum description is necessary to calculate the thermal and electrical conductivities. Earlier calculations by \cite{Canuto1970} and \cite{Itoh1975} concluded that the electron thermal conductivity is strongly suppressed in the direction perpendicular to the magnetic field and increased along the magnetic field lines, which reduces the thermal insulation of the envelope ({\it heat blanketing}). Thus, there is an anisotropic heat transport in the NS's envelope governed by the magnetic field geometry, that produces a non-uniform surface temperature. The anisotropy in the surface temperature of a NS appears to be confirmed by the analysis of observational data from isolated NSs (see \cite{Zavlin2007} and \cite{Haberl2007} for reviews on the current status of theory and observations). The mismatch between the extrapolation to low energy of fits to the X-ray spectra, and the observed Rayleigh-Jeans tail in the optical band ({\it optical excess flux}), cannot be addressed using a unique temperature. Several simultaneous fits to multiwavelength spectra of \rxdieciocho \citep{Pons2002,Truemper2004}, \rbdoce \citep{Schwope2005,Schwope2007}, and \rxcerosiete \citep{Perez2006} are explained by a small hot emitting area $\simeq$ 10--20 km$^2$, and an extended cooler component. Another piece of evidence that strongly supports the nonuniform temperature distribution are pulsations in the X-ray signal of some objects of amplitudes $\simeq$ 5--30 $\%$, some of which have irregular light curves that point towards a non-dipolar temperature distribution. All of these facts reveal that the idealized picture of a NS with a dipolar magnetic field and uniform surface temperature is oversimplified. In a pioneering work, \cite{Greenstein1983} obtained the temperature at the surface of a NS as a function of the magnetic field inclination angle in a simplified plane-parallel approximation. This model was applied to different magnetic field configurations and the observational consequences of a non-uniform temperature distribution were analyzed in the pulsars Vela and Geminga among others \citep{Page1995}. \cite{Potekhin2001} improved the former calculations including realistic thermal conductivities. Nevertheless, the temperature anisotropy as generated in the envelope may be insufficiently to be consistent with the observed thermal distribution and, in this case, should originate deeper inside the NS \citep{Geppert2004,Azorin2006}. Crustal confined magnetic fields could be responsible for the surface thermal anisotropy. In the crust, even if a strong magnetic field is present, the electrons occupy a large number of Landau levels and the classical approximation remains valid during a long time in the thermal evolution. The magnetic field limits the movement of electrons in the direction perpendicular to the field and, since they are the main carriers of the heat transport, the thermal conductivity in this direction is highly suppressed, while remaining almost unaffected along the field lines. Temperature distributions in the crust were obtained as stationary solutions of the diffusion equation with axial symmetry \citep{Geppert2004}. The approach assumes an isothermal core and a magnetized envelope as an inner and outer boundary condition, respectively. The results show important deviations from the crust isothermal case for crustal confined magnetic fields with strengths larger than $10^{13}$ G and temperatures below $10^{8}$ K. Similar conclusions were obtained considering not only poloidal but also toroidal components for the magnetic field \citep{Azorin2006, Geppert2006}. This models succeeded in explaining simultaneously the observed X-ray spectrum, the optical excess, the pulsed fraction, and other spectral features for some isolated NS such as \rxcerosiete \citep{Perez2006} and \rxdieciocho \citep{Geppert2006}. Non-uniform surface temperature in NSs was studied by different authors using simplified models \citep{ShibYak1996,Potekhin2001}. Although these models can provide useful information, a detailed investigation of heat transport in 2D must be completed to obtain more solid conclusions. However, this is not the only effect that must be revisited to study the cooling of NSs. For isolated NSs, different relevant magnetic field dissipation processes were identified \citep{Goldreich1992}. The {\it Ohmic} dissipation rate is determined by the finite conductivity of the constituent matter. In the crust, the electrical resistivity is mainly due to electron-phonon and electron-impurity scattering processes \citep{Flowers1976}, resulting in more efficient Ohmic dissipation than in the fluid interior. The strong temperature dependence of the resistivity leads to rapid dissipation of the magnetic energy in the outermost low-density regions during the early evolution of a hot NS, which becomes less relevant as the star cools down. Joule heating in the crustal layers due to Ohmic decay was thought to affect only the late photon cooling era in old NS ($\geq 10^7$ yr), and to be an efficient mechanism to maintain the surface temperature as high as $\simeq 10^{4-5}$ K for a long time \citep{Miralles98}. \cite{Page2000} studied the 1-D thermal evolution of NSs combined with an evolving Stokes function that defines a purely poloidal, dipolar magnetic field. The Joule heating rate was evaluated averaging the currents over the azimuthal angle. However, for strongly magnetized NS, Joule heating can be important much earlier in the evolution. In a recent work, \cite{Kaminker2006} placed a heat source inside the outer crust of a young, warm, magnetar of field strength $5\times10^{14}$ G. To explain observations, they concluded that the heat source should be located at a density $\lesssim 5\times 10^{11}$ g~cm$^{-3}$, and the heating rate should be $\gtrsim 10^{20}$ erg~cm$^{-3}$~s$^{-1}$ for at least $5\times10^4$ yr. Anisotropic heat transport is neglected in these simulations, which were performed in spherical symmetry, assuming that it will not affect the results in the early evolution. Nevertheless we will show that, in 2D simulations, the effect of anisotropic heat transport is important. In addition to purely Ohmic dissipation, strongly magnetized NSs can also experience a {\it Hall drift} with a drift velocity proportional to the magnetic field strength. Although the Hall drift conserves the magnetic energy and it is not a dissipative mechanism by itself, it can enhance the Ohmic decay by compressing the field into small scales, or by displacing currents to regions with higher resistivity, where Ohmic dissipation is more efficient. Recently, the first 2D-long term simulations of the magnetic field evolution in the crust studied the interplay of Ohmic dissipation and the Hall drift effect \citep{PonsGeppert2007}. It was shown that, for magnetar field strength, the characteristic timescale during which Hall drift influences Ohmic dissipation is of about $10^{4}$ yr. All of these studies imply that both field decay and Joule heating play a role in the cooling of neutron stars born with field strengths $\geq 10^{13}$ G. We will show that, during the neutrino cooling era and the early stages of the photon cooling era, the thermal evolution is coupled to the magnetic field decay, since both cooling and magnetic field diffusion proceed on a similar timescale ($\approx 10^{6} $ yr). The energy released by magnetic field decay in the crust could be an important heat source that modifies or even controls the thermal evolution of a NS. Observational evidence of this fact is shown in \cite{PonsLink2007}. They found a strong correlation between the inferred magnetic field and the surface temperature for a wide range of magnetic fields: from magnetars ($\geq 10^{14}$ G), through radio-quiet isolated neutron stars ($\simeq 10^{13}$ G) down to some ordinary pulsars ($\leq 10^{13}$ G). The main conclusion is that, rather independently from the stellar structure and the matter composition, the correlation can be explained by the decay of currents on a timescale of $\simeq 10^{6}$ yr. The aim of the present work is to study in a more consistent way the cooling of a realistic NS under the effects of large magnetic fields, including the effects of an anisotropic temperature distribution and Joule heating in 2D simulations. As a first step towards a fully coupled magneto-thermal evolution, a phenomenological law for the magnetic field decay is considered. This article is structured as follows. In Sect.~2 we discuss the equations governing the magnetic field structure and evolution, while Sect.~3 is devoted to the thermal evolution equations. Sect.~4 presents the microphysics inputs. Sect.~5 and 6 contain our results for weakly and strongly magnetized NSs, respectively. In Sect. 7, we focus on the effects of field decay and Joule heating on the cooling history of a NS. Finally, in Sect. 8 we present the main conclusions and perspectives of the present work. | We have presented a thorough study of the thermal evolution of neutron stars including some of the most intriguing effects of magnetic fields. Our results were based on two-dimensional cooling simulations of realistic models that account for the anisotropy in the thermal conductivity tensor. In the first part of the paper, we revisited the classical scenario with low magnetic fields and presented the input microphysics, working assumptions, and the baseline models. As an interesting byproduct, we reconsidered the growth of the crust and of the superconducting region in the NS core, and found that there are situations in which both growth rates are comparable. The main body of the work was aimed at the discussion of the two principal effects of magnetic fields: the anisotropic surface temperature distribution and the additional heating by magnetic field decay. We found that, even for purely dipolar fields, an inverted temperature distribution is plausible at intermediate ages. Thus the surface temperature distribution of neutron stars with high magnetic fields, even in the axisymmetric case, may be quite different from the model with two hot polar caps and a cooler equatorial region. The irregular light curves of some isolated neutron stars, for instance RBS1223, \citep{Schwope2005,Kaplan2005} are an indication of such complex structures. The main result of this work is that, in NSs born as magnetars, Joule heating has an enormous effect on the thermal evolution. Moreover, this effect is important for intermediate field stars. If the magnetic field is supported by crustal currents, this effect can reach a maximum because two combined factors enhance the efficiency of the heating process: i) more heat is released into the crust, in the regions of higher resistivity close to the surface, and ii) large non radial components of the field channel the heat towards the surface, instead of being lost by neutrinos in the core. As expected, it becomes clear that magnetic fields and Joule heating are playing a key role keeping magnetars warm for a long time, but it is likely that the same effect, although quantitatively smaller, must be considered in radio--quiet isolated NSs or high magnetic field radio--pulsars. Another aspect that should be considered when we try to explain observations using theoretical cooling curves is that for many objects the age is estimated assuming that the loss of angular momentum is entirely due to dipolar radiation from a magnetic dipole (spin-down age). In the case of a decaying magnetic field, the {\em spin down age}, seriously overestimates the {\em true age} \citep{GO1970}. Therefore, the cooling evolution time should be corrected, according to the prescription for magnetic field decay, to compare our model accurately with observations. A detailed comparison of the cooling curves obtained in this work with observational sources can be found in \cite{Aguilera2008}. Our last striking remark is that the occurrence of direct URCA or, in general, fast neutrino cooling in NS may be hidden by a combination of effects due to strong magnetic fields. Our conclusion is that the most appropiate candidates to monitor as rapid coolers are NSs with fields lower than $10^{13}$ G. Otherwise, we may be misled in our interpretation of the temperature-age diagrams. The main drawback of our work is that we are not yet able to return a fully consistent simulation of the magneto-thermal coupled evolution of temperature and magnetic field. In the near future, we plan to extend this study by coupling our thermal diffusion code to the consistent evolution of the magnetic field in the crust given by the Hall induction equation. That approach will permit the accurate evaluation of the heating rates, including the non-linear effects associated with the Hall--drift in the NS crust. We believe, however, that the phenomenological parameterization employed in this paper, reproduces qualitatively the results expected in a real case. We have provided another step towards understanding the cooling of neutron stars, by pointing out a number of important features that must be more carefully considered in future work. | 7 | 10 | 0710.0854 |
0710 | 0710.0406.txt | We report results from a search for massive and evolved galaxies at $z \ga 5$ in the Great Observatories Origins Deep Survey (GOODS) southern field. Combining HST ACS, VLT ISAAC and Spitzer IRAC broad--band photometric data, we develop a color selection technique to identify candidates for being evolved galaxies at high redshifts. The color selection is primarily based on locating the Balmer--break using the K- and 3.6$\mu$m bands. Stellar population synthesis models are fitted to the SEDs of these galaxies to identify the final sample. We find 11 candidates with photometric redshifts in the range $4.9 \leq z < 6.5$, dominated by an old stellar population, with ages 0.2$-$1.0 Gyr, and stellar masses in the range $(0.5 - 5) \times 10^{11}$ M$_{\odot}$. The majority of the stars in these galaxies were formed at $z > 9$ and the current star formation activity is in all cases, except two, a few percent of the inferred early star formation rate. One candidate has a spectroscopically confirmed redshift, in good agreement with our photometric redshift. The galaxies are very compact, with half--light radii in the observed $K-$band smaller than $\sim 2$ kpc. Seven of the 11 candidates are also detected at 24$\mu$m with the MIPS instrument on Spitzer. %%% By itself, the 24$\mu$m emission could potentially be interpreted as PAH emission from a dusty starburst at $z \sim 2-3$, however, it is also consistent with the presence of an obscured AGN at $z \ga 5$. Indeed, for the $z \ga 5$ solutions, all the MIPS detected galaxies, except two, %BBG\#3348 and JD2, have relatively high internal extinction. While we favor the obscured AGN interpretation, based on the model SED fits to the optical/UV, we define a 'no--MIPS' sample of candidates in addition to the full sample. Results will be quoted for both samples. % We estimate the completeness of the Balmer break galaxy sample to be $\sim$40\% (an upper limit). The comoving number density of galaxies with a stellar mass $\ga 10^{11}$ M$_{\odot}$, at an average redshift $\bar{z} = 5.2$, is $3.9 \times 10^{-5}$ Mpc$^{-3}$ (no--MIPS sample: $1.4 \times 10^{-5}$ Mpc$^{-3}$). The corresponding stellar mass density is $8 \times 10^{6}$ M$_{\odot}$ Mpc$^{-3}$ (no--MIPS sample: $6.2 \times 10^{6}$ M$_{\odot}$ Mpc$^{-3}$). The estimated stellar mass density at $\bar{z} = 5.2$ is $2-3$\% of the present day total stellar mass density and $20-25$\% of the stellar mass density at $z \sim 2$. If the stellar mass estimates are correct, the presence of these massive and evolved galaxies when the universe was $\sim$1 Gyr old could suggest that conversion of baryons into stars proceeded more efficiently in the early universe than it does today. | An important goal of observational cosmology is to understand how stars are assembled into galaxies and how this is related to the evolution of dark matter halos. In prevailing hierarchical models, star formation starts out in low mass systems, which build more massive galaxies through sequential merging (e.g. White \& Rees 1978; Somerville 2004). In this picture, the most massive galaxies are found at relatively low redshifts. Recently, a significant population of galaxies with stellar mass $\sim 10^{11}$ M$_{\odot}$ has been found at $z \sim 2 - 3$ (cf. Franx et al. 2003; Glazebrook et al. 2004; Fontana et al. 2004; Yan et al. 2004; Daddi et al. 2005a; Rudnick et al. 2006; van Dokkum et al. 2006). Stellar population synthesis models combined with broad-band photometric data show that many of these galaxies contain an old stellar population, with ages indicating a star formation phase within 1$-$3 Gyr after the Big Bang. Moreover, a number of submillimeter detected galaxies at $z \sim 2 - 3$, which are known to be massive systems, based on their inferred molecular gas and dynamical mass estimates (cf. Greve et al. 2005), also appear to contain an old stellar population with mass $\sim 10^{11}$ M$_{\odot}$ (Borys et al. 2005). Therefore, a consensus seems to be emerging, that the most massive galaxies seen today, formed the bulk of their stars within the first $\sim$3 Gyr of cosmic history (cf. Cimatti et al. 2004; Daddi et al. 2005b; Juneau et al. 2005). However, it is not known how these stars were assembled into their present host galaxies, whether this was done during multiple merger events, as proposed in hierarchical models, or if the stars and their host galaxy are co-eval. In view of the early formation epoch implied for many of these massive galaxies, the question whether the formation is hierarchical or monolithic becomes a matter of semantics as the merger time scale becomes comparable to the dynamical time scale. Recent ultra--deep surveys, done at wavelengths stretching from the UV to mid--infrared, have resulted in detection of galaxies and AGNs at even higher redshifts, reaching into the era of re-ionization. One example is HUDF--JD2, in the Hubble Ultra Deep Field, which Mobasher et al. (2005) identify as a candidate for a massive, evolved galaxy at $z=6.5$. The age of this galaxy is estimated to be $\ga$600 Myrs, with a stellar mass of $\sim 6 \times 10^{11}$ M$_{\odot}$, much larger than the stellar mass of the Milky Way. The implied age of this galaxy means that the bulk of the stars were formed on a short time scale just a few hundred million years after the recombination era. Other recent studies have used data from the Spitzer Space Telescope to analyze the stellar masses and ages for galaxies at $z > 5$ (cf. Yan et al. 2005, 2006; Eyles et al. 2005, 2006; Stark et al. 2006; Verma et al. 2007). The inferred stellar masses range from $1 - 10 \times 10^{10}$ M$_{\odot}$ and ages of several $10^8$ years. In several cases, galaxies have spectroscopically determined redshifts. Another spectroscopically confirmed galaxy is the gravitationally lensed object HCM06 at $z = 6.56$ (Hu et al. 2002), with a stellar mass of a few $10^{10}$ M$_{\odot}$ and an age of $\sim$300 Myr (cf. Chary, Stern \& Eisenhardt 2005; Schaerer \& Pell\'{o} 2005). The presence of these massive and old galaxies at $z \ga 5$ holds important clues for understanding how the first galaxies formed and how the galaxy population in general has evolved with cosmic time. In order to determine whether a significant population of massive and old galaxies exists at $z > 5$, and to derive the parameters characterizing this population, we need a selection method that specifically targets and selects evolved stellar systems at very high redshifts, using broad-band photometric data available from deep multiwavelength surveys. The presence of old galaxies at high redshift can not efficiently be inferred using the normal Lyman drop-out technique. The drop-out technique has proven to be efficient in detecting galaxies that are actively forming stars and contain relatively small amounts of dust, but it is ineffective for detecting galaxies without strong UV continuum. In this paper we develop a method for selecting galaxies dominated by a stellar population older than $\sim$100 Myr and situated at $z \ga 5$, and discuss the results and its implications. The technique is primarily based on detecting the presence of a well--developed Balmer break, redshifted to $\sim$3$\mu$m, that can be probed by the K$_\mathrm{s} - 3.6\mu$m color index. A second color index is used to further isolate the old high-z galaxies from foreground `contaminants'. The color signature of the Balmer break has previously been used to select galaxies at redshifts $z \sim 1 - 3$ (Franx et al. 2003; Daddi et al. 2005a; Adelberger et al. 2004). By choosing a suitable filter combination, the Balmer break can be used to select galaxies at any redshift, in a manner similar to the Lyman--break technique. In this paper we will refer to the galaxies selected through this technique at $z > 5$ as Balmer--break galaxies (BBGs). The paper is structured as follows: % In Sect~\ref{data} we present our sample and photometric data. % Sect.~\ref{selecting} discusses the Balmer break feature, the stellar population synthesis models used and examines the confidence of the model fitting procedure using Monte Carlo simulations. In this section we also discuss degeneracies and define the final color selection criteria used in this paper. % In Sect.~\ref{results} we present the Balmer--break candidates selected using our color criteria and model fits of synthetic stellar spectra. We derive associated physical parameters from the models and discuss the Spitzer/MIPS 24$\mu$m detections. In this section we also discuss the individual sources and assign a confidence classification to each source based on its likelihood to have the correct redshift. % In Sect.~\ref{testing} we apply our model fitting to galaxies with known spectroscopic redshift and assess the reliability of the estimated parameters. % In Sect.~\ref{errors} we discuss different sources of errors and derive the completeness of our sample. % We discuss our results in Sect.~\ref{discussion} and compare the number density of Balmer--break galaxies with the expected number density of dark matter halos. % Sect.~\ref{summary} gives a summary of our results. % We adopt H$_0 = 72$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_m = 0.3$, and $\Omega_{\Lambda} = 0.7$ throughout this paper. All magnitudes are in the AB system (Oke 1974). | \label{discussion} The existence of massive and evolved galaxies at redshifts $z \ga 5$, when the universe was $\la 1$ Gyr old, seems surprising at first sight. In the hierarchical scenario for galaxy formation, the majority of massive galaxies are assembled at relatively low redshifts. However, the presence of massive galaxies at high redshifts poses a fundamental problem for hierarchical models only if their number density exceeds that of correspondingly massive dark matter halos (e.g. Somerville 2004). In Sect.~\ref{completeness} we derived the number and mass density of the Balmer--break galaxies, using our sample of 11 galaxies, as well as for a more restricted sample only containing those candidates which are not detected with MIPS at 24$\mu$m, the 'no--MIPS sample'. %Dark matter halos By equating the co--moving number density of the Balmer--break galaxies with the expected density of dark matter halos at the same redshift, we can estimate the the maximum halo mass associated with the BBG's. Using the Sheth--Tormen modified Press--Schechter formalism (Sheth \& Tormen 1999), with the dark matter halo concentration predicted for the revised value of the power spectrum normalization $\sigma_{8}=0.74$ (Spergel et al. 2007), and the estimated lower limit to the number density of BBGs ($3.9 \times 10^{-5}\,\eta^{-1}$ Mpc$^{-3}$), we predict a halo mass of $M_{\mathrm{h}} = 1.0 \times 10^{12}$ M$_{\odot}$ (assuming the standard $\Lambda$CDM model)\footnote{The estimated halo mass is a non-linear function of the incompleteness coefficienct $\eta$. For instance, with $\eta = 0.5$, the corresponding halo mass becomes $8 \times 10^{11}$ M$_{\odot}$.}. Using the average stellar mass for the BBGs, we get $M_{*}/M_{\mathrm{h}} \approx 0.20$. Considering the no--MIPS sample, with a number density $1.4 \times 10^{-5}\,\eta^{-1}$ Mpc$^{-3}$, the corresponding halo mass is $1.3 \times 10^{12}$ M$_{\odot}$, giving a ratio $M_{*}/M_{\mathrm{h}} \approx 0.08$. This estimate of the halo mass assumes that all available $\sim 10^{12}$ M$_{\odot}$ halos at $z \sim 5.2$ are associated with Balmer--break galaxies. If a fraction of these halos would host lower mass stellar systems, such as Lyman--break galaxies, the $M_{*}/M_{\mathrm{h}}$ ratio for the Balmer--break galaxies would have to increase accordingly. The ratio of the baryonic--to--total mass can be expressed in terms of a star formation efficiency parameter, $\beta$ ($M_{*} = \beta\,M_{\mathrm{baryon}}$: the fraction of baryons turned into stars over the life time of the galaxy), and the stellar mass, M$_{*}$, as, $$ M_{\mathrm{baryon}}/M_{\mathrm{total}} = \beta^{-1}M_{*}/M_{\mathrm{total}} = \kappa $$ where $M_{\mathrm{total}} = \beta^{-1}\,M_{*} + M_{\mathrm{h}}$. We can then write $$ M_{*}/M_{\mathrm{h}} = \beta\,\kappa/(1 - \kappa). $$ Adopting $\kappa = 0.17$ from the WMAP3 results (Spergel et al. 2007), we get $M_{*}/M_{\mathrm{h}} = 0.20\,\beta$. Klypin, Zhao \& Somerville (2002) estimate the total (virial) and baryonic mass of the Milky Way and M31 galaxies and find $M_{*}/M_{\mathrm{h}} = 0.06 - 0.08$, implying that in this case $\beta = 0.3 - 0.4$. For the Balmer--break galaxies, we find $\beta \approx 0.4 - 1.0$, where the lower value corresponds to the no--MIPS sample. If we only consider the no--MIPS sample, the baryonic fraction is comparable to local galaxies. However, for the full sample, the results suggests that the BBGs at $z \approx 5.2$ contain a higher fractions of baryons than galaxies at $z \approx 0$. Another possible explanation for the high baryonic fraction is that the number density of dark matter halos at high redshift is underestimated by the Sheth--Tormen analysis, or that we have systematically overestimated the stellar masses of the BBGs by a factor $\ga 2$. %%% Using a Chabrier or Kroupa initial mass function with a less steep low--mass end, could lower the estimated stellar masses by a factor 1.5--1.8 (see Sect.~\ref{systematics}). %%% It would be more instructive to compare the $M_{*}/M_{\mathrm{h}}$ ratio to that of massive elliptical galaxies at $z \approx 0$. However, the evidence for dark matter in elliptical galaxies is still circumstancial and limited to the central regions. Using planetary nebulae and globular clusters as kinematic probes, it has been possible to push the analysis to $\sim$5 $R_{\mathrm{eff}}$ (e.g. Romanowsky 2003; Richtler 2004). While the number of ellipticals studied in detail is still small, the general conclusion is that most have surprisingly weak dark matter halos, i.e. large $M_{*}/M_{\mathrm{h}}$ ratios. It remains to be determined whether the inferred $M_{*}/M_{\mathrm{h}}$ ratio for Balmer--break galaxies is consistent with local giant elliptical galaxies. %Stellar mass density The stellar mass density of the universe from redshifts 0 to 6 has been estimated by several groups, using different samples and methods (e.g. Bell et al. 2003; Dickinson et al. 2003; Rudnick et al. 2003, 2006; Fontana et al. 2006; Yan et al. 2006). Some of these results are listed in Table~\ref{stellarmass}, as a comparison with the results obtained for the BBGs. Most of the values listed in Table~\ref{stellarmass} are based only on the objects observed and are lower limits. In a few cases, the mass function has been integrated to obtain the total stellar mass (e.g. Dickinson et al. 2003). In the local universe, the global stellar mass density is $(3-4) \times 10^{8}$\,M$_{\odot}$ Mpc$^{-3}$, while it decreases to $\sim 0.3 \times 10^{8}$ M$_{\odot}$ Mpc$^{-3}$ at $z = 2.5-3$. In Sect.~\ref{completeness} we found that the stellar mass density of the 11 BBG candidates is $8 \times 10^{6}\,\eta^{-1}$ M$_{\odot}$ Mpc$^{-3}$. This is $\sim 2-3$\% of the present day total stellar mass density. Restricting the comparison to large early type galaxies in the local universe, that is, galaxies at least as massive as our BBG sample, the percentage increases to $\sim 4-6$\%. Comparing with the stellar mass density at $z \sim 2$, the BBG sample already comprise $20-25$\% of the total stellar mass found at this redshift. For the no--MIPS sample, the stellar mass density is $\sim$5 times smaller, and in this case the comparison with stellar mass densities at lower redshifts has to be corrected accordingly. The galaxies found in this study are remarkable in that they contain a large stellar mass, have small physical sizes and that their main epoch of star formation occured at $z \ga 10$. Galaxies with similar properties have, however, also been found by others. In a recent paper, McClure et al. (2006) searched for Lyman--break galaxies in the UKIDSS ultra deep survey, and found 9 candidates with $z > 5$. Their stellar masses are $>5 \times 10^{10}$ M$_{\odot}$ and their ages range from 50--500 Myr. Overall, these galaxies have properties similar to our Balmer--break galaxies. The number density for the $z > 5$ galaxies found by McClure et al. is $\sim$4x smaller than what we find in this paper. However, the different selection process, the fact that the UKIDSS sample does not include IRAC data and the large completeness corrections needed, makes a comparison difficult. A number of massive galaxies at $z > 4$ were also found by Fontana et al. (2006) using the GOODS-MUSIC sample. Their broad--band photometric data set consists of 14 bands, including the 4 IRAC bands. The objects were identified by fitting template SED based on Bruzual \& Charlot (2003) models to all galaxies in the sample. The best-fitting SEDs for the high redshift objects suggest that they are passively evolving galaxies, characterized by a very short time scale for star formation or by a constant star formation and a large amount of dust extinction. The stellar masses found are in excess of $10^{11}$ M$_{\odot}$. Hence, massive and passively evolving galaxies at $z \sim 5$ are found in several studies. A direct comparison of the results is presently not practical as different selection criteria are used, and the completeness corrections are presently poorly defined. Another surprising property of the Balmer--break galaxies is their compact sizes. As derived in Sect.~\ref{candidateparameters}, the typical half--light radius is $\la$2 kpc. Although this is larger than what is expected from the size vs. redshift relation derived for UV bright galaxies at similarly high redshift (Ferguson et al. 2004; Bouwens et al. 2004; Dahlen et al. in prep.), the stellar masses of the Balmer--break galaxies are at least 10 times higher. No massive compact galaxies of this type has been found in the local universe. However, compact galaxies with a large stellar mass have been found at $z \sim 1.4$ (Trujillo et al. 2006) and at $z\sim 2.5$ (Daddi et al. 2005b; Zirm et al. 2007; Toft et al. 2007). These galaxies are massive ($M_{*} > 10^{11}$ M$_{\odot}$), with no sign of AGN activity and contain a passively evolving stellar population, similar to the Balmer--break galaxies. The effective radius of these galaxies, measured at rest--frame optical wavelengths, are typically $\la 1$ kpc (Zirm et al. 2007; Toft et al. 2007), or 3--6 times smaller than local counterparts of similar stellar mass. It is hypothesized that the on--set of rapid star formation in these systems quench the star formation process, leading to very compact objects. These galaxies, as well as the Balmer--break galaxies, cannot represent fully assembled systems and must undergo subsequent evolution in their structural parameters in order to resemble local galaxies with the same stellar mass. %Implications for the reionization of the IGM The presence of a population of massive galaxies that underwent a period of intense star formation at $z \sim 10 - 25$ is likely to have ramifications to the reionization of the intergalactic medium (IGM). Panagia et al. (2005) calculated that the star formation associated with the formation of the massive $z = 6.5$ galaxy HUDF--JD2 (Mobasher et al. 2005), could significantly contribute to the reionization of the IGM. The main uncertainties were the escape fraction of the Lyman continuum photons and the volume density of objects similar to JD2. With the new sample of post--starburst galaxies with formation redshifts in the same range as JD2, it is possible to address this question. The integrated output of Lyman continuum photons from the Balmer--break candidates depends only on the total stellar mass and the assumed IMF (Panagia et al. 2005). Because the average stellar mass for the BBG candidates is about a factor 2 smaller than for JD2, assuming the same IMF, the average number of Lyman continuum photons is likewise a factor of 2 lower. Panagia et al. (2005) concluded that if each field of 2\ffam5 $\times$ 2\ffam5 contained a source like JD2, then these sources account for at least $\sim$20\% of the reionization of the IGM. A higher percentage is possible if the escape fraction is higher and/or the IGM is clumped. In the present case, we have a total area that is 25 times larger and a total ionizing photon output $\sim$10 times larger than in the case of JD2. Hence, the implication is that the BBG sources can account for $\sim$10\% or more of the photons needed for reionization, depending on the poorly constrained parameters describing the Lyman continuum escape fraction and the clumpiness of the IGM itself. The implications for reionization are discussed in more detail in Panagia et al. (in prep.). | 7 | 10 | 0710.0406 |
0710 | 0710.1255_arXiv.txt | I~have re-visited the spatial distribution of stars and high-mass brown dwarfs in the $\sigma$~Orionis cluster ($\sim$3\,Ma, $\sim$360\,pc). The input was a catalogue of 340 cluster members and candidates at separations less than 30\,arcmin to $\sigma$~Ori~AB. Of them, 70\,\% have features of extreme youth. I~fitted the normalised cumulative number of objects counting from the cluster centre to several power-law, exponential and King radial distributions. The cluster seems to have two components: a dense core that extends from the centre to $r \approx$ 20\,arcmin and a rarified halo at larger separations. The radial distribution in the core follows a power-law proportional to $r^1$, which corresponds to a volume density proportional to $r^{-2}$. This is consistent with the collapse of an isothermal spherical molecular cloud. The stars more massive than 3.7\,$M_\odot$ concentrate, however, towards the cluster centre, where there is also an apparent deficit of very low-mass objects ($M <$ 0.16\,$M_\odot$). Last, I~demonstrated through Monte Carlo simulations that the cluster is azimuthally asymmetric, with a filamentary overdensity of objects that runs from the cluster centre to the Horsehead~Nebula. | The $\sigma$~Orionis region in the {Ori~OB~1~b} association is finally becoming recognised as one of the most important young open clusters, with an age of only about 3\,Ma. In the discovery paper, Garrison (1967) used the term ``clustering'' to refer to an agglomeration of fifteen B-type stars surrounding and including the multiple star {$\sigma$~Ori}. Afterwards, Lyng\aa~(1981) tabulated $\sigma$~Orionis in his catalogue of open clusters. Since the rediscovery of the cluster by Wolk (1996) and its subsequent study in depth, which has revealed the most numerous and best known substellar population (B\'ejar et~al. 1999; Zapatero Osorio et~al. 2000, 2002; Caballero et~al. 2007), only a few authors have investigated the $\sigma$~Orionis spatial distribution. In particular, B\'ejar et~al. (2004) and Sherry et~al. (2004) analysed the radial distribution of $\sigma$~Orionis cluster members and candidates in annuli of width $\Delta r$ as a function of the separation $r$ to $\sigma$~Ori~AB. To maximise the number of objects per annulus and minimise the Poissonian errors, $\Delta r$ must be wide. This leads to have few annuli (no more than 12 in the $r$ = 0--30\,arcmin interval) to fit to a suitable radial profile (exponential decay -- B\'ejar et~al. 2004; King -- Sherry et~al.~2004). Both studies agree that the cluster may extend only up to $\sim$25--30\,arcmin. The low surface density of cluster members at larger separations, the sharp increase of extinction due to the nearby {Horsehead Nebula}-{Flame Nebula}-{IC~434} complex and the closeness to (or even overlapping with) other stellar populations in the Orion Belt surrounding {Alnitak} ($\zeta$~Ori) and {Alnilam} ($\epsilon$~Ori) prevent from suitably broaden the radial distribution analysis (Caballero 2007a). At the canonical heliocentric distance to $\sigma$~Orionis of 360\,pc (e.g. Brown, de~Geus \& de~Zeeuw 1994), the cluster would have an approximate radius of~3\,pc. In spite of the agreement on the size of $\sigma$~Orionis, the fits and the profiles in B\'ejar et~al. (2004) and Sherry et~al. (2004) seem to be rather incomplete and inappropiate, respectively. On the one hand, the King models were designed for tidally truncated globular clusters (King 1962, 1966; Meylan 1987), and have also been satisfactorily used for describing galaxies (e.g. Kormendy 1977; Binggeli, Sandage \& Tarenghi 1984). These systems have had enough time to be isothermal, on the contrary to very young open clusters like $\sigma$~Orionis, where only gravitational relaxation by initial mixing may have occurred (King 1962). On the other hand, B\'ejar et~al. (2004) exclusively focused on the cluster substellar population. Besides, the exponential fit in B\'ejar et~al. (2004) only accounted for the five innermost annuli, which leaded to a high uncertainty in the derived parameters. Last, in the works by Sherry et~al. (2004) and B\'ejar et~al. (2004), the input list of cluster members and candidates came from $VRI/IZ$ optical surveys. Many sources in both analysis had no near-infrared or spectroscopic follow-up. For a correct study of the spatial distribution in $\sigma$~Orionis, it is therefore necessary to use new fitting radial profiles and an input catalogue as comprehensive as possible. It must cover a wide mass interval. Maximum completeness and minimum contamination of the catalogue are also desired. These requirements are verified by the {\em Mayrit} catalogue, which tabulates 339 $\sigma$~Orionis members and candidates in a 30\,arcmin-radius circular area centred on $\sigma$~Ori~AB (Caballero 2007c). Of them, 241 display features of extreme youth (e.g. OB spectral types, Li~{\sc i} in absorption, H$\alpha$ in strong emission, spectral signatures of low gravity, near- and mid-infrared excesses due to discs). The catalogue covers three orders of magnitude in mass, from the $\sim$18+12\,$M_\odot$ of the O9.5V+B0.5V binary $\sigma$~Ori~AB to the $\sim$0.033\,$M_\odot$ of the brown dwarf {B05~2.03--617} (Caballero \& Chabrier, in~prep.). Accounting for $\sigma$~Ori~A and B as different objects separated by $\sim$0.25\,arcsec, then the equatorial coordinates of 340 young stars, brown dwarfs and cluster member candidates are available. I will use this input catalogue to investigate the radial and azimuthal distribution of objects in the $\sigma$~Orionis~region. | The $\sim$3\,Ma-old $\sigma$~Orionis cluster is a perfect laboratory of star formation. I~have investigated the radial distribution of 340 cluster members and candidates in a 30\,arcmin-radius area centred on $\sigma$~Ori~AB, taken from Caballero (2007c). The analysis has covered a mass interval from the 18+12\,$M_\odot$ of $\sigma$~Ori~AB to the $\sim$0.03\,$M_\odot$ of the faintest brown dwarf detectable by DENIS. The cluster shows a clear radial density gradient, quantified by the ${\mathcal Q}$-parameter, that accounts for the mean separation between members and the Euclidean minimum spanning tree of the cluster. I~have calculated the functional relations between normalised cumulative numbers of objects counting from the cluster centre, $f(r)$, and surface densities, $\sigma(r)$. Cumulative distribution functions as these avoid many problems associated with binning. Among the studied radial (power-law, exponential and King) profiles, the best fit is for a composite power-law distribution of cluster members with a core and a rarified halo. The core extends up to $\sim$20\,arcmin from the cluster centre and is nicely modelled by a surface density $\sigma(r) \propto r^{-1}$, that corresponds to a volume density $\rho(r) \propto r^{-2}$. This volume density matches, in its turn, the radial profile in a cluster formed from the collapse of a self-gravitating, isothermal sphere. The most massive $\sigma$~Orionis stars deviate, however, from the general trend and are much more concentrated towards the cluster centre. There is also an apparent deficit of very low-mass stars and high-mass brown dwarfs (0.16\,$M_\odot \ga M \ga$ 0.035\,$M_\odot$) in the innermost 4\,arcmin and an excess in the annulus at 6--7\,arcmin to the central Trapezium-like system. Last, there is a significant azimuthal asymmetry due to a filament-shape overdensity of objects that connects the cluster centre with a part of the Horsehead Nebula. This discovery supports the formation scenarios that predict burst of star formation in filamentary~gas. | 7 | 10 | 0710.1255 |
0710 | 0710.3911_arXiv.txt | % We follow the evolution of helium stars of initial mass $(2.2 - 2.5)\,M_\odot$, and show that they undergo off-center carbon burning, which leaves behind ${\mathbf \sim 0.01\,M_\odot}$ of unburnt carbon in the inner part of the core. When the carbon-oxygen core grows to Chandrasekhar mass, the amount of left-over carbon is sufficient to ignite thermonuclear runaway. At the moment of explosion, the star will possess an envelope of several $0.1\,M_{\odot}$, consisting of He, C, and possibly some H, perhaps producing a kind of peculiar SN. Based on the results of \citet{Waldman2007} for accreting white dwarfs, we expect to get thermonuclear runaway at a broad range of $\rho_c \approx (1 - 6) \times 10^9 \mathrm{ g\,cm^{-3}}$, depending on the amount of residual carbon. We verified the feasibility of this scenario by showing that in a close binary system with initial masses $(8.5 + 7.7)\,M_{\odot}$\ and initial period of 150 day the primary produces a helium remnant of $2.3\,M_{\odot}$\ that evolves further like the model we considered. | Type Ia supernovae (SN~Ia) have a relatively small dispersion of luminosity (the standard deviation in peak blue luminosity is $\sigma_B \approx 0.4 - 0.5$ mag., \citet{Branch1993ApJL}) and are being used as distance indicators (``standard candles''), having especial significance in the effort of determining the cosmological parameters of our universe. The long-standing explanation of the SN~Ia phenomenon has been the explosive burning of degenerate carbon in the core of a carbon-oxygen white dwarf, which becomes unstable as it grows to Chandrasekhar mass ($M_{Ch}$) either by accretion from a binary companion or by a merger of two white dwarfs, following the angular momentum loss from the system by gravitational wave radiation. However, theoretical models are still far from self-consistently producing an evolutionary path towards the progenitor and reproducing crucial features of the observational data, such as the composition of the ejecta. For a detailed review of the above see, e.g., \citet{Leibundgut2000A&ARv, Hillebrandt&Niemeyer2000ARA&A, Filippenko2005ASSL}. As well, SN~Ia can not be regarded as perfectly homogeneous class, since their Hubble diagram exhibits scatter larger than the photometric errors, while spectroscopic and photometric peculiarities have been noted with increasing frequency in well-observed SN~Ia \citep[e.g., ][]{Filippenko2005ASSL}. Therefore, there is an obvious need for progenitor scenarios that could explain the diversity among SN~Ia. Several explanations have been suggested, such as variations in the metallicity of the progenitor, in the carbon to oxygen ratio at its center, or in the central density at the time of ignition \citep[e.g., ][]{Timmes2003ApJ, Ropke_etal2006A&A, Lesaffre2006MNRAS}. The variation of the latter two parameters is expected to result from the variation in the initial white dwarf mass and in the accretion history. In this work we follow the evolution of helium stars with initial mass $\approx (2.2 - 2.5)\,M_\odot$ and show that they might reach thermonuclear explosion and perhaps account for some of the peculiar SNe. | 7 | 10 | 0710.3911 |
|
0710 | 0710.5126_arXiv.txt | Our campaign of deep monitoring observations with {\it Chandra} of the nearby elliptical galaxy NGC 3379 has lead to the detection of nine globular cluster (GC) and 53 field low mass X-ray binaries (LMXBs) in the joint {\it Hubble}/{\it Chandra} field of view of this galaxy. Comparing these populations, we find a highly significant lack of GC LMXBs at the low (0.3-8~keV) X-ray luminosities (in the $\sim 10^{36}$ to $\sim 4\times10^{37}$ erg s$^{-1}$ range) probed with our observations. This result conflicts with the proposition that all LMXBs are formed in GCs. This lack of low-luminosity sources in GCs is consistent with continuous LMXB formation due to stellar interactions and with the transition from persistent to transient LMXBs. The observed cut-off X-ray luminosity favors a predominance of LMXBs with main-sequence donors instead of ultra-compact binaries with white-dwarf donors; ultra-compacts could contribute significantly only if their disks are not affected by X-ray irradiation. Our results suggest that current theories of magnetic stellar wind braking may work rather better for the unevolved companions of GC LMXBs than for field LMXBs and cataclysmic variables in the Galaxy, where these companions may be somewhat evolved. | Since their discovery in the Milky Way (see Giacconi 1974), the origin and evolution of Low-mass X-ray binaries (LMXBs) has been the subject of much discussion. LMXBs are found in both the stellar field and globular clusters (GCs). Their incidence per unit stellar mass is much higher in GCs than in the field, requiring a special formation mechanism, presumably dynamical (Clark 1975; Katz 1975). The high efficiency of these dynamical processes led to the suggestion that all LMXBs may form in GCs and then disperse in the field (see e.g., Grindlay 1984; Grindlay \& Hertz 1985); however, some primordial field binaries are also expected to evolve into LMXBs, suggesting that there may be two populations of these sources, with distinct formation and evolutionary histories (see review by Verbunt \& van den Heuvel 1995). The discovery of X-ray source populations in early-type galaxies with {\it Chandra} has provided a wider and more diverse observational basis for the study of LMXBs, their association with GCs and the role that GCs may play in LMXB formation (see review Fabbiano 2006). However, until recently, only the most luminous extra-Galactic LMXBs with ($\sim$0.3-8~keV) $L_X \geq$ a few $10^{37}$erg s$^{-1}$ have been observed, and therefore the study of the GC-LMXB association has been limited to systems with X-ray luminosities in the upper range of Galactic LMXB luminosities. With our deep 337 ks {\it Chandra} ACIS-S3 observations of the unperturbed elliptical galaxy NGC 3379 in the nearby poor group Leo (D$\sim$11 Mpc), we can now pursue this study in a luminosity range more typical of the well-studied Galactic LMXBs. In NGC 3379, The GC-LMXB association has been previously studied by Kundu, Maccarone \& Zepf (2007, hereafter KMZ), using the first shorter {\it Chandra} observation of this galaxy, which has a typical source detection threshold of $\sim 1-2 \times 10^{37}$ erg s$^{-1}$ (KMZ). Our data set is $\sim$10 times deeper (337 ks), allowing the detection of LMXBs at luminosities of $\sim 10^{36}$ erg s$^{-1}$. | Our campaign of deep monitoring observations of the nearby elliptical galaxy NGC 3379 with {\it Chandra} ACIS-S3 has led to the detection of nine GC LMXBs in the field studied optically with {\it Hubble} WFPC2, seven of which were previously detected in a much shorter {\it Chandra} exposure (1/10 of the exposure time; KMZ). The comparison of GC and field LMXB statistics in the joint {\it Hubble}-{\it Chandra} field of view demonstrates a relative lack of GC LMXB at luminosities below $\sim 4\times10^{37}$ erg s$^{-1}$; field and GC LMXBs instead are know to have closely matching XLFs above this luminosity (Kim E. et al 2006; KMZ). The dearth of low-luminosity GC LMXBs in NGC 3379 is consistent with a similar suggested behavior in NGC 3115 (KMZ), and with the XLFs of field and GC LMXBs in the Milky Way and M31 (Voss \& Gilfanov 2007). These differences between low-luminosity field and GC XLFs falsify suggested theories that {\it all} LMXBs may have been originated in GCs. The luminosity-dependent differences of field and GC XLFs cannot be explained by high luminosity GCs containing multiple LMXBs, because we find clear evidence of source variability for seven out of the nine sources invalidating this hypothesis. Persistent behavior of high luminosity GC sources, compared with transient field sources of similar high luminosity may explain the discrepancy as excess of luminous GC LMXBs. However, the detection of three candidate transient source (with peak luminosity greater than $10^{37}$ erg s$^{-1}$) in the GC LMXB population of NGC 3379 may not support this explanation. The lack of low-luminosity sources in GCs is consistent with the prediction of Bildsten and Deloye (BD4) based on the transition of sources from persistent to transients due to the thermal disk instability. However, the value of the observed XLF cut-off is not consistent with their suggestion that GC LMXBs are dominated by ultra-compact binaries, but instead favors LMXBs with H-rich MS donors. The $\sim 4\times10^{37}$ erg s$^{-1}$ luminosity cut-off is also consistent with current theories of magnetic stellar wind braking, suggesting that this effect may work rather better for the unevolved companions of GC LMXBs than for field LMXBs and cataclysmic variables in the Galaxy, where these companions may be somewhat evolved. While our results firmly establish a dearth of GC sources in NGC3379 at low luminosity, the accurate luminosity (and uncertainty) of the GC XLF cut-off will need the formal analysis of the NGC3379 LMXB luminosity function (Kim et al in preparation). The forthcoming analysis of the very deep {\it Chandra} observations of NGC4278, a GC-rich galaxy (see Kim, D.-W. et al 2006), will provide in the near future stronger constraints on the potential 'universality' and value of the GC LMXB XLF cut-off luminosity. | 7 | 10 | 0710.5126 |
0710 | 0710.0968_arXiv.txt | {Gamma-ray binaries have been established as a new class of sources of very high energy (VHE, $>$100~GeV) photons. These binaries are composed of a massive star and a compact object. The gamma-rays are probably produced by inverse Compton scattering of the stellar light by VHE electrons accelerated in the vicinity of the compact object. The VHE emission from LS~5039 displays an orbital modulation. } {The inverse Compton spectrum depends on the angle between the incoming and outgoing photon in the rest frame of the electron. Since the angle at which an observer sees the star and electrons changes with the orbit, a phase dependence of the spectrum is expected.} {A procedure to compute anisotropic inverse Compton emission is explained and applied to the case of \ls. The spectrum is calculated assuming the continuous injection of electrons close to the compact object: the shape of the steady-state distribution depends on the injected power-law and on the magnetic field intensity.} {Compared to the isotropic approximation, anisotropic scattering produces harder and fainter emission at inferior conjunction, crucially at a time when attenuation due to pair production of the VHE gamma-rays on star light is minimum. The computed lightcurve and spectra are very good fits to the HESS and EGRET observations, except at phases of maximum attenuation where pair cascade emission may be significant for HESS. Detailed predictions are made for a modulation in the GLAST energy range. The magnetic field intensity at periastron is 0.8$\pm0.2$~G. } {The anisotropy in inverse Compton scattering plays a major role in \ls. A simple model reproduces the observations, constraining the magnetic field intensity and injection spectrum. The comparison with observations, the derived magnetic field intensity, injection energy and slope suggest emission from a rotation-powered pulsar wind nebula. These results confirm gamma-ray binaries as promising sources to study the environment of pulsars on small scales.} | Gamma-ray binaries have been established in the past couple of years as a new class of sources of very high energy (VHE, $>$100~GeV) photons. They are characterized by a large gamma-ray luminosity above an MeV, at the level of or exceeding their X-ray luminosity. At present, all three such systems known (\object{\ls}, \object{\psrb}\ and \object{\lsi}, recently possibly joined by \object{Cyg X-1}) comprise a massive star \citep{2005Sci...309..746A,2005A&A...442....1A,2006Sci...312.1771A,2007arXiv0706.1505M}. The compact object in \psrb\ is a 48-ms, young radio pulsar. The VHE emission arises from the interaction of the relativistic wind from this pulsar, extracting rotational energy from the neutron star, with the stellar wind from its companion \citep{1994ApJ...433L..37T}. Particles gain energy at the shock between the winds, resulting in a small-scale pulsar wind nebula \citep{1981MNRAS.194P...1M}. The particles radiate away their energy as they are entrained in the shocked flow, forming a comet-like trail of emission behind the pulsar \citep{2006A&A...456..801D}. The nature of the compact object and origin of the VHE emission remains controversial in \ls\ and \lsi, although recent observations indicate the radio emission of \lsi\ behaves like the comet tail expected in the pulsar scenario \citep{2006smqw.confE..52D}. Alternatively, the VHE emission could originate from particles accelerated in a relativistic jet, the energy source being accretion onto a black hole or neutron star \citep{2006ApJ...643.1081D,2006A&A...451..259P}. The rationale being that there is evidence for particle acceleration in the jets of microquasars and active galactic nuclei. However, hard evidence for accretion occuring in either \ls\ or \lsi\ has been hard to come by \citep[e.g.][]{2005A&A...430..245M} and the similarities between the three systems (and differences with the usual microquasars) do not argue in favour of the accretion/ejection scenario \citep{2006A&A...456..801D}. Regardless of the actual powering mechanism, some particles must be accelerated to high energies to generate the VHE gamma-rays. If these particles are leptons, the only viable gamma-ray radiation mechanism is inverse Compton scattering on the stellar photons. The massive stars in gamma-ray binaries have effective temperatures of several tens of thousand K and radii of about 10~$R_\odot$, yielding luminosities of the order of $10^{39}$~erg~s$^{-1}$. This provides a huge density of stellar photons in the UV band that VHE leptons may up-scatter, much greater than any other possible source of target photons (e.g. synchrotron or bremsstrahlung emission). The emitted VHE photons also have enough energy to produce $e^+e^-$ pairs with the UV stellar photons. Most of the VHE flux may therefore be lost to the observer if the source is behind the star and VHE photons have to travel through the stellar light. Gamma-ray attenuation has been shown to lead to a modulation of the VHE flux with minimum absorption (maximum) at inferior (superior) conjunction \citep{2005ApJ...634L..81B,2006A&A...451....9D}. HESS observations have indeed shown a stable modulation of the VHE flux from \ls\ on the orbital period with a maximum around inferior conjunction. This suggests attenuation plays a role and that the source of VHE gamma-rays cannot be more than about an AU from the binary (or attenuation would be too weak to modulate the flux). However, a non-zero flux is detected at superior conjunction where a large attenuation is expected, possibly because of pair cascading. Moreover, the spectral changes that are reported do not fit with an interpretation based on pure attenuation of a constant VHE source spectrum \citep{2006A&A...460..743A}. Inverse Compton scattering also has a well-known dependence on the angle $\Theta$ between incoming and outgoing photon. The photon flux from the star being anisotropic, the resulting inverse Compton emission will depend on the angle at which it is observed. Hence, a phase-dependent VHE spectrum will be observed even if the distribution of particles is isotropic and remains constant throughout. This effect has previously been investigated in \psrb\ by \citet{2000APh....12..335B} who calculated the radiative drag on the unshocked pulsar wind from scattering of stellar light, using results from \citet{1989ApJ...343..277H}. The drag produces a Compton gamma-ray line with a strong dependence on viewing angle. This work purports to explain the HESS observations of LS~5039 using a combination of anisotropic inverse Compton scattering and attenuation in the simplest way possible. The aim is to constrain the underlying particle distribution and/or powering mechanism. \S2 derives the main equations governing anisotropic Compton scattering in the context of gamma-ray binaries and discusses the principal characteristics to expect. \S3 presents the application to the case of LS~5039. The lightcurve and spectra observed by the HESS collaboration are reproduced by a model taking into account the photon field anisotropy and the attenuation due to pair creation. \S4 concludes on the origin of the VHE emission from this system. | The anisotropic behaviour of inverse Compton scattering has a major influence on the emission from gamma-ray binaries. In these sources, the massive star provides a large source of seed photons with energies around an electron-volt. If high energy electrons are accelerated in the vicinity of the compact object, then the angle between the star, compact object and observer changes with orbital phase. The variation in viewing angle leads to a strong modulation in both the intensity and spectral shape of the scattered radiation. Scattering stellar photons to the TeV range requires very energetic electrons with Lorentz factors $\gamma_{\rm e}\approx 10^6-10^7$. The scattering therefore occurs in the Klein-Nishina regime. In this case, the anisotropy results, at inferior conjunction, in a harder and fainter spectrum than predicted using an isotropic approximation for the incoming photons. Crucially, inferior conjunction also corresponds to the phase at which the produced VHE gamma-rays are less attenuated by pair production on stellar photons. At other phases the emitted spectrum is close to the one obtained using the isotropic photon field approximation and can be severely attenuated by pair production. The result is a complex interplay that reduces the amplitude of the variations expected from a pure attenuation model and a hardening at inferior conjunction. The \ls\ lightcurve and spectra were modelled using a simple-minded leptonic model. The electrons are assumed to be accelerated efficiently in a small zone in the vicinity of the compact object with a standard $\gamma_{\rm e}^{-p}$ power-law. Radiative losses due to inverse Compton emission and synchrotron emission generate a distinctive steady-state electron distribution in this environment dominated by stellar photons. The distribution has a prominent hardening between the energy at which inverse Compton losses enter the Klein-Nishina regime ($\gamma_{\rm KN}\approx 6~10^4$ in \ls) and the energy at which synchrotron losses take over ($\gamma_{\rm S}\approx 10^7$ for a 1~G field). This is for instance the distribution found in the vicinity of the pulsar wind shock but it applies equally well to any leptonic model where particles are accelerated close to the compact object. The magnetic field was allowed to vary as the inverse of the orbital separation, as expected from a pulsar wind nebula. The model has only three parameters: the intensity of the magnetic field, the normalization of the electron distribution and the slope $p$ of the injected power-law $\gamma_{\rm e}^{-p}$. The cutoff in the very high energy gamma-ray spectrum is very sensitive to the magnetic field intensity, via the location of $\gamma_{\rm S}$ in the electron distribution. Fitting the high-state spectrum seen by HESS gives a rather constrained magnetic field intensity at periastron of 0.8$\pm0.2$~G. This value compares well with the values found using simple pulsar wind models which give 5 $(\dot{E}_{36} \sigma_3)^{1/2} R_{11}^{-1}$~G, where $\dot{E}_{36}$ is the pulsar spindown power in units of 10$^{36}$~erg~s$^{1}$, $\sigma_3$ is the ratio of magnetic to kinetic energy in the pulsar wind in units of 10$^{-3}$ and $R_{11}$ is the distance of the shock to the pulsar in units of 10$^{11}$~cm. Fitting the HESS high-state spectrum also sets the injection slope to $p=2\pm0.3$, close to the canonical value for shock acceleration. The normalization of the electron distribution implies an injection rate of 10$^{36}$~erg~s$^{-1}$ for a radiative zone of 3~$10^{11}$~cm. These results are remarkably consistent with the expectations for a pulsar wind model. The spectrum is also found to fit extremely well the EGRET observations, adding credence to the reliability of this simple approach. The model predicts a strong variation in the GLAST band with a softening from high to low flux below a GeV (where synchrotron emission dominates the spectrum) but a hardening above a GeV (where inverse Compton emission dominates the spectrum). The HESS low-state spectrum is not explained to satisfaction. The model fits nicely the EGRET measurements but produces too many gamma-rays at 5-10~TeV. A possible solution is a more complex orbital phase-dependence of the electron distribution at selected phases. Another solution is that the low-state spectrum corresponds to phases of strong attenuation and that emission from the created pairs contribute significantly to the spectrum. Additional HESS observations near minimum flux would be welcomed. The orbital modulation of the HESS emission is easily reproduced. A well-defined peak is predicted between phases 0.7-0.9 for which evidence may already be seen in the data. The lightcurve at GLAST energies is anti-correlated with the HESS lightcurve and has a peak at periastron, where the stellar photon density is maximum, and a minimum at inferior conjunction because of the anisotropic effects in inverse Compton scattering. The GLAST spectrum below 1~GeV should be influenced by the tail of the synchrotron emission from the highest energy electrons. The peak synchrotron emission is at about $100$~MeV for maximally accelerated electrons, regardless of magnetic field. Hence, if this component is detected, it will provide evidence that electrons are indeed accelerated with extreme efficiency in this source. Similar results for the magnetic field intensity and particle energy are found when a lower inclination is used, i.e. implying a black hole compact object rather than a neutron star. In this case, the emission is thought to arise from a relativistic jet powered by accretion onto the black hole. Within the assumptions of this work on the particle distribution, it is difficult to argue that a significant part of the emission occurs far along a jet since this does not naturally reproduce neither the spectrum nor the lightcurve measured by HESS. Most of the emission should still occur close to the compact object. However, unlike in the case of a pulsar wind nebula, there is no independent theoretical expectations in support of the magnetic field intensity (certainly smaller than its equipartition value in the accretion flow) and particle energy that are derived. Therefore, the pulsar wind nebula model appears favoured independently of other possible considerations. Despite the complexity of the phenomena involved in pulsar wind nebula emission, it is found that the peculiar environment of a gamma-ray binary, most prominently the enormous luminosity of the massive companion, severely constrains the number of degrees-of-freedom in the model. A simple model suffices to reproduce most of the observations. The value of the magnetic field at the shock is found to be tightly constrained by the HESS observations to 0.8$\pm$0.2~G and the injection spectrum slope to $p=2\pm0.3$. These results confirm that gamma-ray binaries are promising sources to study the environment of pulsars on very small scales. | 7 | 10 | 0710.0968 |
0710 | 0710.3193_arXiv.txt | In this paper a neutron star with an inner core which undergoes a phase transition, which is characterized by conformal degrees of freedom on the phase boundary, is considered. Typical cases of such a phase transition are e.g. quantum Hall effect, superconductivity and superfluidity. Assuming the mechanical stability of this system the effects induced by the conformal degrees of freedom on the phase boundary will be analyzed. We will see that the inclusion of conformal degrees of freedom is not always consistent with the staticity of the phase boundary. Indeed also in the case of mechanical equilibrium there may be the tendency of one phase to swallow the other. Such a shift of the phase boundary would not imply any compression or decompression of the core. By solving the Israel junction conditions for the conformal matter, we have found the range of physical parameters which can guarantee a stable equilibrium of the phase boundary of the neutron star. The relevant parameters turn out to be not only the density difference but also the difference of the slope of the density profiles of the two phases. The values of the parameters which guarantee the stability turn out to be in a phenomenologically reasonable range. For the parameter values where the the phase boundary tends to move, a possible astrophysical consequence related to sudden small changes of the moment of inertia of the star is briefly discussed. | One of the most intriguing features of neutron stars is the possibility of having an inner core which undergoes a (colour) superconductivity and/or superfluidity phase transition due to the extremely high pressure and density. This fact was first pointed out by Migdal \cite{Mi59} (there is a huge amount of literature on this subject with a little hope of providing one with a complete list of references: for two detailed reviews, see\ \cite% {Pet92} \cite{Ra99} and references therein; for general relativistic formalisms suitable for dealing with neutron stars within these scenarios see \cite{CL98} \cite{AC07} and references therein). Such phase transitions inside neutron stars could lead to interesting observational effects relevant in cosmology \cite{Sigl:2006ur} and through the quasi normal modes of the stars as well as the related gravitational radiation (see, for instance, \cite{ST02} \cite{MP03}; detailed reviews on the subjects are \cite% {KS99} \cite{Ste03} \cite{ST03}). Recently, there has been also pointed out the intriguing possibility to have a quantum hall phase of gluons and quarks in the inner core of a neutron star (see, in particular, \cite{IM05} \cite{Iw05}% ). The discovery of \textit{magnetars} \cite{DT92}, highly magnetized neutron stars whose magnetic fields can reach $10^{18}G$ in the core, makes the arising of Quantum Hall features in the inner core of a neutron star not unlikely. All the here described types of phase transitions have in common to be characterized by the presence of conformal degrees of freedom on the phase boundary predicted by QFT\footnote{% The fact that the boundary degrees of freedom are conformal can be seen "heuristically" in the case of a superconducting phase. Due to the Meissner effect, the current must circulate on the phase boundary of the superconductive phase. Phenomenologically, there is no dissipation, this means that there is no physical scale over which the boundary excitations can dissipate energy. The lack of a characteristic scale for the boundary theory is related to conformal symmetry. In the case of Quantum Hall Effect, this can be easily seen from the fact that it is well described by a Chern-Simons action in the bulk. Thus, the theory induced on the boundary is the Wess-Zumino-Witten action which also has conformal symmetry.} (see, for instance, \cite{Wi90} \cite{We96}). The main goal of this paper is to study under which conditions the phase boundary inside a neutron star, which undergoes such a phase transition with conformal boundary degrees of freedom, can be static in general relativity, assuming mechanical equilibrium of the system\footnote{ The analysis of the mechanical stability of neutron stars has been performed in a huge number of papers. To provide one with a complete list of references is a completely hopeless task; representative papers are, for instance, \cite{EKO91} \cite% {ABR00} \cite{CL98} and references therein)}. Indeed also if the system is mechanically stable i.e. there is no pressure inducing the collapse or expansion of the core, there can be another sort of instability: there may be the tendency of one phase to prevail over the other near the boundary. This would imply that the boundary gets shifted. This situation is analogous to a bubble chamber where a small perturbation induces a phase transition. To study this phenomenon one must take into account the contribution of the nontrivial traceless boundary stress tensor to the Einstein field equations. This problem is equivalent to solving Israel's junction conditions. It is necessary to assume a dynamical phase boundary and to check under which conditions such a configuration is a solution of the Einstein field equations (by solving the junction conditions). It turns out that the equation of motion for the dynamical phase boundary is equivalent to the dynamics of a classical point particle in an effective potential. There exist a static stable equilibrium configuration if the effective potential has a local minimum for a suitable negative value of the potential. It is shown that the existence of a local minimum is regulated both by the difference in the mass density between the two phases and by difference of the slope of the two density profiles. A third parameter which determines the form of the potential is the energy of the shell. The non trivial effects of conformal boundary degrees of freedom have been pointed out, in the context of black hole physics, in \cite% {Cao7} in which they are related to the arising (after the event horizon is formed) of Quantum Hall features \cite{Lau99} \cite{Wi90}\ related to the strong attractive nature of the gravitational field acting on Fermions inside a collapsed neutron star. The structure of the paper is the following: in Section 2, the assumptions of the present paper are explained. In Section 3 the junction conditions derived from the Einstein equations for a neutron star with a phase boundary are solved. In section 4 the range of the physical parameters, characterizing our configuration, for which the phase boundary is in stable equilibrium is determined. In section 5 an astrophysical implication is discussed. Section 6 the conclusions and perspectives are drawn. | We have studied a model of a neutron star with an inner core which undergoes a phase transition. It has been shown that, even assuming mechanical equilibrium, it is not always consistent to assume a static phase boundary once conformal boundary degrees of freedom are taken into account. In order to study the staticity of such a phase boundary one must take into account the general relativistic effects of these boundary degrees of freedom by including a nontrivial stress tensor on the junction. To the best of the authors knowledge, such a consistency analysis with conformal boundary degrees of freedom has not been performed previously. The astrophysical consequences of the non-static regime, related to sudden changes of the moment of inertia of the star, are worth further investigating. | 7 | 10 | 0710.3193 |
0710 | 0710.1196_arXiv.txt | HI features near young star clusters in M81 are identified as the photodissociated surfaces of Giant Molecular Clouds (GMCs) from which the young stars have recently formed. The HI column densities of these features show a weak trend, from undetectable values inside $R = 3.7$\ kpc and increasing rapidly to values around $3 \times 10^{21}\ \mbox{cm}^{-2}$ near $R \approx 7.5$\ kpc. This trend is similar to that of the radially-averaged HI distribution in this galaxy, and implies a constant area covering factor of $\approx 0.21$ for GMCs throughout M81. The incident UV fluxes \Gnaught\ of our sample of candidate PDRs decrease radially. A simple equilibrium model of the photodissociation-reformation process connects the observed values of the incident UV flux, the HI column density, and the relative dust content, permitting an independent estimate to be made of the total gas density in the GMC. Within the GMC this gas will be predominantly molecular hydrogen. Volume densities of $1 < n < 200\ \mbox{cm}^{-3}$ are derived, with a geometric mean of $17\ \mbox{cm}^{-3}$. These values are similar to the densities of GMCs in the Galaxy, but somewhat lower than those found earlier for M101 with similar methods. Low values of molecular density in the GMCs of M81 will result in low levels of collisional excitation of the CO(1-0) transition, and are consistent with the very low surface brightness of CO(1-0) emission observed in the disk of M81. | \label{sec:intro} Giant molecular clouds (GMCs) in the interstellar medium (ISM) are generally accepted as the birthplaces of new stars. The most massive of these new stars produce copious amounts of far ultraviolet (FUV) radiation which will, in turn, photodissociate the molecular gas of the parent GMCs, producing ``blankets'' of HI (and other atomic species) on their surfaces. \citet{all1986} were the first to present evidence that major features in the HI distribution of the nearby spiral galaxy M83 (NGC 5236), namely, the inner HI spiral arms observed with the Very Large Array (VLA), were the result of photodissociation of \Htwo on galactic scales. \citet{all1997} confirmed that HI features existed in M81 (NGC 3031) on scales of $\approx 150$ pc which were qualitatively consistent with the expected morphology of large, low density photodissociation regions (PDRs), and explicitly related those HI features to nearby bright sources of Far-UV (FUV) radiation found on images of the galaxy from the Ultraviolet Imaging Telescope (UIT). \citet{smi2000} used a simple, but quantitative, model for the equilibrium physics of photodissociation regions and applied it to VLA-HI and UIT-FUV observations of M101 (NGC 5457). Their work showed that this approach provides an entirely new method for determining the volume density of molecular gas in star-forming GMCs of galaxies, a method which is no longer dependent on the (often poorly known) excitation conditions for line emission by specific molecular tracers. In this paper we return to another spiral galaxy, M81, and carry out a quantitative analysis of the GMCs in this galaxy using the methods described by \citet{smi2000}. This new analysis has been made possible by new data on M81 which was not available to \citet{all1997}; new, higher-resolution, high-sensitivity VLA-HI observations, and sensitive new FUV imagery from the GALEX satellite. We have also sought an independent verification that the very basis of our approach is valid, namely, that the HI in the immediate vicinity of FUV concentrations is indeed produced in PDRs. To this end we have examined the Spitzer/IRAC data on M81 for evidence of mid-IR emission by polycyclic aromatic hydrocarbons (PAHs) which are also thought to be tracers of PDRs. M81 has been probed extensively for the molecular tracer CO. No global CO map of M81 has been published to this date, and the detected CO emission is found to be very weak and spotty (see e.g. \citet{kna2006} and references therein). Since PDRs also produce CO emission \citep[e.g.][]{all2004}, comparing the CO results to the results in this paper is therefore of considerable interest. The outline of this paper is as follows: Section \ref{sec:data} contains a description of the data we used, followed by a brief theoretical description of PDRs in Section \ref{sec:theory} and the application of the method in Section \ref{sec:method}. The results are presented in Section \ref{sec:results}. The results are discussed and the conclusions are briefly summarized in Section \ref{sec:discussionconclusions}. | \label{sec:discussionconclusions} The candidate PDRs in M81, which were selected on the basis of their FUV emission, seem to fit the photodissociation model well. Our results show no systematically different properties of the parent GMCs in different parts of M81. The total hydrogen volume density is roughly constant, even as the underlying HI, FUV and dust-to-gas ratio vary. The cloud densities we find are lower than the range of values (30 - 1000 cm$^{-3}$) found in M101 \citep{smi2000}. The observed values of $N_{HI}$ of individual PDRs are similar to those seen in M101. The observed HI columns in both galaxies abruptly increase at the same normalized radius (0.3) and appear to decline somewhat beyond 0.7, assuming an $R_{25}$ of 30.3 kpc for M101 \citep{vil1992}. The range of observed HI columns is also the same. The downward trend in \Gnaught\ also is consistent with \citet{smi2000}, when no internal extinction correction is applied to their data. Figure \ref{fig:Rnorm_n} shows that the densities of the GMCs in M81 do not appear to change with galactocentric radius, consistent with the M101 results. \citet{kau1999} modeled the expected CO intensity from a PDR for a range of incident FUV fluxes and cloud densities. The range of our results would be consistent with modeled CO intensities below $5 \times 10^{-8}~ \mbox{ergs cm}^{-2}~ \mbox{s}^{-1}~ \mbox{sr}^{-1}$, or $1 \times 10^{-8}~ \mbox{ergs cm}^{-2}~ \mbox{s}^{-1}~ \mbox{sr}^{-1}$ for the vast majority of sources (6.4 K km s$^{-1}$). Figure 1 in \citet{all2004} shows that for the range of values in Figure \ref{fig:selection}, the modeled CO intensities are independent of $n$. The low volume densities we find are consistent with a lack of CO emission in M81 as discussed by \citet{kna2006}, and which those authors attribute to insufficient excitation. We note that the volume number densities of colliding \Htwo\ molecules in the GMCs of M81 is 1/2 of the values for $n$ calculated here. The mean value we have found for $n$ in the GMCs of M81 translates into a mean number density for $n_{H_2} \approx 10 \mbox{cm}^{-3}$, well below the values required for collisional excitation of CO molecules. In this case, the CO emission in M81 is in general \textit{subthermal}. Their reported upper limit for $I_{CO}$ is 1.03 K km s$^{-1}$ for the regions they investigated, near our source no. 18. The \Gnaught\ and $n$ that we find there are consistent to that value of $I_{CO}$ (and somewhat higher) per Figure 1 in \citet{all2004}, when beam dilution is taken into account. The low CO emission in M81 is also explored in \citet{cas2007}, who again point to a lack of excitation of the molecular gas. They find no molecular gas in the nucleus of M81. The absence of FUV sources to excite the gas is consistent with that finding. After accounting for observational and projection effects (see the previous Section), we note that in the nearby galaxies in general on which our method is applicable, it is most sensitive to a combination of low \Gnaught\ and low $n$. Summarizing, our conclusions are: \begin{itemize} \item We selected a number of discrete FUV sources in M81, which we consider to be potential PDRs on the surfaces of the parent GMCs. \item The total hydrogen volume densities of GMCs close to clusters of young stars in M81 are in the range of 1 $< n <$ 200 cm$^{-3}$ with a geometric mean of 17 cm$^{-3}$. This is approximately ten times lower than GMCs in M101 studied with the same method. \item The low GMC volume densities are consistent with a lack of CO emission in M81. \item M81's GMCs have a filling factor of $\approx$ 0.21 within $R_{25}$. \item No candidate PDRs are found in M81 within $R < 0.3 R_{25}$. \item We have provided a thorough analysis of the observational selection effects on our results and conclude that, while such effects are (necessarily) present in our results, our main conclusions as to the range and values of the total volume densities of GMCs in M81 are not affected. \item The presence of PAH emission in the neighborhood of our candidate PDRs lends support to our view that the HI patches near FUV sources are indeed produced by photodissociation. PAH emission occurs near almost all UV sources. PAH and HI emission coincide in more than half of our sources. \end{itemize} | 7 | 10 | 0710.1196 |
0710 | 0710.1475_arXiv.txt | Gamma ray bursts have been divided into two classes, long-soft gamma ray burst and short-hard gamma ray burst according to the bimodal distribution in duration time. Due to the harder spectrum and the lack of afterglows of short-hard bursts in optical and radio observations, different progenitors for short-hard bursts and long-soft bursts have been suggested. Based on the X-ray afterglow observation and the cumulative red-shift distribution of short-hard bursts, \citet{Nak06} found that the progenitors of short-hard bursts are consistent with old populations, such as mergers of binary neutron stars. Recently, the existence of two subclasses in long-soft bursts has been suggested after considering multiple characteristics of gamma-ray bursts, including fluences and the duration time. In this work, we extended the analysis of cumulative red-shift distribution to two possible subclasses in L-GRBs. We found that two possible subclass GRBs show different red-shift distributions, especially for red-shifts $z > 1$. Our results indicate that the accumulative red-shift distribution can be used as a tool to constrain the progenitor characteristics of possible subclasses in L-GRBs. | Traditionally, gamma-ray bursts (denoted as GRBs) have been divided into two classes which are separated by the duration time of gamma-ray bursts $T_{90}=2$ sec \citep{Kou93}. Since the short-duration gamma-ray bursts had relatively hard observed spectra than long-duration gamma-ray bursts, two classes are usually called as long-soft gamma ray burst (denoted as L-GRB) and short-hard gamma-ray burst (denoted as SHB). From the afterglow observations of L-GRBs, the association between the L-GRBs and supernovae/hypernovae has been suggested. These observations suggest that the progenitors of L-GRBs are massive giant stars which are collapsing. One of the best candidate is the Collapsar model \citep{Woo93,Mac99} in which rapidly rotating black hole with Kerr parameter about 0.8 is formed after GRB. Recently, spinning black holes with high Kerr parameters, $a_\star \approx 0.8$, have been observed in soft X-ray black hole binaries \citep{McC06,Sha06}. These observations support the idea that the rapidly rotating black holes in soft X-ray black hole binaries are the remnants of L-GRBs \citep{Bro00,Lee02}, providing the natural sources for Collapsars. Due to the differences in the observational characteristics of SHBs compared to those of L-GRBs, such as lack of afterglow observations in optical and radio band and harder spectrum, etc., different progenitors for SHBs have been suggested. Mergers of binary neutron stars \citep{Eic89,Nar92} are among the best candidates. After recent X-ray afterglow observations by Swift and HETE-II, the associations between SHBs and the host galaxies were reported (see Table~\ref{tab1} and references there in). Based on the X-ray afterglow observation of SHBs and the cumulative red-shift distribution of SHBs, \citet{Nak06} found that the progenitors of SHBs are consistent with old populations. This finding supports mergers of binary neutron stars as progenitors of SHBs. Recently, the existence of two subclasses in L-GRBs has been suggested after considering both fluences and the duration time \citep{Cha07}. They divided L-GRBs into Clusters II and III which were separated by the fluences combined with the duration time. They classified L-GRBs with higher fluences as Cluster III and suggested that their progenitors are the massive collapsing stars, such as Collapsar \citep{Woo93,Mac99}. As a possible origin of Cluster II GRBs, they suggested neutron-star, white-dwarf binaries. If there exist two subclasses in L-GRBs as \citet{Cha07} suggested and only Cluster III GRBs are associated with massive collapsing giants which are the progenitors of supernovae or hypernovae, the observed red-shift distribution of two subclasses might show different behaviour. In this work, we extended the work of \citet{Nak06} to two possible subclasses in L-GRBs. We found that Clustters II and III show different red-shift distribution at high red-shift $z>1$. Our results indicate that the accumulative red-shift distribution can be used as a tool to constrain the lifetimes of possible subclasses in GRBs. In Sec.~\ref{sec-rate}, the numerical methods which we used in our analysis are summarized \citep{Nak06}. The estimated lifetime of SHBs are summarized in Sec.~\ref{sec-SHB}. We confirmed that the progenitors of SHBs are consistent with old population with lifetime $\tau_\ast=6.5$ Gyr, indicating that SHBs are from the mergers of binary neutron stars. In Sec.~\ref{sec-LGRB}, the cumulative red-shift distributions of two possible subclasses in L-GRB have been summarized. Our results show that there are clear differences in the red-shift distributions of two possible subclasses in L-GRBs for the red-shift $z>1$. This suggests that the cumulative red-shift distribution can be used as a tool to distinguish the subclasses in L-GRBs. Our final conclusion follows in Sec.~\ref{sec-con}. | \label{sec-con} In this work, by extending the work of Nakar et al. \citep{Nak06,Nak07}, we investigated the constraints on the progenitor lifetime of SHBs. We confirmed that the SHBs are consistent with old population with mean time $\tau_\ast=6.5$ Gyr. We also confirmed that this conclusion is independent of the details of the cosmology models. This results support the conclusion that the progenitors of SHBs are consistent with old population, such as merging compact star binaries (neutron star--neutron star binaries or neutron star--black hole binaries) \citep{Nak06,Nak07,Bet07,Lee07}. We also investigated two subclasses Cluster II and III L-GRBs \citep{Cha07} using the same method of \citet{Nak06}. We found that the power-law distribution is more consistent with Cluster III GRBs than with Cluster II GRBs. We also found that the cumulative red-shift distribution of two clusters show clear difference in the red-shift region $z>1$. This result supports the existence of two subclasses in L-GRBs. However, in order to draw firm conclusions, many effects which are not included in our analysis, e.g., the connection between GRBs and the metallicity of the host galaxy, different characteristics of cosmological GRBs and subluminous GRBs, observability premium for different type of GRB progenitors, etc., have to be considered. | 7 | 10 | 0710.1475 |
0710 | 0710.1828_arXiv.txt | We present high signal-to-noise spectrophotometric observations of seven luminous \HII\ galaxies. The observations have been made with the use of a double-arm spectrograph which provides spectra with a wide wavelength coverage, from 3400 to 10400\,\AA\ free of second order effects, of exactly the same region of a given galaxy. These observations are analysed applying a methodology designed to obtain accurate elemental abundances of oxygen, sulphur, nitrogen, neon, argon and iron in the ionized gas. Four electron temperatures and one electron density are derived from the observed forbidden line ratios using the five-level atom approximation. For our best objects errors of 1\% in t$_e$([O{\sc iii}]), 3\% in t$_e$([O{\sc ii}]) and 5\% in t$_e$([S{\sc iii}]) are achieved with a resulting accuracy of 7\% in total oxygen abundances, O/H. The ionisation structure of the nebulae can be mapped by the theoretical oxygen and sulphur ionic ratios, on the one side, and the corresponding observed emission line ratios, on the other -- the $\eta$ and $\eta$' plots --. The combination of both is shown to provide a means to test photo-ionisation model sequences currently applied to derive elemental abundances in \HII\ galaxies. | When studying evolution two types of ages should be distinguished: the chronological and the evolutionary ages. In the case of galaxies, estimates of the chronological age can be obtained analyzing, for example, the age distribution of their stellar population while the evolutionary age can be estimated from, for example, the metal content of their interstellar medium. \HII\ galaxies, the subclass of Blue Compact Dwarf galaxies (BCDs) which show spectra with strong emission lines similar to those of giant extragalactic \HII\ regions \citep[GEHRs;][]{1970ApJ...162L.155S,1980ApJ...240...41F}, have the lowest metal content of any starforming galaxy suggesting that they are among the youngest or less evolved galaxies known \citep{2007ApJ...654..226R, 1972ApJ...173...25S}. After the findings that a considerable number of the objects observed at intermediate and high redshifts seem to have properties similar to the \HII\ galaxies we know in the Local Universe, it has been suggested that these objects might have been very common in the past and some of them may have evolved to other kind of objects \citep{1995ApJ...440L..49K}. In order to detect these evolutionary effects we need to compare the properties of \HII\ galaxies both in the Local Universe and at higher redshifts. We therefore need to know the true distribution functions of their properties among which the chemical abundances are of the greatest relevance. Spectrophotometry of bright \HII\ galaxies in the Local Universe allows the determination of abundances from methods that rely on the measurement of emission line intensities and atomic physics. This is referred to as the "direct" method. In the case of more distant or intrinsically fainter galaxies, the low signal-to-noise obtained with current telescopes precludes the application of this method and empirical ones based on the strongest emission lines are required. The fundamental basis of these empirical methods is reasonably well understood \citep[see e.g.][]{2005MNRAS.361.1063P}. The accuracy of the results however depends on the goodness of their calibration which in turn depends on a well sampled set of precisely derived abundances by the "direct" method so that interpolation procedures are reliable. Enlarging the calibration range is also important since, at any rate, empirically obtained relations should never be used outside their calibration validity range. The precise derivation of elemental abundances however is not a straightforward matter. Firstly, accurate measurements of the emission lines are needed. Secondly, a certain knowledge of the ionisation structure of the region is required in order to derive ionic abundances of the different elements and in some cases photoionisation models are needed to correct for unseen ionisation states. An accurate diagnostic requires the measurement of faint auroral lines covering a wide spectral range and their accurate (better than 5\%) ratios to Balmer recombination lines. These faint lines are usually about 1\% of the H$\beta$ intensity. The spectral range must include from the UV [O{\sc ii}]\,$\lambda$\,3727\,\AA\ doublet, to the near IR [S{\sc iii}] $\lambda\lambda$\,9069,9532\,\AA\ lines. This allows the derivation of the different line temperatures: T$_e$([O{\sc ii}]), T$_e$([S{\sc ii}]), T$_e$([O{\sc iii}]), T$_e$([S{\sc iii}]), T$_e$([N{\sc ii}]), needed in order to study the temperature and ionisation structure of each \HII\ galaxy considered as a multizone ionised region. Unfortunately most of the available \HII\ galaxy spectra have only a restricted wavelength range (usually from about 3600 to 7000\,\AA), consequence of observations with single arm spectrographs, and do not have the adequate S/N to accurately measure the intensities of the weak diagnostic emission lines. Even the Sloan Digital Sky Survey (SDSS; Stoughton et al.\ 2002) \nocite{2002AJ....123..485S} spectra do not cover simultaneously the 3727 [O{\sc ii}] and the 9069 [S{\sc iii}] lines, they only represent an average inside a 3\,arcsec fibre and reach the required S/N only for the brightest objects. It is important to realise that the combination of accurate spectrophotometry and wide spectral coverage cannot be achieved using single arm spectrographs where, in order to reach the necessary spectral resolution, the wavelength range must be split into several independent observations. In those cases, the quality of the spectrophotometry is at best doubtful mainly because the different spectral ranges are not observed simultaneously. This problem applies to both objects and calibrators. Furthermore one can never be sure of observing exactly the same region of the nebula in each spectral range. To avoid all these problems the use of double arm spectrographs is required. In this work we present simultaneous blue and red observations obtained with the double arm TWIN spectrograph at the 3.5m telescope of the Spanish-German Observatory of Calar Alto. These data are of a sufficient quality as to allow the detection and measurement of several temperature sensitive lines and add to the still scarce base of precisely derived abundances. In the next section we describe some details regarding the selection of the sample as well as the observations and data reduction. The results are presented in section 3. Sections 4 and 5 are devoted to the analysis of these results which are compared with previous data in section 6. Section 7 is devoted to the discussion of our results and finally, our conclusions are summarized in section 8. \begin{table*} \centering \caption[]{Journal of observations.} \label{jour} \begin{tabular} {@{}c c c c c c} \hline \multicolumn{1}{c}{Object ID} & spSpec SDSS & hereafter ID & Date & Exposure (s) & Seeing (\arcsec)\\ \hline \uno & spSpec-52790-1351-474 & \unoc & 2006 June 25 & 5 $\times$ 1800 & 0.9-1.2 \\ \dos & spSpec-52721-1050-274 & \dosc & 2006 June 23 & 4 $\times$ 1800 & 0.8-1.1 \\ \tres & spSpec-52765-1293-580 & \tresc & 2006 June 22 & 4 $\times$ 1800 & 0.8-1.2 \\ \cuatro & spSpec-52072-0617-464 & \cuatroc & 2006 June 24 & 6 $\times$ 1800 & 1.0-1.4 \\ \cinco & spSpec-52377-0624-361 & \cincoc & 2006 June 23 & 5 $\times$ 1800 & 0.8-1.1 \\ \seis & spSpec-52791-1176-591 & \seisc & 2006 June 25 & 5 $\times$ 1800 & 0.9-1.2 \\ \siete & spSpec-51818-0358-472 & \sietec & 2006 June 22 & 5 $\times$ 1800 & 0.8-1.1 \\ \hline \end{tabular} \end{table*} \begin{table*} \centering \caption[]{Right ascension, declination, redshift and SDSS photometric magnitudes obtained using the SDSS explore tools$^a$.} \label{obj} \begin{tabular} {l c c c c r r c c} \hline \multicolumn{1}{c}{Object ID} & \multicolumn{1}{c}{RA} & \multicolumn{1}{c}{Dec} & redshift & \multicolumn{1}{c}{u} & \multicolumn{1}{c}{g} & \multicolumn{1}{c}{r} & \multicolumn{1}{c}{i} & \multicolumn{1}{c}{z} \\ \hline \unoc & 14$^h$\,55$^m$\,06\fs06 & 38$^{\circ}$\,08\arcmin\,16\farcs67 & 0.028 & 18.25 & 17.57 & 17.98 & 18.23 & 18.18 \\ \dosc & 15$^h$\,09$^m$\,09\fs03 & 45$^{\circ}$\,43\arcmin\,08\farcs88 & 0.048 & 18.57 & 17.72 & 18.19 & 17.87 & 17.94 \\ \tresc & 15$^h$\,28$^m$\,17\fs18 & 39$^{\circ}$\,56\arcmin\,50\farcs43 & 0.064 & 18.54 & 17.88 & 18.17 & 17.52 & 17.99 \\ \cuatroc & 15$^h$\,40$^m$\,54\fs31 & 56$^{\circ}$\,51\arcmin\,38\farcs98 & 0.011 & 19.11 & 18.91 & 18.97 & 19.53 & 19.46 \\ \cincoc & 16$^h$\,16$^m$\,23\fs53 & 47$^{\circ}$\,02\arcmin\,02\farcs36 & 0.002 & 16.84 & 16.45 & 16.77 & 17.35 & 17.43 \\ \seisc & 16$^h$\,57$^m$\,12\fs75 & 32$^{\circ}$\,11\arcmin\,41\farcs42 & 0.038 & 17.63 & 17.03 & 17.27 & 17.15 & 17.15 \\ \sietec & 17$^h$\,29$^m$\,06\fs56 & 56$^{\circ}$\,53\arcmin\,19\farcs40 & 0.016 & 18.05 & 17.26 & 17.21 & 17.38 & 17.24 \\ \hline \multicolumn{9}{l}{$^a$http://cas.sdss.org/astro/en/tools/explore/obj.asp} \end{tabular} \end{table*} | We have performed a detailed analysis of newly obtained spectra of seven \HII\ galaxies selected from the Sloan Digital Sky Survey Data Release 3. The spectra cover from 3400 to 10400 \AA\ in wavelength at a FWHM resolution of about 2000 in the blue and 1500 in the red spectral regions. The high signal-to-noise ratio of the obtained spectra allows the measurement of four line electron temperatures: T$_e$([O{\sc iii}]), T$_e$([S{\sc iii}]), T$_e$([O{\sc ii}]) and T$_e$([S{\sc ii}]), for all the objects of the sample with the addition of T$_e$([N{\sc ii}]) for one of the objects. These measurements and a careful and realistic treatment of the observational errors yield total oxygen abundances with accuracies between 7 and 12\%. The fractional error is as low as 1\% for the ionic O$ ^{2+} $/H$ ^{+} $ ratio due to the small errors associated with the measurement of the strong nebular lines of [O{\sc iii}] and the derived T$_e$([O{\sc iii}]), but increases to up to 30\% for the O$^{+}$/H$^{+}$ ratio. The accuracies are lower in the case of the abundances of sulphur (of the order of 25\% for S$^+$ and 15\% for S$^{2+}$) due to the presence of larger observational errors both in the measured line fluxes and the derived electron temperatures. The error for the total abundance of sulphur is also larger than in the case of oxygen (between 15\% and 20\%) due to the uncertainties in the ionisation correction factors. This is in contrast with the unrealistically small errors quoted for line temperatures other than T$_e$([O{\sc iii}]) in the literature, in part due to the commonly assumed methodology of deriving them from the measured T$_e$([O{\sc iii}]) through a theoretical relation. These relations are found from photoionization model sequences and no uncertainty is attached to them although large scatter is found when observed values are plotted; usually the line temperatures obtained in this way carry only the observational error found for the T$_e$([O{\sc iii}]) measurement and does not include the observed scatter, thus heavily understimating the errors in the derived temperature. In fact, no clear relation is found between T$_e$([O{\sc iii}]) and T$_e$([O{\sc ii}]) for the existing sample of objects confirming our previous results. A comparison between measured and model derived T$_e$([O{\sc ii}]) shows than, in general, model predictions overestimate this temperature and hence underestimate the O$^+$/H$^+$ ratio. This, though not very important for high excitation objects, could be of some concern for lower excitation ones for which total O/H abundances could be underestimated by up to 0.2 dex. It is worth noting that the objects observed with double-arm spectrographs, therefore implying simultaneous and spatially coincident observations over the whole spectral range, show less scatter in the T$_e$([O{\sc iii}])\,-\,T$_e$([O{\sc ii}] plane clustering around the N$_e$\,=\,100 cm$^{-3}$ photo-ionisation model sequence. On the other hand, this small scatter could partially be due to the small range of temperatures shown by these objects due to possible selection effects. This small temperature range does not allow either to investigate the metallicity effects found in the ralations between the various line temperatures in recent photo-ionisation models by Izotov et al.\ (2006). Also, the observed objects, as well as those in Paper I, though showing Ne/O and Ar/O relative abundances typical of those found for a large \HII\ galaxy sample (P\'erez-Montero et al.\ 2007), show higher than typical N/O abundance ratios that would be even higher if the [O{\sc ii}] temperatures would be found from photo-ionisation models. We therefore conclude that approach of deriving the O$^+$ temperature from the O$^{2+}$ one should be discouraged if an accurate abundance derivation is sought. These issues could be addressed by re-observing the objects in Table \ref{temp} , which cover an ample range in temperatures and metal content, with double arm spectrographs. This sample should be further extended to obtain a self consistent sample of about 50 objects with high S/N and excellent spectrophotometry covering simultaneously from 3600 to 9900\,\AA\ This simple and easily feasible project would provide important scientific return in the form of critical tests of photoionisation models. The O$^{+} $/O$^{2+} $ and S$^{+} $/S$ ^{2+} $ ratios for all the observed galaxies, except one, cluster around a value of the ``softness parameter" $\eta$ of 1 implying high values of the stellar ionising temperature. For the discrepant object, showing a much lower value of $\eta$, the intensity of the [O{\sc ii}]\,$\lambda\lambda$\,7319,25\,\AA\ lines are affected by atmospheric absorption lines. When the observational counterpart of the ionic ratios is used, this object shows a ionisation structure similar to the rest of the observed ones. This simple exercise shows the potential of checking for consistency in both the $\eta$ and $\eta$' plots in order to test if a given assumed ionisation structure is adequate. In fact, these consistency checks show that the stellar ionising temperatures found for the observed \HII\ galaxies using the ionisation structure predicted by state of the art ionisation models result too low when compared to those implied by the corresponding observed emission line ratios. Therefore, metallicity calibrations based on abundances derived according to this conventional method are probably bound to provide metallicities which are systematically too high and should be revised. | 7 | 10 | 0710.1828 |
0710 | 0710.4956_arXiv.txt | We present Submillimeter Array observations of the z=3.91 gravitationally lensed broad absorption line quasar \apm\, which spatially resolve the 1.0~mm (200~$\mu$m rest-frame) dust continuum emission. At 0\fasec4 resolution, the emission is separated into two components, a stronger, extended one to the northeast ($46\pm5$~mJy) and a weaker, compact one to the southwest ($15\pm2$~mJy). We have carried out simulations of the gravitational lensing effect responsible for the two submm components in order to constrain the intrinsic size of the submm continuum emission. Using an elliptical lens potential, the best fit lensing model yields an intrinsic (projected) diameter of $\sim$80~pc, which is not as compact as the optical/near-infrared (NIR) emission and agrees with previous size estimates of the gas and dust emission in \apm. Based on our estimate, we favor a scenario in which the 200~$\mu$m (rest-frame) emission originates from a warm dust component (T$_{\rm d}$=150-220~K) that is mainly heated by the AGN rather than by a starburst (SB). The flux is boosted by a factor of $\sim$90 in our model, consistent with recent estimates for \apm. | \apm\, (=APM08) is a strongly lensed broad absorption line (BAL) quasar \citep{irwin98,lewis98} with a very powerful active galactic nucleus \citep[AGN;][]{soifer87}. This combination makes it an extremely bright object despite its redshift of $z$=3.91 \citep{downes99}. Its bolometric luminosity is thought to be $\sim$5$\times$10$^{13}$~L$_{\odot}$, amplified by up to a factor of 100 by a foreground galaxy which has yet to be identified. Due to the large amplification, broad absorption lines have been detected even in the X-ray \citep{chartas02}. Ground-based, {\it Chandra} and HST observations have suggested that \sapm\, consists of at least two components separated by $\sim$$0\farcs4$. \citet[=I99]{ibata99} and \citet[=E00]{egami00} later detected a third image (their image C), and argued that it is likely a third lensed image of the background QSO instead of the lensing galaxy. \cite{lewis02b} obtained optical spectra of the three images using HST/STIS, and showed that the three spectra are quite similar. This indicates that the image C is indeed a third lensed image of \sapm. Due to the magnification by the lens, several molecular lines have been detected in \sapm\, \citep[e.g.,][] {weiss07,guelin07,riecher06,wagg06,garcia06,wagg05,lewis02a,downes99}, suggesting a molecular gas mass of M$_{\rm gas}$(H$_2$)=8$\times$10$^{10}$~$m^{-1}$~M$_\odot$ \citep[][$m$$\equiv$magnification factor]{riecher06}. The dust mass is estimated to be M$_{\rm dust}$=5$\times$10$^8$~$m^{-1}$~M$_\odot$ \citep[][W07]{downes99,weiss07}. However, despite the increasing amount of mm-data, no tight constraints have been set on the size of the dust and/or gas emission which may help to characterize the heating source of the dust and gas, i.e., whether it is in the form of an AGN and/or starburst (SB). Only recently have W07 presented some indirect evidence through brightness temperature arguments that the molecular line and dust emission originate from a compact region with a radius of 100-200~pc. In this {\it Letter}, we present observations from the Submillimeter Array (SMA) with sufficient resolution (0\fasec4) to resolve the 1.0~mm dust continuum emission in \sapm, and lensing models that constrain its intrinsic source size. | \label{con} We have detected the 200~$\mu$m rest-frame continuum emission in \sapm\, using the SMA with an angular resolution of $\sim$0\fasec4. The two (main) lensed images are clearly separated with a combined flux of $\sim$60$\pm$12~mJy. Simulations of the gravitational lensing effect in this system yield a diameter of the intrinsic (submm) continuum emitting region of $\sim$80~pc and magnification factor of 90. Our data and simulations seem to be only consistent with a lensing scenario including a third lensed image of \sapm\, close to NE if the position of the compact optical/NIR emission and the position of the extended submm emission are offset from each other before lensing by $\sim$0\fasec003 ($\equiv$21 pc). Further (sub)mm observations may be beneficial to determine whether such a positional offset is indeed necessary or more complicated lens potentials have to be considered. Given our size estimate, we favor a scenario in which the 200~$\mu$m emission originates from a warm dust component that is supposed to be mainly heated by the AGN rather than by a SB. | 7 | 10 | 0710.4956 |
0710 | 0710.1097_arXiv.txt | We report extensive photometry of the dwarf nova V419 Lyr throughout its 2006 July superoutburst till quiescence. The superoutburst with amplitude of $\sim 3.5$ magnitude lasted at least 15 days and was characterized by the presence of clear superhumps with a mean period of $P_{\rm sh}=0.089985(58)$ days ($129.58\pm0.08$ min). According to the Stolz-Schoembs relation, this indicates that the orbital period of the binary should be around 0.086 days i.e. within the period gap. During the superoutburst the superhump period was decreasing with the rate of $\dot{P}/P_{sh}=-24.8(2.2)\times10^{-5}$, which is one of the highest values ever observed in SU UMa systems. At the end of the plateau phase, the superhump period stabilized at a value of 0.08983(8) days. The superhump amplitude decreased from 0.3 mag at the beginning of the superoutburst to 0.1 mag at its end. In the case of V419 Lyr we have not observed clear secondary humps, which seems to be typical for long period systems. \noindent {\bf Key words:} Stars: individual: V419 Lyr -- binaries: close -- novae, cataclysmic variables | Dwarf novae -- a subclass of Cataclysmic Variable stars -- are quite well studied interacting binary systems composed of late-type red dwarf secondary and white dwarf primary stars (Warner 1995, Hellier 2001). Matter transferred from the red dwarf forms an accretion disc around the white dwarf. Although in the last decade significant progress has been made in explaining the behaviour of dwarf novae light curves, some physical processes ongoing in these systems are still not fully understood (see for example Smak 2000, Schreiber and Lasota 2007). In particular, the thermal-tidal instability model of Osaki (1996, 2005) describing the phenomenon of superoutbursts and superhumps may be tested by examination of SU UMa-type dwarf novae light curves. Additionally, objects near and inside the so called period gap are very important from an evolutionary point of view. Those systems give us an unprecedented opportunity to study the evolution of dwarf novae. V419 Lyr is a poorly studied cataclysmic variable discovered by Kurochkin (1990) and originally classified as a Z Cam-type dwarf nova. Later, Nogami et al. (1998) caught this object in outburst and found superhumps in its light curve. Detection of superhumps together with characteristic properties of the outburst allowed them to classify V419 Lyr as a SU UMa-type dwarf nova, but short coverage of the eruption did not allow accurate determination of the superhump period. Nevertheless, there was a strong suggestion that V419 Lyr has one of the longest orbital periods known among SU UMa variables. This object has been monitored at various photometric bands by the Variable Star Network (VSNET) (see for example Kato et al. 2004a). The observations from that program enabled a tentative determination of the supercycle period to be about $\sim 340$ days (Katysheva and Pavlenko 2003). Moreover, Morales-Rueda and Marsh (2002) obtained a spectrum of V419 Lyr during outburst showing a relatively broad absorption feature around 430-440 nm. In this work we present an analysis of photometric data collected during the 2006 July superoutburst of V419 Lyr. The data are much richer than previous studies and provide us with an opportunity to determine parameters describing this system more precisely. | The orbital period of V419 Lyr is unknown. However it is possible to estimate its value using the relation in Stolz and Schoembs (1984) connecting the period excess $\epsilon$ defined as $P_{\rm sh}/P_{\rm orb}-1$ with the orbital period of the binary. This empirical relation is as follows: \begin{equation} \epsilon = 0.858(11) \cdot P_{\rm orb} - 0.0282(2) \end{equation} Using the definition of $\epsilon$ and knowing $P_{\rm sh}$ for V419 Lyr, we were able to estimate the orbital period as $P_{\rm orb}\approx 0.086$ days. This is slightly longer than two hours which indicates that V419 Lyr is a dwarf nova in the period gap. Many characteristics of V419 Lyr are typical of SU UMa stars. It goes into superoutburst every year or so, the eruption lasts about two weeks and has an amplitude of $\sim 3.5$ mag. Superhumps appear shortly after the beginning of the superoutburst and have a maximum amplitude of 0.3 mag, which decreases to 0.1 mag at the end of the outburst. In addition to its long orbital period, V419 Lyr is unusual in two other properties. Its superhump period derivative has one of the largest negative values known and it shows only a weak trace of secondary humps in the final stages of the superoutburst. \bigskip \noindent {\bf Acknowledgments.} ~We acknowledge generous allocation of the Warsaw Observatory 0.6-m telescope time. This work used the online service of the VSNET and AAVSO. We would like to thank Prof. J\'ozef Smak for fruitful discussions. | 7 | 10 | 0710.1097 |
0710 | 0710.1574_arXiv.txt | {We present the results of a {\it Chandra} soft X-ray observation of the spectacular ionization cone in the nearby Seyfert~2 galaxy NGC~5252. As almost invariably observed in obscured AGN, the soft X-ray emission exhibits a remarkable morphological concidence with the cone ionized gas as traced by HST O[{\sc iii}] images. Energy-resolved images and high-resolution spectroscopy suggest that the X-ray emitting gas is photoionized by the AGN, at least on scales as large as the innermost gas and stellar ring ($\le$3~kpc). Assuming that the whole cone is photoionized by the AGN, we reconstruct the history of the active nucles in the last $\sim$10$^5$~years.} | 7 | 10 | 0710.1574 |
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0710 | 0710.0602_arXiv.txt | We describe the discovery of HAT-P-4b, a low-density extrasolar planet transiting BD+36~2593, a $V=11.2$\,mag slightly evolved metal-rich late F star. The planet's orbital period is $3.056536\pm0.000057$\,d with a mid-transit epoch of $2,454,245.8154\pm0.0003$ (HJD). Based on high-precision photometric and spectroscopic data, and by using transit light curve modeling, spectrum analysis and evolutionary models, we derive the following planet parameters: \mpl$=0.68\pm0.04$\,\mjup, \rpl$=1.27\pm0.05$\,\rjup, \rhopl$=0.41\pm0.06$\,\gcmc\ and $a=0.0446\pm0.0012$\,AU\@. Because of its relatively large radius, together with its assumed high metallicity of that of its parent star, this planet adds to the theoretical challenges to explain inflated extrasolar planets. | \label{sec:intro} In the course of our ongoing wide field planetary transit search program HATNet \citep{bakos04}, we have discovered a large radius and low density planet orbiting an $11$th magnitude star BD+36~2593. This planet is the fifth member of a group of low-density transiting exoplanets. The combination of its low mass and the relatively high metallicity and age of the parent star makes theoretical interpretation of its large radius difficult. In this Letter we describe the observational properties of the system and derive the physical parameters both for the host star and for the planet. We also briefly comment on the theoretical status of inflated extrasolar planets. | \label{sec:disc} We presented the discovery data and derived the physical parameters of HAT-P-4b, an inflated planet orbiting BD+36~2593. Among the 20 transiting planets discovered so far, there are five with \rhopl~$\lesssim0.4$\,\gcmc. All others have at least $50$\% higher densities. For ease of comparison, \tabr{puffy} lists the relevant properties of the five inflated planets. It is remarkable how similar these planets are (except for TrES-4 that has distinctively low density). With its Safronov number of $0.036$, HAT-P-4b belongs to the Class II planets according to the recent classification of \cite{hansen07} and (together with TrES-4) further strengthens the mysterious dichotomy of the known transiting planets in this parameter. The parent star of HAT-P-4b is among the largest radii, largest mass, lowest gravity and highest metallicity transiting planet host stars. \notetoeditor{This is the intended place of \tabr{puffy}} \begin{deluxetable}{lcccccc} \tabletypesize{\scriptsize} \tablecaption{ Comparison of the properties of inflated planets.\tablenotemark{a} \label{tab:puffy}} \tablewidth{0pt} \tablehead{ \colhead{Name} & \colhead{P} & \colhead{a} & \colhead{M} & \colhead{R} & \colhead{$\rho$} & \colhead{$\log g$} \\ \colhead{} & \colhead{(d)} & \colhead{(AU)} & \colhead{(M$_J$)} & \colhead{(R$_J$)} & \colhead{(\sc{cgs})} & \colhead{(\sc{cgs})} } \startdata WASP-1b & 2.52 & 0.038 & 0.87 & 1.40 & 0.39 & 3.04 \\ HAT-P-4b & 3.06 & 0.045 & 0.68 & 1.27 & 0.41 & 3.02 \\ HD~209458b & 3.53 & 0.045 & 0.64 & 1.32 & 0.35 & 2.96 \\ TrES-4 & 3.55 & 0.049 & 0.84 & 1.67 & 0.22 & 2.87 \\ HAT-P-1b & 4.47 & 0.055 & 0.53 & 1.20 & 0.38 & 2.96 \\ \enddata \tablenotetext{a} {Data from \cite{shporer07}, this paper, \cite{burro07}, \cite{mandu07} and \cite{winn07}. From top to bottom, metallicities for the parent stars are: $0.23$ \citep{stempels07}, $0.24$, $0.02$ $0.0$ (adopted) and $0.13$.} \end{deluxetable} Current models of irradiated giant planets are able to match the observed radii of most of the planets without invoking any additional heating mechanism. Higher metallicity cases, such as the present one, however, may pose problems (assuming that the planet and star have similar metallicities). More metals imply two opposite effects on the radius: (i) inflating it due to higher opacities in the envelope; (ii) shrinking it due to the higher molecular weight of the interior and the possible development of a large high density core. These effects have been discussed recently by \cite{burro07}. Since WASP-1b is similar in several aspects (i.e., irradiance, metallicity) to HAT-P-4b, we consider the coreless models of WASP-1b as shown in Fig.~7 of \cite{burro07}. It seems that HAT-P-4b can be fitted by near solar metallicity coreless models, assuming that its age is not too much greater that $4$\,Gyr. We also refer to the layered convective mechanism of \cite{chabrier07} that gives an alternative explanation for planets with inflated radii. We conclude that more definite statements on the relation of the observations and planet structure theories can be made only by reaching higher accuracy in the observed star/planet parameters. Nevertheless, HAT-P-4b (together with WASP-1b) does not seem to support the existence of a simple relation between host star metallicity and planet's core mass \citep[see][]{guillot06,burro07}. | 7 | 10 | 0710.0602 |
0710 | 0710.2741_arXiv.txt | We study the cosmological evolution of an induced gravity model with a self-interacting scalar field $\sigma$ and in the presence of matter and radiation. Such model leads to Einstein Gravity plus a cosmological constant as a stable attractor among homogeneous cosmologies and is therefore a viable dark-energy (DE) model for a wide range of scalar field initial conditions and values for its positive $\gamma$ coupling to the Ricci curvature $\gamma \sigma^{2}R$. | Several years ago a model for a varying gravitational coupling was introduced by Brans and Dicke \cite{BD}. The model consisted of a massless scalar field whose inverse was associated with the gravitational coupling. Such a field evolved dynamically in the presence of matter and led to cosmological predictions differing from Einstein Gravity (EG) in that one generally obtained a power-law dependence on time for the gravitational coupling. In order to reduce such a strong time dependence in a cosmological setting, while retaining the Brans-Dicke results in the weak field limit, several years ago a simple model for induced gravity \cite{CV,TV} involving a scalar field $\sigma$ and a quartic $\lambda\sigma^{4}/4$ potential was introduced. This model was globally scale invariant (that is did not include any dimensional parameter) and the spontaneous breaking of scale invariance in such a context led to both the gravitational constant and inflation, through a non-zero cosmological constant \cite{Zee}. The cosmological consequences of introducing matter as a perturbation were studied leading to a time dependence for the scalar field and consistent results. Since the present observational status is compatible with an accelerating universe dominated by dark energy (DE) we feel that the model should be studied in more detail. \\ In an EG framework quadratic or quartic potentials for canonical scalar fields are consistent with DE only with an extremely fine tuning in the initial conditions leading to slow-roll evolution until the present time. Indeed, a massive or a self-interacting scalar field in EG, respectively behave as dust or radiation during the coherent oscillation regime. In contrast with this, induced gravity with a self-interacting $\lambda\sigma^4/4$ potential has as attractor EG plus a cosmological constant on breaking scale invariance. The simple model we consider illustrates how non-trivial and non perturbative dynamics can be obtained in the context of induced gravity DE, and more generally within scalar-tensor DE. | We have shown how a simple model of self-interacting induced gravity \cite{CV,TV} is a viable DE model, for tiny values of the self-coupling ($\lambda \sim {\cal O}(10^{-128})$). The model has a stable attractor towards EG plus $\Lambda$ and can be very similar to the $\Lambda$CDM model for the homogeneous mode, on taking into account the Solar system constraints quoted in \cite{Bertotti:2003rm,Eubanks,gdotref}. At the cosmological level, in the presence of radiation and dust, it is interesting that for such a simple potential (quartic for induced gravity or interacting for the equivalent Brans-Dicke model) the model has an attractor towards EG plus $\Lambda$, very differently from the case of a massless scalar - i.e. $\lambda = 0$ -, for which there is no mechanism of attraction towards EG. The attractor mechanism towards GR and the onset to acceleration are both inevitably triggered by the same mechanism, i.e. scale symmetry breaking in this induced gravity model. \begin{figure}[t!!] \centering \epsfig{file=ws_phi4.eps, width=8.5 cm} \caption{Evolution of $w_{\rm DE}$ as defined in (\ref{rhopfake}) for different choices of $\gamma$.\label{ws}}. \end{figure} In contrast with EG, the choice of a runaway potential \cite{PR} for quintessence is not mandatory (see however \cite{others} for a study of these potential in the context of scalar-tensor theories). If the final attractor is an accelerated universe, the only constraints on parameters and initial conditions come from observations. We have shown that late time cosmology - after recombination, for instance - is mostly insensitive with respect to the initial time derivative of $\sigma$. For this reason, the full set of parameters and initial conditions of the model are fully specified on taking the observed values for $G, H_0, \Omega_\Lambda$.\\ We have discussed in detail such model in the context of Einstein gravity, i.e. keeping the Newton constant (approximately) fixed at its actual value. The model predicts an equation of state of the equivalent Einstein model with $w_{DE}$ slightly less than $-1$ at present and homogeneous cosmology by itself seems able to constrain $\gamma \lesssim {\cal O}(10^{-3})$, although a full statistical analysis is clearly beyond the scope of this work.\\ It is interesting to also study also what other potentials, besides the simple potential $\lambda \sigma^4/4$ employed in this paper, will also be compatible with the observed evolution of the universe.\\ \\ {\bf Acknowledgment.} We wish to thank the Referee for helpful and constructive criticism. | 7 | 10 | 0710.2741 |
0710 | 0710.0058_arXiv.txt | The Wilkinson Microwave Anisotropy Probe (WMAP) three year results are used to constraint non-minimal inflation models. Two different non-minimally coupled scalar field potentials are considered to calculate corresponding slow-roll parameters of non-minimal inflation. The results of numerical analysis of parameter space are compared with WMAP3 data to find appropriate new constraints on the values of the non-minimal coupling. A detailed comparison of our results with previous studies reveals the present status of the non-minimal inflation model after WMAP3.\\ {\bf PACS}:\, 04.50.+h,\, 98.80.-k\\ {\bf Key Words}: Scalar-Tensor Gravity,\, Inflation,\, WMAP3 Data | Non-minimal coupling (NMC) of scalar field with gravity is necessary in many situations of physical and cosmological interest. There are several compelling reasons to include an explicit non-minimal coupling in the action ( see for example [1,2,3,4] and references therein ). NMC arises at the quantum level when quantum corrections to the scalar field theory are considered. It is necessary also for the renormalizability of the scalar field theory in curved space. In most theories used to describe inflationary scenarios, it turns out that a non-vanishing value of the coupling constant is unavoidable [2]. In general relativity, and in all other metric theories of gravity ( in which the scalar field is not part of the gravitational sector ), the coupling constant necessarily assumes a non-vanishing value[5]. The study of the asymptotically free theories in an external gravitational field with a Gauss-Bonnet term shows a scale dependent coupling parameter. For instance, asymptotically free grand unified theories have a non-minimal coupling depending on a renormalization group parameter that converges to the value of $\frac{1}{6}$ or to any other initial conditions depending on the gauge group and on the matter content of the theory[6]. In view of the above results and several other motivations( see for example [7,8] and references therein), it is then natural to incorporate an explicit NMC between scalar field and Ricci scalar in the inflationary paradigm. Generally, with non-minimally coupled scalar field it is harder to achieve accelerated expansion of the universe[2,7]. Part of this difficulty is related to the more sophisticated machinery of fine tuning. Nevertheless, over the last decade several non-minimal inflation scenario have been proposed to find reliable framework for issues such as graceful exit of inflationary phase and the observational constraints on the values that non-minimal coupling can attain in order to have successful inflationary scenario are discussed [8-16]. The recent astronomical observations, specially high precision data of WMAP3 [17] have opened new doors in the field of observational cosmology. As a result, these data have significant impact on the inflation paradigm. In this regard, these high precision data will give more accurate bounds on the values of non-minimal coupling in a typical non-minimal inflation model. The purpose of this letter is to study impact of WMAP3 and non-minimal inflation. Considering some well-known inflationary potentials, we explore new observational constraints on the values of non-minimal coupling to have successful non-minimal inflation. By definition of an effective scalar field potential, we show that there is a region in parameter space that inflation is driven by the non-minimal coupling term. A detailed study of spectral index and its running shows the spontaneous exit of inflationary phase ( without any mechanism ) in a suitable region of the parameter space. We also compare our results with the results of previous studies. This comparison reveals the present status of non-minimal inflation paradigm after WMAP3. | 7 | 10 | 0710.0058 |
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0710 | 0710.3167_arXiv.txt | In this brief proceedings article I summarize the review talk I gave at the IAU 246 meeting in Capri, Italy, glossing over the well-known results from the literature, but paying particular attention to new, previously unpublished material. This new material includes a careful comparison of the apparently contradictory results of two independent methods used to simulate the evolution of binary populations in dense stellar systems (the direct $N$-body method of \cite{2007ApJ...665..707H} and the approximate Monte Carlo method of \cite{2005MNRAS.358..572I}), that shows that the two methods may not actually yield contradictory results, and suggests future work to more directly compare the two methods. | Globular clusters are observed to contain significant numbers of binary star systems---so many, in fact, that they must have born with binaries (\cite{1992PASP..104..981H}). Their presence in clusters is important for two complementary reasons. Through super-elastic dynamical scattering interactions, they act as an energy source which may postpone core collapse, and may be the dominant factor in setting the core radii of observed Galactic globulars. Similarly, the dense stellar environment and increased dynamical interaction rate in cluster cores is responsible for the high specific frequency of stellar ``exotica'' found in clusters, including low-mass X-ray binaries (LMXBs), cataclysmic variables (CVs), blue straggler stars (BSSs), and recycled millisecond pulsars (MSPs). | \label{sec:summary} In this proceedings article I have very briefly discussed the connection between binary stars and globular cluster dynamics, moving from basic physics to current research in the span of a few paragraphs. A thorough, easily readable, and fairly recent discussion of the material can be found in \cite{2003gmbp.book.....H}. The primary new material presented here is a comparison in phase space of the seemingly contradictory binary population evolution simulations of \cite{2005MNRAS.358..572I} and \cite{2007ApJ...665..707H}, showing that they may in fact both represent the same underlying physics. In other words, new simulations must be performed to better compare the two very different methods. | 7 | 10 | 0710.3167 |
0710 | 0710.3351_arXiv.txt | We present the results of \spi\ mid-infrared spectroscopic observations of two highly-obscured massive X-ray binaries: \igrj\ and \gx. Our observations reveal for the first time the extremely rich mid-infrared environments of this type of source, including multiple continuum emission components (a hot component with $T$ $>$ 700~K and a warm component with $T$ $\sim$ 180~K) with apparent silicate absorption features, numerous \ion{H}{1} recombination lines, many forbidden ionic lines of low ionization potentials, and pure rotational \hh\ lines. This indicates that both sources have hot and warm circumstellar dust, ionized stellar winds, extended low-density ionized regions, and photo-dissociated regions. It appears difficult to attribute the total optical extinction of both sources to the hot and warm dust components, which suggests that there could be an otherwise observable colder dust component responsible for the most of the optical extinction and silicate absorption features. The observed mid-infrared spectra are similar to those from Luminous Blue Variables, indicating that the highly-obscured massive X-ray binaries may represent a previously unknown evolutionary phase of X-ray binaries with early-type optical companions. Our results highlight the importance and utility of mid-infrared spectroscopy to investigate highly-obscured X-ray binaries. | \label{sec_intro} Recently, a large number of highly-obscured (e.g., \nh\ $\ge$ 10$^{23}$ cm$^{-2}$) massive X-ray binaries have been discovered by the \integ\ hard X-ray ($\ge$ 15 keV) satellite \citep{wet03}. The prototypical case is \igrj\ which shows variable high obscuration in the X-ray, sometimes reaching \nh\ $\simeq$ 2 $\times$ 10$^{24}$ \citep{coet03, walteret03}. Its bright optical and near-infrared (IR) counterpart is an early B-type supergiant star with numerous emission lines \citep{fc04}. Interestingly the obscuration toward \igrj\ obtained in the optical and near-IR wavebands (\av\ $\sim$ 18) is almost two orders of magnitude smaller than that inferred from the X-rays, which suggests that the extreme obscuration seen in the X-ray is intrinsic only to the X-ray source. However, \av\ $\sim$ 18 is still greater than the interstellar obscuration, indicating the existence of substantial circumstellar material around the supergiant companion. In order to fully understand the implications of the \integ\ discoveries --- specifically if they imply existence of a separate class of highly-obscured X-ray binaries --- we must investigate any similarities between the new \integ\ sources and previously known sources. As already suggested by \citet{ret03}, the X-ray pulsar \gx\ appears similar to the new \integ\ highly-obscured massive X-ray binaries: it has variable X-ray obscuration of \nh\ $\simeq$ 10$^{23}$--10$^{24}$ cm$^{-2}$, the compact source is a neutron star, and the optical companion is an early B-type supergiant \citep[or hypergiant;][]{ket95}, like \igrj. Recently, \citet[][; hereafter Paper~I]{kmr06}, using optical, near-, and mid-IR ($\le$~20~$\mu$m) spectral energy distributions (SEDs), have found that both sources have strong mid-IR excesses that they identify as continuum emission from hot dust. This is suggestive that they have very similar circumstellar material which may be related to their strong X-ray obscuration. In this {\em Letter}, we present the results of \spi\ mid-IR spectroscopic observations of \igrj\ and \gx, showing that both sources indeed have very similar rich mid-IR properties previously unknown for X-ray binaries. | \label{sec_dis_sum} Our \spi\ spectroscopic observations of \igrj\ and \gx\ have revealed for the first time the rich mid-IR environment of highly-obscured X-ray binaries. This includes two dust components with prominent silicate absorption, numerous \ion{H}{1} recombination lines, many forbidden ionic lines, and pure rotational \hh\ lines. Based on the observed spectra, we infer the following components for \igrj\ and \gx: (1) hot ($T$ $>$ 700~K) and warm ($T$ $\sim$ 180~K) circumstellar dust; (2) ionized stellar winds responsible for the \ion{H}{1} lines; (3) extended low-density ionized regions for the forbidden lines; and (4) photo-dissociated regions associated with the PAH, \hh\ and possibly the \silii\ line emission. For the forbidden lines, all the detected lines have relatively low ionization potentials like in starburst galaxies where the radiation is relatively soft (compared with active galactic nuclei, for instance). This may indicate that the illumination of hard X-rays from the central compact X-ray source is not primarily responsible for the forbidden line emission. However, the inferred temperature for the radiation exciting the \neii\ and \neiii\ lines is hotter than the stellar photospheres, so there can be some contribution from the compact object. Considering that \niii\ and \feii\ were detected only in \igrj\ while \feiii\ was only in \gx, the radiation field of \igrj\ may be softer than \gx, as demonstrated by the small temperature difference between the two sources (see \S~\ref{sec_highres}). Perhaps the most natural explanation for the origin of the hot dust component relies on dust formation in the dense outflows from the early-type companions of \igrj\ and \gx, as is the case in most sgB[e] stars (i.e., B-type supergiants with forbidden emission lines). However, the origin of the warm circumstellar dust component is very uncertain, and B[e] stars seldom show evidence for this type of warm dust. If both the hot and warm dust components have spherical shell geometres around the central star, then the associated optical extinctions are: \av\ $\simeq$ 4 $\times$ 10$^{-4}$ $Q_{\rm abs}$ ($M_{\rm d}$/10$^{-6}$ \msun) ($T_{\rm d}$/100 K)$^4$ ($L_{\rm UV}$/10$^{39}$ ergs s$^{-1}$), where $M_{\rm d}$ and $T_{\rm d}$ are the mass and temperature of the dust components, $Q_{\rm abs}$ $<$ 1 is the dust absorption coefficient, and $L_{\rm UV}$ is the ultra-violet luminosity of the central star. This gives the optical extinctions \av\ $\ll$ 1 for both the hot and warm dust components of \igrj\ and \gx. (Here we use 1 $\times$ 10$^{39}$ ergs s$^{-1}$ for $L_{\rm UV}$ for both sources.) Therefore, under the assumption of spherical shell geometry, both the hot and warm dust components of \igrj\ and \gx\ contribute very little to the total optical extinction, suggesting that the hot and warm dust components are {\em NOT} strongly associated with the silicate absorption features. This also applies to the dust associated with the warm extended \hh\ gas since the optical extinction from this component is tiny (i.e., Av $\ll$ 1; see \S~\ref{sec_highres}). What's then the origin of the silicate absorption features? If it is due to the foreground ISM, we would expect to see the $\rm CO_2$ ice feature around 15 $\mu$m, especially for \igrj, based on the intensity ratio between the silicate absorption feature and ice feature found in the ISM \citep{ket05}. The absence of the ice feature in our spectra supports the interpretation that the silicate absorption features are probably not associated with the foreground ISM. One possibility may be the existence of an undisclosed colder (e.g., $\ll$ 100~K) circumstellar dust component which is responsible for the silicate absorption features and most of the optical extinction. We need further longer wavelength observatons to confirm this possibility. Considering that the 9.7~$\mu$m silicate absorption feature represents the oxygen-rich material, the existence of the 9.7 micron silicate absorption feature may indicate that the origin of the potential colder dust component is related to the nucleosynthesis of the progenitors of \igrj\ and \gx\ (or their companions). The optical/near-IR companion of \igrj\ is a sgB[e] star. Such stars are known to have hot ($T$ $\sim$ 1000~K) circumstellar dust, and probably are evolving into Luminous Blue Variables (LBVs) or Wolf-Rayet stars. Our mid-IR spectra of \igrj\ and \gx\ are very similar to that of the LBV P~Cygni which shows many \ion{H}{1} lines and forbidden ionic lines \citep{let96b}. The difference is that while the mid-IR continuum of P~Cygni is due to the free-free emission in stellar winds, the main mid-IR emission of \igrj\ and \gx\ is from multiple dust continua. (The SEDs of both sources observed here are inconsistent with the free-free emission, as mentioned in Paper~I.) However, we note that some LBV stars have also been observed to have thermal mid-IR dust emission \citep{let96a}. Based on the fact that B[e] stars seldom show mid-IR forbidden line emission, the B-type supergiant (or hypergiant) companions of \igrj\ and \gx\ may be in the evolutionary track to LBVs, implying that the extremely high obscuration seen in some massive X-ray binaries may be a phenomenon associated with the evolutionary phase. This scenario is also consistent with the fact that the highly-obscured massive X-ray binary Cygnus~X-3 has a Wolf-Rayet star companion, together with the circumstellar dust emission of $T$ $\sim$ 250~K \citep{ket02}. | 7 | 10 | 0710.3351 |
0710 | 0710.4601_arXiv.txt | We constructed some main-sequence mergers from case A binary evolution and studied their characteristics via Eggleton's stellar evolution code. Both total mass and orbital angular momentum are conservative in our binary evolutions. Assuming that the matter from the secondary homogeneously mixes with the envelope of the primary and that no mass are lost from the system during the merger process, we found that some mergers might be on the left of the zero-age main sequence as defined by normal surface composition (i.e helium content $Y=0.28$ with metallicity $Z=0.02$ for Pop I) on a colour-magnitude diagram(CMD) because of enhanced surface helium content. The study also shows that central hydrogen content of the mergers is independent of mass. Our simple models provide a possible way to explain a few blue stragglers (BSs) observed on the left of zero-age main sequence in some clusters, but the concentration toward the blue side of the main sequence with decreasing mass predicted by Sandquist et al. will not appear in our models. The products with little central hydrogen in our models are probably subgiants when they are formed, since the primaries in the progenitors also have little central hydrogen and will likely leave the main sequence during merger process. As a consequence, we fit the formula of magnitude $M_{\rm v}$ and $B-V$ of the mergers when they return back to thermal equilibrium with maximum error 0.29 and 0.037, respectively. Employing the consequences above, we performed Monte Carlo simulations to examine our models in an old open cluster NGC 2682 and an intermediate-age cluster NGC 2660. Angular momentum loss (AML) of low mass binaries is very important in NGC 2682 and its effect was estimated in a simple way. In NGC 2682, binary mergers from our models cover the region with high luminosity and those from AML are located in the region with low luminosity, existing a certain width. The BSs from AML are much more than those from our models, indicating that AML of low mass binaries makes a major contribution to BSs in this old cluster. Our models are corresponding for several BSs in NGC 2660. At the region with the most opportunity on the CMD, however, no BSs have been observed at present. {\bf Our results are well-matched to the observations if there is $\sim 0.5M_\odot$ of mass loss in the merger process, but a physical mechanism for this much mass loss is a problem.} | Much evidence shows that primordial binaries make an important contribution to blue stragglers (BSs) \cite{fer03,dpa04,map04}. At present, a few BSs, i.e. F190, $\theta$ Car, have already been confirmed to be in binaries by observations, and their formation may be interpreted by mass transfer between the components of a binary. Whereas in intermediate-age and old open and globular clusters, the number of observed close binaries among well-studied BSs is consistent with the hypothesis of binary coalescence. For example, Mateo et al. \shortcite{mat90} made a comparison of the number of close binaries with the total number of BSs in NGC 5466 and found that it is an acceptable claim that all non-eclipsing BSs are formed as the result of mergers of the components in close binaries, though the possibility of other mechanisms to produce BSs cannot be ruled out due to the large uncertainties in their analysis. Monte-Carlo simulations of binary stellar evolution \cite{pol94} also show that binary coalescence may be an important channel to form BSs in some clusters (e.g. with an age greater than 40 Myr). Meanwhile, the arguments in theory show that W UMa binaries (low-mass contact binaries) must eventually merge into a single star \cite{web76,web85,ty87,mat90}. Observationally, the lack of radial velocity variations for most BSs further indicates that binary coalescence may be more important than mass transfer for BS formation \cite{str93,pol94}. FK Comae stars are generally considered to be direct evidence for binary coalescence \cite{str93}. The smallest mass ratio of components among observed W UMa systems to date is about 0.06. All of the above show that it is important to study the remnants of close binaries. However the merge process is complicated and the physics during the process is still uncertain. Recently, Andronov, Pinsonneault \& Terndrup \shortcite{apt06} studied the mergers of close primordial binaries by employing the angular momentum loss rate inferred from the spindown of open cluster stars. Their study shows that main sequence mergers can account for the observed number of single BSs in M67 and that such mergers are responsible for at least one third of the BSs in open clusters older than 1 Gyr. The physics of mergers are limiting case treatments in the study of Andronov, Pinsonneault \& Terndrup \shortcite{apt06}. Based on previous studies of contact binaries and some assumptions, we construct a series of merger models in this paper, to study the structure and evolution of the models and show some comparisons with observations. Case A binary evolution has been well studied by Nelson \& Eggleton \shortcite{nel01}. They defined six major subtypes for the evolution (AD, AR, AS, AE, AL and AN) and two rare cases (AG and AB). Three of the subtypes (AD, AR, AS) lead the binary contact as both components are main--sequence stars and two cases (AE and AG) reach contact with one or both components having left the main sequence. As there is no description for weird objects except for two merged main--sequence stars, merger products (except for two main-sequences stars) are generally assumed to have terminated their evolution \cite{pol94}, i.e. they have left the main sequence and cannot be recognized as BSs. Here we are interested in the cases of two main-sequence stars, i.e. cases AS, AR and AD. If $t_{\rm dyn}$, $t_{\rm KH}$, $t_{\rm MS}$ represent the dynamic timescale, thermal timescale and main sequence timescale of the primary (the initial massive star,*1), respectively, the following shows a simple definition of the three evolutionary cases: AD--dynamic Roche lobe overflow (RLOF), $\dot{M}>M/t_{\rm dyn}$; AR--rapid evolution to contact, $\dot{M}>M/t_{\rm KH}, t_{\rm contact} -t_{\rm RLOF}(*1)<0.1t_{\rm MS}(*1)$; AS--slow evolution to contact, $t_{\rm contact} -t_{\rm RLOF}(*1)>0.1t_{\rm MS}(*1)$, where $t_{\rm RLOF}$ and $t_{\rm contact}$ are the ages at which RLOF begins and the binary comes into contact, respectively. In case AD, the core of the secondary spirals in quickly and stays in the center of the merger. The merger then has a chemical composition similar to that of the primary, resembling the result of smoothed particle hydrodynamic calculations \cite{lrs96,sill97,sill01}. We therefore studied just the systems in cases AR and AS for this work. \section {Assumptions} Using the stellar evolution code devised by Eggleton \shortcite{egg71,egg72,egg73}, which has been updated with the latest physics over the last three decades \cite{han94,pol95,pol98}, we re-calculate the models of cases AS and AR with primary masses between 0.89 and $2M_\odot$ until the systems become contact binaries. The structures of the primaries and the compositions of the secondaries are stored to construct the merger remnants. Before the system comes into contact, the accreting matter is assumed to be deposited onto the surface of the secondary with zero falling velocity and distributed homogeneously all over the outer layers. The change of chemical composition on the secondary's surface caused by the accreting matter is \begin{equation} {\partial X_i / \partial t }={(\partial M /\partial t)/[(\partial M /\partial t){\rm d}t+M_{\rm s}}] \cdot (X_{i{\rm a}}-X_{i{\rm s}}), \end {equation} where $\partial M /\partial t$ is the mass accretion rate, $X_{i{\rm a}}$ and $X_{i{\rm s}}$ are element abundances of the accreting matter and of the secondary's surface for species $i$, respectively, and $M_{\rm s}$ is the mass of the outermost layer of the secondary. The value of $M_{\rm s}$ will change with the moving of the non-Lagrangian mesh as well as the chosen model resolution, but it is so small ($\sim 10^{-9}-10^{-12} M_{\odot}$) in comparison with $(\partial M /\partial t){\rm d}t$ ($\sim 10^{-3}-10^{-5} M_{\odot}$) during RLOF that we may ignore the effect of various $M_{\rm s}$ on element abundances. Before and after RLOF, we get $\partial X_{\rm i}/\partial t =0$ from the equation, which is reasonable in the absence of mixing \cite{ch04}. The merger models are constructed based on the following assumptions: (i) contact binaries with two main-sequence components coalesce finally and the changes of structures of individual components during coalescence are ignored; (ii) the matter of the secondary is homogeneously mixed with that of the primary beyond the core-envelope transition point, which separates the core and the envelope of the mass donor; (iii) the system mass is conserved. Firstly, we present a brief discussion on these assumptions. Webbink \shortcite{web76} studied the evolutionary fate of low-mass contact binaries, and found that a system cannot sustain its binary character beyond the limits set by marginal contact evolution ($\mu =M_1/(M_1+M_2)=1.0$). He stated that a contact binary will very likely coalesce as the primary is still on the main sequence in a real system. Up to now, it is widely believed that case AD probably leads to common envelope, spiral-in, and coalescence on quite a short timescale. The final consequences of AS and AR are not very clear, but Eggleton \shortcite{egg00} pointed out that systems undergoing AR or AS evolution may maintain a shallow contact (perhaps intermittently) as the mass ratio becomes more extreme, and finally coalesce. Recent study on W UMa (Li, Zhang \& Han, 2005) also shows that these systems will be eventually coalescence. The merged timescale, i.e. the time from a binary contact to coalescence, is important here. If it is too long, the structures of both components will change remarkably and the system may have not completed coalescence within the cluster age. There are many {\bf conflicting estimates} for the timescale, however, from observations and theoretical models of the merger process. {\bf Early observational estimates range from} $10^7$--$10^8$ yr in various environments \cite{van79,eggen89}. The following study explored the average age about $5 \times 10^8$ yr \cite{van94,dry02}. Bilir et al. \shortcite{bil05} pointed out that the age difference between field contact binaries and chromospherically active binaries, 1.61 Gyr, is likely an upper limit for the contact stage by assuming an equilibrium in the Galaxy, whereas the study of W UMa by Li, Han \& Zhang \shortcite{lhz04} suggested a much longer timescale, about 7 Gyr. We adopt the empirically estimated values in this paper {\bf (i.e from $5 \times 10^7$ to $1 \times 10^9$ yrs)} and ignore the changes of structure of individual components during merger process. For low-mass contact binaries, the common envelope is convective \cite{web77}, and the matter in it is thus homogeneous. If a system mimics shallow contact during coalescence, it is reasonable to assume that the matter of the secondary mixes with the envelope homogeneously. Van't Veer \shortcite{van97} found that the mass loss from the system during coalescence is at a rate of about $2 \times 10^{-10}M_\odot{\rm yr}^{-1}$ by observations. If we consider that the coalescence time is $5 \times 10^8$ yr in a binary, only $0.1M_\odot$ is lost from the system as the binary finally becomes a single star. We then roughly assume that the mass is conservative during coalescence. However mass loss might be an important way to carry orbital angular momentum away from the binary in this process. Secondly, we discuss the choice of the core-envelope transition point which separates the core and the envelope in the primary. Many characteristics of the merger are relevant to the choice, e.g. the chemical composition in the envelope, evolutionary track on Hertzsprung-Russel diagram, and some observational characteristics. Unfortunately, one cannot find the core-envelope transition point in a main-sequence star as easily as in evolved stars because the density profile, as well as many other thermodynamic quantities (entropy, pressure, temperature etc.), is smooth and does not have a deep gradient for main-sequence stars. Chen \& Han \shortcite{ch05} studied the influences of core-envelope transition point on the mergers of contact binaries with two main-sequence components. They found that one may ignore the effects which result from different choices of the transition point on colours and magnitudes of the merger if it is outside the nuclear reaction region of the primary, which is commonly considered as the nearest boundary of the secondary reaches in cases AS and AR. In this paper, the core-envelope transition is determined as the point within which the core produces 99 per cent of total luminosity. This choice is generally outside the nuclear reaction regions and has little effect on the final results. Finally the merger remnant is constructed as follows: it has the total mass of the system and a chemical composition within $M_{\rm 1c}$ similar to the core of the primary. The chemical composition in the envelope of the merger is given by \begin{equation} X_i=(M_{i2}+M_{i1\rm b})/(M_2+m_{\rm b}), \end{equation} where $M_{i2}$ and $M_{i1\rm b}$ denote the total masses of species $i$ of the secondary and of the primary's envelope, respectively. $m_{\rm b}$ is the envelope mass of the primary. There might be a region in which the helium abundance is less than that of the outer region. The matter in this region then has a lower mean molecular weight than that in the outer region. This results in secular instability and thermohaline mixing \cite{kip80,ulr72}. We include it as a diffusion process in our code \cite{ch04}. In the models of Nelson \& Eggleton \shortcite{nel01}, both total mass and angular momentum are conservative. It was mentioned by the authors, however, that these assumptions were only reasonable for a restricted range of intermediate masses, i.e spectra from about G0 to B1 and luminosity class III-V. Observationally, some low mass binaries with late-type components show clear signs of magnetic activity, which indicates that the systems evolve by way of a scenario implying angular momentum loss (AML) by magnetic braking \cite{mes84}. Magnetized stellar winds probably do not carry off much mass, but they are rich in angular momentum because of magnetic linkage to the binaries. For close binaries, rotation is expected to synchronize with orbital period, so AML is at the expense of the orbital angular momentum, resulting in orbital decaying and the components approaching each other. A detached binaries, then, may become contact and finally coalesce at or before the cluster age \cite{ste95}. There are a number of subjects including the treatment of AML \cite{lhz04,ste06,mk06,dek06}. For simplicity, the conservative assumption is also adopted in our binary evolutions. In old clusters, however, AML of low mass binaries is very important and {\bf we estimate its importance in another way (see section 4.2).} | In Sect.4, we notice that the timescale from contact to complete coalescence, $t_{\rm cc}$, strongly affects the initial parameter space of primordial binaries which eventually produce single BSs in a cluster. On the other hand, there are {\bf some conflicting estimates} for $t_{\rm cc}$ from observations and theoretical models. {\bf In this paper we adopt empirical values, i.e. $t_{\rm cc}$ is short in comparison to the evolution timescale of both components in a binary, and ignore the changes during the merger process. In this section, we will first discuss the consequences of a long $t_{\rm cc}$. Because of evolution of both components during the merger process, the primaries have lower central hydrogen content and the matter from the secondaries have larger He content. The former results in a redder colour for the mergers while the latter makes the mergers bluer. So the final positions of the mergers are possibly similar to those shown in this paper, except that the primaries have left the main sequence at final coalescence. This case will appear in the mergers with little central hydrogen. For example, a star with $2M_\odot$ may evolve from ZAMS to exhausted of central hydrogen in $10^9$ yr, and then none of the mergers from binaries with primary' masses larger than $2M_\odot$ will be on the main sequence if $t_{\rm cc}= 1 \times 10^9$ yr. For the primaries with very little hydrogen in the center at contact, the mergers may never be on the main sequence even in the cases of short $t_{\rm cc}$. The long $t_{\rm cc}$ also delays the appearance of the mergers and shortens their timescales on the main sequence. The latter has not been exactly expressed in our models, and therefore we just see that the mergers from a long $t_{\rm cc}$ have larger luminosities as shown in section 4.} {\bf In our binary evolutions, we have not included AML, which exists in low-mass binaries and may be the main course making the binaries change from detached to contact and finally coalesce, resulting in a large contribution to BSs in old clusters, e.g NGC 2682. In young and intermediate-age clusters, however, AML has little contribution to the birthrate of BSs, since (a) the time is not long enough for binaries to go from detached to contact and (b) the mass of the mergers is probably less than the turnoff of the cluster even though their parents may coalesce in the cluster age. So, we simply estimated the effect of AML in NGC 2682 while negecting it in NGC 2660.} The mass loss during the merger process can also affect our result, mainly the location on the CMD of the products. As shown in NGC 2660, no BSs have been observed in the region with the most opportunity from our models. Because of mass loss, the mergers will be fainter than those given in the paper. However the faintness will be slight since the mass loss is not vast from both observations and smooth particle hydrodynamic simulations \cite{lrs96,sill97,sill01}. The lost mass may carry some angular momentum out from the parent binary. By analyzing the BS spectra from Hubble Space Telescope (there is an apparent continuum deficit on the short-wavelength side of Balmer discontinuity ), De Marco et al. \shortcite{dm04} argued that some BSs might be surrounded by a circumstellar disk. However, Porter \& Townsend \shortcite{pt05} showed that the flux deficits may be attributed wholly to rapid rotation. The rotation rates needed are of the order of those found in the study of De Macro et al. \shortcite{dm05}. Whether the flux deficits shortward of the Balmer jump are induced by a circumstellar disk or rapid rotation, it provides a possible explanation for the orbital angular momentum of the system after coalescence. Such a large mass loss as shown in NGC 2660 (about $0.5 M_\odot$), however, is a problem and should be explained reasonably in physics. Based on some assumptions, we studied the mergers of close binaries from AS and AR evolution by detailed evolutionary calculations. The products from our models may stay on the left of the ZAMS and have no central concentration with decreasing mass. Because of the {\bf development} of the convective core, the mergers with little central hydrogen (less than 0.01) in our models have unusually long timescales on the main sequence ($\sim 10^8$ yrs). These objects are probably subgiants as they are formed, since the primaries in the progenitors also have little central hydrogen and may have left the main sequence during merger process. The mergers from our models stay on the main sequence for a timescale in order of $10^8$ yrs. Some low-mass mergers may stay on the MS for about $10^9$ yrs. The timescale is similar to that of W UMa stars from observations, and therefore we may roughly estimate the contribution to BSs from AS and AR via the number of W UMa systems in a cluster. The estimation, however, is not absolutely since both of the two timescales have wide ranges and large uncertainties, and we cannot rule out other methods for creating W UMa systems except for AS and AR. Comparison to observations indicates that our models (binary coalescence from AS and AR) are not important for the produce of BSs in old open clusters, while likely play a critical role in some younger open clusters. We performed Monte Carlo simulations to examine our models in an old open cluster NGC 2682 and in an intermediate-age cluster NGC 2660. The effect of AML was estimated in NGC 2682 in a simple way, where the mergers are replaced with ZAMS models. In NGC 2682, binary mergers from our models cover the region with high luminosity and those from AML are located in the region with low luminosity, existing a certain width. The BSs from AML are much more than those from our models, indicating that AML of low mass binaries makes a major contribution to BSs in this cluster. Our models are corresponding for several BSs in NGC 2660. In the region with the most opportunity on CMD, however, no BSs have been observed. {\bf Our results are well-matched to the observations if there are $\sim 0.5M_\odot$ of mass loss in the merger process, but a physical mechanism for this much mass loss is a problem.} | 7 | 10 | 0710.4601 |
0710 | 0710.5607_arXiv.txt | Extreme gravitational lensing refers to the bending of photon trajectories that pass very close to supermassive black holes and that cannot be described in the conventional weak deflection limit. A complete analytical description of the whole expected phenomenology has been achieved in the recent years using the strong deflection limit. These progresses and possible directions for new investigations are reviewed in this paper at a basic level. We also discuss the requirements for future facilities aimed at detecting higher order gravitational lensing images generated by the supermassive black hole in the Galactic center. | All known astrophysical cases of bending of photon trajectories by gravitational fields can be interpreted using the conventional weak deflection paradigm, originally formulated by Einstein \cite{Ein}. Yet it is clear that photons passing very close to black holes suffer very large deflections, which must be addressed in a full general relativistic context. However, integrating null geodesics in full general relativity usually leads to very complicated analytical formulae or heavy numerical codes, which typically obscure the basic physical interpretation. The Strong Deflection Limit (SDL) is an analytical tool that allows to derive simple analytical formulae describing the higher order images appearing in extreme gravitational lensing. In this work we review the main progresses achieved in the recent years in this technique and its application to the most interesting physical case: the black hole in the Galactic center (Sgr A*). All relevant formulae derived in previous works are recalled and restated here in the simplest possible form. This work is structured as follows: \S~2 explains the basic phenomenology of extreme lensing; \S~3 introduces the SDL for spherically symmetric black holes; \S~4 generalizes the method to spinning black holes; \S~5 applies the formulae to Sgr A*, examining possible sources for extreme gravitational lensing; \S~6 generalizes the method to sources very close to the black hole; \S~7 contains the conclusions. | Extreme gravitational lensing is a very spectacular though elusive phenomenon, which demands a great effort in order to be observed. The most promising extreme lens is represented by the supermassive black hole in the Galactic center, which anyway requires resolutions of order microarcseconds for the direct observation of higher order images. Suitable known sources for gravitational lensing are giant stars in the IR band and low mass X-ray binaries in the X-ray band. In this paper we have reviewed recent results on gravitational lensing in the Strong Deflection Limit, which is an approximation devoted to the analytical determination of the properties of higher order images. Within this framework, it has been proved that the logarithmic divergence in the deflection angle is a universal feature of all black hole metrics. There exists a general method to determine the SDL coefficients of the deflection angle for any given spherically symmetric black hole metric. This method has been applied to numerous metrics. Higher order images are formed by photons performing one or more loops around the black hole before reaching the observer. The position and the flux ratios of higher order images depend on the specific metric through the SDL coefficients and are thus able to track any possible deviations from general relativity in the strong field regime. The SDL can be extended to spinning black holes, where several new features emerge, such as extended and shifted caustics and additional images. The position and the shape of the caustics can be expressed by simple analytical formulae and the lens equation for sources near to caustics, which represent the most physically relevant situation, can be solved. The extension of the SDL to the case of sources very close to black holes opens the way to even more interesting investigations of extreme gravitational lensing. In fact, whereas direct observation of higher order images is very difficult and probably very far to come in the future, the contribution of higher order images to currently observed phenomena involving sources very close to black holes is already measurable at present time. Light curves of flares born in the accretion disk and spectral measurements of Iron K-lines are heavily influenced by the presence of higher order images. With the SDL setup in its updated version, it is possible to study such phenomena, bearing in mind that the large model dependence remains the main uncertainty in all theoretical attempts to interpret the complicated physics of the supermassive black holes environment. | 7 | 10 | 0710.5607 |
0710 | 0710.5431_arXiv.txt | Warm dark matter is consistent with the observations of the large-scale structure, and it can also explain the cored density profiles on smaller scales. However, it has been argued that warm dark matter could delay the star formation. This does not happen if warm dark matter is made up of keV sterile neutrinos, which can decay into X-ray photons and active neutrinos. The X-ray photons have a catalytic effect on the formation of molecular hydrogen, the essential cooling ingredient in the primordial gas. In all the cases we have examined, the overall effect of sterile dark matter is to facilitate the cooling of the gas and to reduce the minimal mass of the halo prone to collapse. We find that the X-rays from the decay of keV sterile neutrinos facilitate the collapse of the gas clouds and the subsequent star formation at high redshift. | Both cold and warm dark matter models agree with the observed structure on the large scales. However, there are several inconsistencies between the predictions of the cold dark matter (CDM) model and the observations~\cite{cdm_problems}. The low cutoff in dark matter contents of dwarf spheroids, the smoothness of our dark matter halo, and the old globular clusters (observed in Fornax) resisting the infall into the center by dynamical friction~\cite{cdm_problems}, all can be explained by warm dark matter (WDM) because it suppresses the structure on scales that are smaller than the free-streaming length. While the suppression of the small-scale structure is desirable, it has been argued that ``generic'' WDM (for example, gravitino) can slow down structure formation and delay reionization of the universe, which can lead, in turn, to an inconsistency with the reionization redshift obtained by the WMAP \cite{Yoshida}. This problem can be alleviated in the case of the WDM in the form of sterile neutrinos with mass of several keV and a small mixing angle with the ordinary neutrino \cite{bier,stas,stas07} because such sterile neutrinos can decay and produce photons that catalyze the formation molecular hydrogen and speed up the star formation. In the absence of metals, gas cooling is mainly due to the collisional excitation of H$_{2}$, its subsequent spontaneous de-excitation, and photon emission. In the primordial gas clouds, hydrogen molecules can be formed only in reactions involving $e^{-}$ or H$^+$ as a catalyst. Thus, an X-ray radiation can increase the production of the H$_{2}$ by enhancing the ionization fraction, which subsequently leads to speed up of the gas cooling and star formation. Although sterile neutrinos are stable on cosmological time scales, they nevertheless decay. The decay channel important for us is that of decay into one active neutrino and one photon, i.e., $\nu_s \rightarrow \nu_a \gamma$, where the photon energy is half of the sterile neutrino mass, $E_{0} \approx m_{s}c^2/2$. These decays produce an X-ray background radiation that increases the production of molecular hydrogen and can induce a rapid and prompt star formation at high redshift. Sterile dark matter has a firm motivation from particle physics~\cite{dw,Kusenko:2006rh}. The discovery of the neutrino masses implies the existence of right-handed gauge-singlet fields, all or some of which can be lighter than the electroweak scale. These sterile neutrinos can be produced in the early universe by different mechanisms, for example, from neutrino oscillations~\cite{dw} or from the Higgs decays~\cite{Kusenko:2006rh}, or from the couplings to a low-scale inflaton~\cite{Shaposhnikov:2006xi}. The same particles, produced in a supernova, could account for the supernova asymmetries and the pulsar kicks~\cite{kus}, and can play a role in the formation of super-massive black holes in the early universe~\cite{puzzle}. We will examine the thermal evolution of the gas clouds, taking into account both effects of the sterile neutrino decays, namely, the ionization and heating of the gas. We follow the evolution of the baryonic top-hat overdensity, the gas temperature and the H$_{2}$ and $e^{-}$ fraction. In order to perform the calculation we have incorporated to our previous code \cite{stas}, the effects of sterile neutrino decays within collapsing halos and absorption of the X-ray background from sterile neutrinos by He atoms in the intergalactic medium. Our goal is to juxtapose the evolution of the gas temperature in the primordial clouds in the CDM model and the WDM model with keV sterile neutrinos and estimate of the minimal mass able to collapse at a given redshift. | 7 | 10 | 0710.5431 |
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0710 | 0710.2027_arXiv.txt | { The exciting results from H.E.S.S. point to a new population of $\gamma$-ray sources at energies E$>$10~TeV, paving the way for future studies and new discoveries in the multi-TeV energy range. Connected with these energies is the search for sources of PeV cosmic-rays (CRs) and the study of multi-TeV $\gamma$-ray production in a growing number of astrophysical environments. {\em TenTen} is a proposed stereoscopic array (with a suggested site in Australia) of modest-sized (10 to 30m$^2$) Cherenkov imaging telescopes with a wide field of view (8$^\circ$ to 10$^\circ$ diameter) optimised for the E$\sim$10 to 100~TeV range. {\em TenTen} will achieve an effective area of $\sim$10~km$^2$ at energies above 10~TeV. We outline here the motivation for {\em TenTen} and summarise key performance parameters. } \email{[email protected]} \begin{document} | We have outlined the motivation for a new array of IACTs achieving 10~km$^2$ at $E>10$~TeV and described some important performance parameters. This array, known as {\em TenTen}, could also be considered complementary to future MeV to TeV $\gamma$-ray instruments such as GLAST, HESS-II and MAGIC-II. Studies are currently underway to further optimise individual telescopes (optics, electronics, camera design), overall layout parameters, and site potential in Australia. \scriptsize | 7 | 10 | 0710.2027 |
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0710 | 0710.5788_arXiv.txt | The observed increase in star formation efficiency with average cloud density, from several percent in whole giant molecular clouds to $\sim30$\% or more in cluster-forming cores, can be understood as the result of hierarchical cloud structure if there is a characteristic density as which individual stars become well defined. Also in this case, the efficiency of star formation increases with the dispersion of the density probability distribution function (pdf). Models with log-normal pdf's illustrate these effects. The difference between star formation in bound clusters and star formation in loose groupings is attributed to a difference in cloud pressure, with higher pressures forming more tightly bound clusters. This correlation accounts for the observed increase in clustering fraction with star formation rate and with the observation of Scaled OB Associations in low pressure environments. ``Faint fuzzie'' star clusters, which are bound but have low densities, can form in regions with high Mach numbers and low background tidal forces. The proposal by Burkert, Brodie \& Larsen (2005) that faint fuzzies form at large radii in galactic collisional rings, satisfies these constraints. | \label{sect:intro} Stars form in concentrations with a range of densities, from star complexes, OB associations, and T Tauri associations at the low end to compact clusters and super-star clusters (SSC) at the high end. The overall structure is usually hierarchical (Scalo 1985; Feitzinger \& Galinski 1987; Ivanov et al. 1992; Gomez et al. 1993; Battinelli, Efremov \& Magnier 1996; Bastian, et al. 2005, 2007; Elmegreen et al. 2006; see reviews in Efremov 1995; Elmegreen et al. 2000, Elmegreen 2005), and this hierarchy continues even inside the youngest clusters (Testi et al. 2000; Heydari-Malayeri et al. 2001; Nanda Kumar, Kamath, \& Davis 2004; Smith et al. 2005a; Gutermuth et al. 2005; Dahm \& Simon 2005; Stanke, et al. 2006; see review in Allen et al. 2006). Most likely, the hierarchy in stars comes from a hierarchy in the gas (St\"utzki et al. 1998; Dickey et al. 2001), which is the result of turbulence compression and gravitational contraction that is self-similar over a wide range of scales (see review in Mac Low \& Klessen 2004). Clusters form in the densest parts of this gas and lose their initial sub-structure as the stellar orbits mix (for simulations of star formation in clusters, see Klessen \& Burkert 2000; Bonnell, Bate, \& Vine 2003; Li, et al. 2004; Tilley \& Pudritz 2004; Klessen et al. 2005; V\'azquez-Semadeni, Kim, \& Ballesteros-Paredes 2005; Bonnell \& Bate 2006; Li \& Nakamura 2006). Hunter (1999) and Ma\'iz-Apell\'aniz (2001) noted that some massive star-forming regions (which they called ``scaled OB associations'', or SOBA's) do not form dense clusters while others with the same total mass do (e.g., the SSC's). We would like to understand this difference. Obviously the density of the gas is involved, as dense clusters require dense gas, but the distinction between SSCs and SOBAs should also be related to the efficiency of star formation, because clusters forming at low efficiency disperse quickly when the gas leaves (Lada, Margulis \& Dearborn 1984). At very low efficiency, stars form individually without passing through an embedded cluster phase. Recent observations of giant molecular clouds (GMCs) show star formation at both high and low densities, with some embedded stars in dense clusters and others more dispersed (Megeath et al. 2004; J{\o}rgensen et al. 2006, 2007). Larsen \& Brodie (2000) discovered ``faint fuzzies'' at intermediate radii in the disks of the S0 galaxies NGC 1023 and NGC 3384, and suggested they are old clusters with unusually large radii (7-15 pc) and low densities. They are gravitationally bound because of their large ages (8-13 Gyr; Brodie \& Larsen 2002), and they are as massive as SSC's and halo globular clusters. Such clusters appear to represent an intermediate stage between dispersed and bound star formation. They appear to be too low in average density to have had time for core collapse and envelope expansion as in models of globular clusters by Baumgardt et al. (2002). Burkert, Brodie \& Larsen (2005) suggested they formed in a collisional ring interaction between two galaxies. What determines the relative proportion of dispersed and clustered star formation? Larsen \& Richtler (2000) showed that clustering on a galactic scale, measured as the fraction of uv light in the form of massive young clusters, increases as the star formation rate increases. This could be the result of a selection effect if starbursts are active for less than a cluster destruction time. On the other hand, the clustering fraction could depend on pressure. Higher pressure makes the cool phase of gas denser, which promotes more clustering, and locally high pressures trigger star formation on the periphery of GMCs, making clusters in the dense gas (e.g., Comer\'on, Schneider, \& Russeil 2005; Zavagno et al. 2006). The Larsen \& Richtler correlation could then follow from the mutual correlation between pressure and star formation rate with gas column density. Faint fuzzies are a counter example, however: they formed bound but presumably at low pressure to have such low central column densities now. Can all of these clustering types be understood as a continuum of properties in a universal physical model? There have been several attempts to explain the difference between clustered and distributed star formation based on numerical simulations. Klessen, Heitsch, \& Mac Low (2000) suggested that clusters form in non-magnetic gas when the turbulence driving scale is large. Heitsch, Mac Low \& Klessen (2001) noted that non-magnetic turbulence driven on large scales produces a clustered collapse, while magnetic turbulence in supercritical clouds produces a more distributed collapse. Mac Low (2002) suggested that stars form in clusters when there is no turbulent support and they form disbursed when there is. V\'azquez-Semadeni, Ballesteros-Paredes, \& Klessen (2003) suggest this transition from no global turbulent support to support corresponds to an increase in the Mach number and a decrease in the sonic scale, which is the length where the size-linewidth relation gives a Mach number of unity. Large sonic scale (low Mach number) corresponds to a sonic mass larger than the thermal Jeans mass, which means a lack of global support and the formation of a cluster. Low sonic scale corresponds to the dispersed formation of stars, one for each tiny compressed region where the mass exceeds the local thermal Jeans mass. Li, Klessen \& Mac Low (2003) suggested that the equation of state determines the stellar clustering properties: soft equations produce dense clusters while hard equations produce isolated stars. These suggestions all apply to initially uniform media. External compression of a cloud into a massive dense core can also make a cluster; most young clusters are in high-pressure regions like OB associations. Simulations of turbulent media produce stars in compressed regions that act as seeds for the small-scale gravitational collapses that follow (V\'azquez-Semadeni, Ballesteros-Paredes, \& Klessen 2003; Clark \& Bonnell 2005). These simulations also have probability distribution functions (pdfs) for density that are either log-normal or log-normal with a power-law tail at high density, especially when self-gravity is important (e.g., Li, Klessen \& Mac Low 2003). The efficiency of star formation is then proportional to the fraction of the gas in a dense form. Here we examine variations in this fraction as functions of average density and velocity dispersion, and as a function of the local density inside a cloud. The efficiency is taken to be the ratio of the stellar mass to the gas mass during a complete star-forming event. It generally increases with cloud density from a few percent in GMCs (Williams \& McKee 1997) to several tens of percent in cluster-forming cores (e.g., Lada \& Lada 2003). It may reach $\sim50$\% or more inside the densest star-forming cores. We explain this increase as a result of hierarchical structure, regardless of the dynamics and mechanisms of star formation, and we show that for log-normal or similar density pdf's, as expected in turbulent media, the mass fraction of regions with high efficiency increases with the Mach number and, independently, with the average density. This result may explain the Larsen \& Richtler (2000) correlation as well as the observed variations in clustering properties with pressure. We also show that at high Mach number, bound clusters can form with relatively low average densities, thereby explaining faint fuzzies. These are all consequences of star formation in the dense cores of clouds that are structured by turbulence. They result primarily from the geometry of the gas, which is somewhat universal, and should be nearly independent of the gas dynamics or the strength of the magnetic field. We make an important assumption that gravitational contraction and star formation can occur in regions that are either larger or smaller than the sonic scale. This means we assume that contraction to one or more stars can occur in a supersonically turbulent region. V\'azquez-Semadeni, Ballesteros-Paredes, \& Klessen (2003) suggest that if a cloud is supported by turbulence, then only regions smaller than a sonic scale and more massive than the thermal Jeans mass are unstable to form stars. Padoan (1995) was the first to consider this condition. However, clouds are probably not supported for any significant time by turbulence, and even if they were, it is only necessary that a clump mass exceed the turbulent Jeans mass for self-gravitational forces to exceed inertial forces. GMCs for example, have comparable self-gravitational and turbulent energy densities and yet are much larger than the sonic length. Our assumption is contrary to that in Krumholz \& McKee (2005), who assume the same as V\'azquez-Semadeni et al.. We are consistent with McKee \& Tan (2003), however, as they consider the collapse of a highly turbulent core to make a massive star. Saito, et al. (2006), for example, observe star formation in massive turbulent cores. Thus the sonic length should not provide a threshold for star formation. Our primary condition for star formation is that the gas density exceed some fixed value, taken here to be $10^5$ cm$^{-3}$. This is the density at which HCN gives a nearly constant star formation rate per unit gas mass (Gao \& Solomon 2004a,b; Wu et al. 2005) and at which a variety of microscopic processes conspire to shorten the magnetic diffusion time (Elmegreen 2007). Regions with this density should have a wide range of Mach numbers but a nearly universal efficiency, according to the Gao \& Solomon and Wu et al. observations. The density of $\sim10^5$ cm$^{-3}$, when converted to 5900 M$_\odot$ pc$^{-3}$, is also typical for star clusters, as most of those surveyed by Tan (2007), which span a factor of $10^5$ in mass, have about this average density. The cluster density equals the gas density times the efficiency, and the efficiency has to exceed $\sim10-30$\% for a bound cluster to form. Thus, the gas density for both cluster formation and high efficiency appears to be around $10^5$ cm$^{-3}$ or, possibly, $10^6$ cm$^{-3}$. We note that the long timescale derived for HCN gas as the ratio of the total HCN mass divided by the total star formation rate (Gao \& Solomon 2004b; Wu et al. 2005) is not the duration of star formation in any one place, but is the HCN consumption time. As long as a new HCN region is formed somewhere each time an old HCN region disperses, the HCN consumption time can be long even when each region of star formation lasts for a short time (see Elmegreen 2007 for a discussion of star formation timescales). As a result, the average efficiency of star formation in any one HCN region can be moderately large even if the average HCN consumption time is long. An average efficiency of $\sim5-10$\% would be reasonable considering that most star-forming regions leave unbound clusters after the gas leaves (Lada \& Lada 2003), most regions are observed at half their total ages, and star formation typically accelerates over time (Palla \& Stahler 2000). With constant acceleration, only $\sim1/4$ of the total stars form in the first 1/2 of the total time. The efficiency that determines whether a bound cluster will remain is taken here to be 14.4\% (see below), and the peak efficiency in a single-star core is taken to be 50\% at $10^5$ cm$^{-3}$ density. These values are uncertain and are used here only to illustrate how cluster formation might scale with the velocity dispersion and density of the cool component of the interstellar medium. | Stars form bound clusters where the total efficiency is high. The efficiency should be proportional to the mass fraction of a cloud in the form of dense cores where the individual stars form. This mass fraction increases with the cloud density in hierarchical clouds. It also increases with the Mach number of the turbulence because stronger shocks at higher Mach numbers compress the gas to a wider range of densities. As a result, a log-normal density pdf becomes flatter below the threshold density of star formation, and this means the mass fraction is higher at each density below this value. When the Mach number is high (really, when the dispersion of the density pdf is high), the efficiency is high even at a fairly low average cloud density, and so a high fraction of star formation ends up bound. An extreme example of this trend is the type of cluster called a faint fuzzie. We propose that faint fuzzies form in moderately low density clouds with moderately high Mach numbers. The galactic tidal ring environment proposed by Burkert et al. (2005) is an example of a region that would have such clouds. The mass fraction of clumps at a particular high density of star formation also increases with the average density of the ISM because then the whole density pdf shifts toward higher values. The combination of high densities and high Mach numbers, characteristic of starburst regions, makes for a high fraction of star formation in bound clusters. On the other hand, relatively low Mach numbers and/or low densities should produce stars in a more dispersed way, as in the low-density regions of molecular clouds or in low surface brightness galaxies and regions of galaxies. This combination of parameters corresponds to a low ISM pressure, so we infer that low pressure regions, which means those with a low gas column densities and low star formation rates per unit area, should produce relatively fewer bound clusters and relatively more unbound associations. There is a physical explanation for the trends discussed here. Consider a moderately low density cloud with a low turbulent Mach number. The compressions inside that cloud will be modest and most of them will not reach a level of density or enhanced magnetic diffusion rate that allows gravitational collapse before the compression ends. Then few stars will form and the efficiency will be low (unless the cloud continues to makes these weak compressions for a very large number of crossing times, which seems unlikely). Now consider this same cloud with the same size and average density but with a higher Mach number (this will require a higher ISM pressure). The stronger compressions will more easily produce high-density cores in which gravity is important and magnetic diffusion is fast. More stars will form and the efficiency will be high. The only difference between these two examples is the Mach number, which affects the range of densities in the compressed regions. In terms of the density pdf, higher Mach numbers produce a broader pdf at the same average density. In terms of star formation in bound clusters, high pressure regions having clouds with high Mach number produce a higher fraction of their stars in bound clusters. The model predicts that young bound clusters in starburst regions, or in regions of high ISM pressures or Mach numbers, will have a wider range of densities than bound clusters in more quiescent, low-pressure regions. This is partly because the cloud densities should be higher in starburst regions, so the resulting clusters can be denser overall, but it is also because the lower-density parts of starburst clouds can form stars with a high efficiency, leaving bound clusters when the gas leaves rather than dispersed OB associations. As there is generally more mass at low density than at high density, the net distribution of cluster density could shift toward lower values when the ISM density pdf is broad. According to the model, this density shift is relative to the mean ISM density, and it should be measured only before significant cluster expansion. Low surface brightness galaxies should make a preponderance of unbound OB associations rather than bound clusters, and the young clusters they form should have a relatively narrow range of densities compared to normal. | 7 | 10 | 0710.5788 |
0710 | 0710.0955_arXiv.txt | {} {We present the optical classification and redshift of 348 X-ray selected sources from the XMM-Newton Bright Serendipitous Survey (XBS) which contains a total of 400 objects (identification level = 87\%). About 240 are new identifications. In particular, we discuss in detail the classification criteria adopted for the Active Galactic Nuclei population.} {By means of systematic spectroscopic campaigns and through the literature search we have collected an optical spectrum for the large majority of the sources in the XBS survey and applied a well-defined classification ``flow-chart''. } {We find that the AGN represent the most numerous population at the flux limit of the XBS survey ($\sim$10$^{-13}$ erg cm$^{-2}$ s$^{-1}$) constituting 80\% of the XBS sources selected in the 0.5-4.5 keV energy band and 95\% of the ``hard'' (4.5-7.5 keV) selected objects. Galactic sources populate significantly the 0.5-4.5 keV sample (17\%) and only marginally (3\%) the 4.5-7.5 keV sample. The remaining sources in both samples are clusters/groups of galaxies and normal galaxies (i.e. probably not powered by an AGN). Furthermore, the percentage of type~2 AGN (i.e. optically absorbed AGNs with A$_V>2$mag) dramatically increases going from the 0.5-4.5 keV sample (f=N$_{AGN 2}$/N$_{AGN}$=7\%) to the 4.5-7.5 keV sample (f=32\%). We finally propose two simple diagnostic plots that can be easily used to obtain the spectral classification for relatively low redshift AGNs even if the quality of the spectrum is not good. } {} | In the last few years {\it XMM-Newton} and {\it Chandra} telescopes have represented an excellent tool to survey the hard X-ray sky at all fluxes, from relatively bright (10$^{-13}$ erg cm$^{-2}$ s$^{-1}$, e.g. Della Ceca et al. 2004 and references therein), to medium (10$^{-13}$ erg cm$^{-2}$ s$^{-1}$-10$^{-14}$ erg cm$^{-2}$ s$^{-1}$, e.g. Barcons et al. 2007 and references therein) and deep (10$^{-14}$-10$^{-16}$ erg cm$^{-2}$ s$^{-1}$, Brandt \& Hasinger 2005; Worsley et al. 2005 and references therein) fluxes. At the energies ($\sim$0.5-10 keV) covered by the instruments on board these two telescopes, Active Galactic Nuclei (AGN) can be efficiently selected and studied even when affected by large levels of absorption (up to N$_H\sim$10$^{24}$ cm$^{-2}$, corresponding to an optical absorption of A$_V\sim$500 mag). This important characteristic, combined with the good/excellent spatial and energy resolutions of the detectors, makes the ongoing surveys a fundamental tool for AGN studies. At the same time, these new surveys represent an observational challenge at wavelengths different from the X-ray ones: multiwavelength follow-ups of X-ray sources, particularly in the optical domain, are decisive to derive the distance and to understand the properties of the selected objects but they also require large fractions of dedicated observing time at different telescopes. Probably one of the most challenging and time-consuming efforts is the optical spectroscopic follow-up of the selected sources. One of the primary goals of all these hard X-ray surveys is to explore the population of absorbed AGN and, to this end, an optical classification that can reliably separate between optically absorbed and non-absorbed objects is always required. Two important limits, however, affect the spectroscopic follow-ups of deep and, in part, medium surveys: first, the optical counterparts are often too faint to be spectroscopically observed even at the largest optical telescopes currently available; second, even when a spectrum can be obtained, its quality is not always good enough to provide the critical pieces of information that are required to assess a reliable optical classification. These two problems often limit the final scientific results that are based on the optical classification of medium/deep surveys. On the contrary, bright surveys offer the important possibility of obtaining a reliable optical classification for virtually all (with some exceptions, as discussed in the next sections) the selected sources. The disadvantage of dealing with shallow/wide-angle samples is that the techniques to observe efficiently many sources at once, like Multi-objects or fibers-based methods, cannot be applied for the optical follow-up, given the low space-density of sources at bright X-ray fluxes. The only suitable method, the ``standard'' long-slit technique, requires many independent observing nights to achieve the completion of the optical follow-up. In this paper we present and discuss in detail the optical classification process of the {\it XMM-Newton} Bright Serendipitous Survey (XBS, Della Ceca et al. 2004), which currently represents the widest (in terms of sky coverage) among the existing {\it XMM-Newton} or {\it Chandra} surveys for which a spectroscopic follow-up has been almost completed. The aim of the paper, in particular, is to provide not only a generic classification of the sources and their redshift but also a quantification, in the limits of the available data, of the corresponding threshold in terms of level of optical absorption. The paper is organized as follows: in Section~2 we describe the XBS survey, in Section~3 we describe the process of identification of the optical counterpart, in Section~4 and Section~5 we respectively summarize our own spectroscopic campaigns carried out to collect the data as well as the data obtained from the literature. In Section~6 we shortly discuss the data reduction and analysis of the optical spectra and in Section~7 we give the details on the classification criteria adopted for the sources in the XBS survey. In Section~8 we propose two diagnostic plots that can be used to easily classify the sources into type~1 and type~2 AGN. The resulting catalogue is presented in Section~9 while in Section~10 we briefly discuss the optical breakdown and the redshift distribution of the sources. The conclusions are finally summarized in Section~11 . Throughout this paper H$_0$=65 km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}$=0.7 and $\Omega_M$ = 0.3 are assumed. | We have presented the details of the identification work of the sources in the XBS survey, which is composed by two complete flux limited samples, the BSS and the HBSS sample, selected in the 0.5-4.5 keV and 4.5-7.5 keV band respectively. We have secured a redshift and a spectroscopic classification for 348 (including data from the literature) out of 400 sources, corresponding to 87\% of the total list of sources and to 87\% and 97\% considering the BSS and HBSS samples separately. The results of the identification work can be summarized as follows: \begin{itemize} \item We have quantified the criteria used to distinguish optically absorbed AGN (i.e. type~2) from optically non-absorbed (or moderately absorbed) AGN (type~1) and we have shown that the adopted dividing line between the two classes of sources corresponds to an optical extinction of A$_V\sim$2 mag, which translates into an expected column density of N$_H\sim$4$\times$10$^{21}$ cm$^{-2}$, assuming a Galactic A$_V$/N$_H$ relationship. \item About 10\% of the extragalactic sources (35 objects in total) show an optical spectrum which is highly contaminated by the star-light from the host galaxy. These sources have been studied in detail in a companion paper (Caccianiga et al. 2007). Using the X-ray data we have found an elusive AGN in 33 of these objects and we have classified them into type~1 and type~2 AGN according to the value of N$_H$ measured from the X-ray spectrum. To this end, we have used a N$_H$=4$\times$10$^{21}$ cm$^{-2}$ dividing value which matches (assuming the standard Galactic A$_V$/N$_H$ relation) the value of A$_V$ (=2 mag) adopted with the optical classification. \item We have then proposed two simple diagnostic diagrams. The first one, based on the 4000\AA\ break and the [OIII]$\lambda$5007\AA\ equivalent width, can reliably distinguish between type~1 and type~2 AGN if the host galaxy does not dominate the optical spectrum. The second uses the H$\alpha$ and [OIII]$\lambda$5007\AA\ line equivalent widths to classify into type~1 and type~2 the elusive AGN sources in which a possibly broad H$\alpha$ emission line is detected. \item We find that the AGN represent the most numerous population at the flux limit of the XBS survey ($\sim$10$^{-13}$ erg cm$^{-2}$ s$^{-1}$) constituting 80\% of the XBS sources selected in the 0.5-4.5 keV energy band and 95\% of the ``hard'' (4.5-7.5 keV) selected objects. Galactic sources populate significantly the 0.5-4.5 keV sample (17\%) and only marginally (3\%) the 4.5-7.5 keV sample. The remaining sources in both samples are clusters/groups of galaxies and normal galaxies (i.e. probably not powered by an AGN). \item As expected, the percentage of type~2 AGN dramatically increases going from the 0.5-4.5 keV sample (f=N$_{AGN 2}$/N$_{AGN}$=7\%) to the 4.5-7.5 keV sample (f=32\%). A detailed analysis on the intrinsic (i.e. taking into account the selection effects) relative fraction of type~1 and type~2 AGN will be be presented in a forthcoming paper (Della Ceca et al. 2007, in prep.). \end{itemize} | 7 | 10 | 0710.0955 |
0710 | 0710.2175_arXiv.txt | Voids are a dominant feature of the low-redshift galaxy distribution. Several recent surveys have found evidence for the existence of large-scale structure at high redshifts as well. We present analytic estimates of galaxy void sizes at redshifts $z \sim 5 - 10$ using the excursion set formalism. We find that recent narrow-band surveys at $z \sim 5 - 6.5$ should find voids with characteristic scales of roughly $20$ comoving $\Mpc$ and maximum diameters approaching $40 \Mpc$. This is consistent with existing surveys, but a precise comparison is difficult because of the relatively small volumes probed so far. At $z \sim 7 -10$, we expect characteristic void scales of $\sim 14 - 20$ comoving $\Mpc$ assuming that all galaxies within dark matter haloes more massive than $10^{10} \Msol$ are observable. We find that these characteristic scales are similar to the sizes of empty regions resulting from purely random fluctuations in the galaxy counts. As a result, true large-scale structure will be difficult to observe at $z \sim 7 -10$, unless galaxies in haloes with masses $\la 10^9 \Msol$ are visible. Galaxy surveys must be deep and only the largest voids will provide meaningful information. Our model provides a convenient picture for estimating the ``worst-case'' effects of cosmic variance on high-redshift galaxy surveys with limited volumes. | The complex network of filaments and voids observed in the present-day Universe is believed to have formed from an initially homogeneous distribution of matter. In hierarchical models of structure formation, tiny perturbations seeded by the inflationary epoch grew through gravitational instability, collapsing first on smaller scales to form haloes. The subsequent merging and clustering of smaller haloes resulted in the formation of highly structured large-scale systems. Perhaps the most striking characteristic of the Universe today is the prevalence of large and nearly spherical voids in the galaxy distribution. The scales of these voids can be enormous. Indeed, \citet{HandV2004} report characteristic radii of $R \sim 15h^{-1}~\Mpc$ with maximum scales approaching $R \sim 25 h^{-1} \Mpc$ in the 2dF Galaxy Redshift Survey. The characteristics of voids and the galaxies that populate them have been the subject of numerous theoretical and observational studies throughout the years \citep{gregory78, kirshner81, delapparent86, vogeley94, HandV2004, Conroy2005}. To date, these studies have mostly focused on low redshifts. However, it is clear that voids should appear at higher redshifts also. The DEEP2 survey indicates that voids exist at redshifts of $z \sim 1$ \citep{Conroy2005}. Surprisingly, a handful of recent Lyman $\alpha$ emitter (LAE) surveys have found hints of large-scale structure at redshifts around $z \sim 5$ \citep{Shimasaku2003, Shimasaku2006, Hu2004, Ouchi2005}. The modelling of voids poses an interesting theoretical problem. There have been numerous studies utilising $N$-body simulations \citep{MandW2002,Benson2003,Gottlober2003,Colberg2005}. While these simulations are invaluable tools for understanding the details of void dynamics, they are computationally expensive due to the large volumes and high dynamic range required to include a representative sample of voids while also resolving the much smaller galaxies that define them. Analytic methods provide a useful alternative. Perhaps the most promising analytic model of void abundances is the excursion set approach taken by \citet{SandV2004}. They argue that voids actually provide deeper insight into large-scale structure than halo formation itself. Their assertion is based on the fact that underdense regions generally tend to evolve toward a spherical geometry, making the idealisation of spherical expansion more reasonable. In contrast, gravitationally bound objects typically have geometries that are far from spherical. Approximating gravitational collapse with the spherical model may be highly inaccurate, which partially explains the discrepancies between the \citet{PandS1974} halo mass function and simulations \citep{SandT1999,Sheth2001,Jenkins2001}. A key disadvantage to the approach of \citet{SandV2004} is the difficulty in relating their definition of voids to observational studies. As we will discuss in \S \ref{VoidDef}, \citet{SandV2004} use the dark matter underdensity to define voids. \citet{FandP2006} extend their model to define voids in terms of the local \emph{galaxy} underdensity. They predict characteristic void sizes of $R \sim 10 h^{-1} \Mpc$ at the present day -- nearly as large as observed voids. In what follows, we present analytic estimates of void size distributions at redshifts between $z \sim 5 - 10$. Our aim is to provide a convenient basis of comparison for current and future high-redshift observations -- presumably, though not limited to, LAE surveys. As such, we consider the effects of statistical fluctuations in the galaxy counts and the abundance of Ly $\alpha$ emitting galaxies on void observations. LAE surveys have become an invaluable tool in cosmological studies. In addition to building larger samples at $z \sim 5$, observers have pushed the threshold to redshifts as high as $z \sim 7 - 10$ \citep{WandC2005,Cuby2007,Stark2007,Ota2007}. Surveys at these redshifts could potentially reveal important information on the epoch of reionization. Indeed, the observed clustering properties of LAEs could someday be a powerful probe of the epoch \citep{furl04-lya, Furlanetto2006,McQuinn2006,McQuinn2007, mesinger07}. Regions of ionized hydrogen grow quickly around clustered galaxies as reionization progresses. When these regions are large enough, Ly $\alpha$ photons are sufficiently redshifted before they reach neutral hydrogen gas, allowing them to avoid absorption in the intergalactic medium (IGM). Sources within overdense regions are therefore more likely to be observed relative to void galaxies, resulting in a large-scale modulation of the number density. One method to quantify such clustering is with void statistics, as first attempted by \citet{McQuinn2007}. This provides comparable power to correlation function measurements of the galaxies. However, taking full advantage of this technique requires a deeper understanding of voids in the underlying galaxy distribution; our calculations aim to provide such a baseline model. The remainder of this paper is organized in the following manner. In \S \ref{ExcursionSet}, we briefly review the basic principles of the excursion set formalism. In section \ref{VoidDef}, we present the \citet{FandP2006} definition of voids in terms of the local galaxy underdensity. Section \ref{VoidDistributions} contains the main results of this paper: void size distributions at $z = 4.86 - 10$. In \S \ref{StochasticVoids}, we estimate the typical sizes of voids that result from random fluctuations in the galaxy distribution and develop an alternative definition of voids. In \S \ref{VisibleFrac}, we explore the assumption that only a certain fraction of galaxies are actually visible in LAE surveys. Section \ref{Observations} contains a rough comparison of our calculations to high redshift Ly $\alpha$ surveys. Finally, we offer concluding remarks in \S \ref{Discussion}. In what follows, we assume a cosmology with parameters $\Omega_m = 0.24,~\Omega_{\Lambda} = 0.76,~\Omega_b = 0.042,~H = 100h~\mathrm{km~s}^{-1}~\Mpc^{-1}$ (with $h = 0.73$), $n = 0.96$, and $\sigma_8 = 0.8$, consistent with the latest measurements \citep{Spergel2007}. All distances are reported in comoving units. | \label{Discussion} We have calculated void size distributions at $z = 4.86$--$10$ using the excursion set model developed by \citet{SandV2004} and \citet{FandP2006}. The latter found characteristic void radii of $R \approx 7-14 \Mpc$ at $z = 0$. For the observational sensitivities assumed in this paper, we obtained characteristic void radii that are very similar: $R \approx 7-10 \Mpc$ for redshifts between $z = 4.86$ and $z = 10$. These results are virtually independent of the void-crushing barrier (for any reasonable choice). We have shown that characteristic void scales actually increase with redshift for a fixed halo mass threshold due to a decreased number density and increased bias with respect to the underlying matter density. Following recent studies on the abundances of low-redshift LAEs, we explored the possibility that only a fraction $f_{vis}$ of galaxies are sampled in LAE surveys. This has only a small effect on the void size distribution but increases the contamination of void samples by stochastic fluctuations. In section \ref{StochasticVoids}, we have explored stochastic fluctuations in the galaxy distribution. These fluctuations, although inherently different from the "real" voids we model in this paper, will result in large empty regions in the sky. Stochastic voids can therefore contaminate real void samples and lead to erroneous conclusions on the formation of large-scale structure. We have estimated the typical scale of these regions to be slightly smaller than the characteristic scale of true voids at $z \sim 5$. At $z \sim 10$, the situation depends on the particular choice of $m_{min}$. For $m_{min} \sim 10^{10} \Msol$, stochastic voids are typically the same scale as real voids. The increased importance of stochastic fluctuations will make the identification of large-scale structure at this redshift difficult. Attempts to do so must observe halos near the minimum mass to form stars, $\sim 10^8$--$10^9 \Msol$, in order for true voids to dominate the observed distribution. We found that a large fraction of real voids in our fiducial model contain no visible galaxies, adding to the difficulties in differentiating them from stochastic fluctuations. We have presented a modified definition of voids that incorporates both stochastic and real voids and so is easier to compare to the limited observational samples thus far available. In our new approach, we defined voids in terms of the probability for a region to be empty. We found that the modified void distributions are more sharply peaked and have characteristic scales that are comparable to the fiducial model. We have also attempted to visually compare our results to the most recent narrow-band filter surveys at $z = 5.7$. While we found no inconsistencies, it is difficult to draw any decisive conclusions because of small-number statistics, projection effects, and lower-redshift contaminants. Obviously, a more systematic approach is required. Future surveys promise to provide better statistics and increased sample volumes for studies on high-redshift voids. In the context of next generation surveys for high-redshift galaxies, our model is useful for gauging the impact of cosmic variance. Consider a fictitious survey at $z = 10$ with a detection threshold of $m_{min} = 10^{10} \Msol$. Figure \ref{PoissonPlot} illustrates that stochastic voids with $R \sim 10 - 20 \Mpc$ will dominate the sky distribution. Therefore, one must either search for voids with $R > 20 \Mpc$ or search deeper for significantly smaller sources. The latter may be possible if the sources observed by \citet{Stark2007} are indeed at $z \sim 9$, in which case they imply that halos near $\la 10^9 \Msol$ are visible \citep{mesinger07}. However, high-redshift galaxies are so highly biased that even with deep observations, a substantial fraction of the Universe is filled with empty or nearly-empty regions. For example, at $z=10$ and $m_{min} = 10^{10} \Msol$, $\sim 37\%$ of space is filled by regions that are at least 80\% underdense in galaxies and at least 20 Mpc across -- or fully 7 arcmin. With the small fields of view available to near-infrared detectors, this suggests that either many independent fields must be observed or a large contiguous volume surveyed to be guaranteed of detecting a reasonable number of sources. Finally, we have neglected reionization and its effect on the appearance of large-scale structure. Regions of neutral hydrogen are expected to modulate the LAE density on large scales and accentuate the appearance of structure \citep{furl04-lya, Furlanetto2006,McQuinn2006,McQuinn2007, mesinger07}. Although the precise time frame is currently unknown, quasar observations and cosmic microwave background measurements have provided some evidence that reionization occurred between $z \sim 6 - 10$ (e.g, \citealt{Fan2006,Page2007,MandH2004,MandH2007}). Interestingly, \citet{Kashikawa2006} found a significant high-luminosity suppression in the LAE luminosity function between $z = 5.7$ and $z = 6.5$. Whether or not reioinization is responsible for this effect is currently unclear \citep[no such suppression was observed by][]{dawson07}. Because IGM absorption modulates the LAE density on large scales, we would expect reionization to have a substantial effect on the observed void sizes in such narrow-band surveys (it should not affect galaxies identified through broadband effects). Of course, the plots in Figure \ref{Z7to10} provide analytic estimates only of the \emph{intrinsic} void size distributions. They provide a basis for comparison with high-redshift surveys in order to determine whether the observed features are easily attributable to the large-scale clustering alone. It therefore helps illuminate efforts to use voids to constrain the IGM properties during reionization, as first attempted by \citet{McQuinn2007}. | 7 | 10 | 0710.2175 |
0710 | 0710.3345_arXiv.txt | Determining the final spin of a black-hole (BH) binary is a question of key importance in astrophysics% . Modelling this quantity in general is made difficult by the fact that it depends on the 7-dimensional space of parameters characterizing the two initial black holes. However, in special cases, when symmetries can be exploited, the description can become % simpler. For black-hole binaries with unequal masses but with equal spins which are aligned with the orbital angular momentum, we show that the use of recent simulations and basic but exact constraints derived from the extreme mass-ratio limit allow to model this quantity with a simple analytic expression. Despite the simple dependence, the expression models very accurately all of the available estimates, with errors of a couple of percent at most. We also discuss how to use the fit % to predict when a Schwarzschild BH is produced by the merger of two spinning BHs, when the total angular momentum of the spacetime ``flips'' sign, or under what conditions the final BH is ``spun-up'' by the merger. % Finally, suggest an extension of the fit % to include unequal-spin binaries, thus potentially providing a \textit{complete} description of the final spin from the coalescence of generic black-hole binaries with spins aligned to the orbital angular momentum. | \label{intro} The determination of the final spin of a BH binary is a question of key importance in astrophysics. Modelling this in general is made difficult by the fact that it depends on the 7-dimensional space of parameters characterizing the two initial BHs. However, in special cases, when symmetries can be exploited, the description can be much simpler. Several recent studies have shed light on the remnant of the merger process. Using conservation principles, Hughes and Blandford~\citep{Hughes_Blandford:2003} argued that mergers rarely lead to rapidly rotating objects. \citet{Gonzalez:2006md} numerically evolved a sequence of non-spinning unequal-mass BHs, arriving at detailed estimates of the radiated energy and angular momentum. In a series of papers~\citep{Koppitz:2007ev, Pollney:2007ss, Rezzolla_etal:2007a} we have studied the parameter space of mergers of equal-mass BH binaries whose spins are aligned with the orbital angular momentum but otherwise arbitrary. The findings agree well with independent numerical evolutions~\citep{Campanelli_etal:2007,Herrmann:2007ac}, as well as more recent studies of models with initial spins up to $J/M^2=0.8$~\citep{Marronetti_etal:2007}. An important result of these studies has been the determination of simple (quadratic) fitting formulas for the recoil velocity and spin of the merger remnant as a function of the initial BH parameters~\citep{Rezzolla_etal:2007a}. A number of analytical approaches have been developed over the years to determine the final spin of a binary coalescence~\citep{ Damour:2001tu, Buonanno_Damour:2000, Buonanno_etal:2006, Damour_Nagar:2007a, Boyle:2007sz}. Very recently, an interesting method, inspired by the dynamics of a test particle around a Kerr BH, has been proposed for generic binaries~(\citet{Buonanno_etal:2007b}, BKL hereafter). The approach assumes that the angular momentum of the final BH is the sum of the individual spins and of the orbital angular momentum of a test particle on the last-stable orbit of a Kerr BH with the same spin parameter as that of the final BH. Here, we combine the data obtained in recent simulations to provide a phenomenological but analytic estimate for the final spin in a binary BH system with arbitrary mass ratio and spin ratio, but in which the spins are constrained to be parallel to the orbital angular momentum. Our numerical simulations have been carried out using the CCATIE code~\citep{Pollney:2007ss}. In addition to the data presented in \citet{Rezzolla_etal:2007a}, we add three simulations of equal-mass, high-spin binaries and three simulations of unequal-mass, spinning binaries (see Table~\ref{tableone}). Other data is taken from % unequal-mass, nonspinning binaries~\citep{Gonzalez:2006md,Berti_etal:2007,Buonanno_etal:2007a}, and of equal-mass, spinning binaries~\citep{Rezzolla_etal:2007a,Marronetti_etal:2007}; all of the AEI data is summarized in Table~\ref{tableone}. To avoid the possible contamination from the errors associated with high-spin binaries reported by~\citet{Marronetti_etal:2007}, we have not considered binaries with initial spin $|J/M^2| \geq 0.75$ reported in the literature~\citep{Campanelli_etal:2007,Marronetti_etal:2007}. We have, however, considered estimates of high-spin binaries (\textit{cf.}, Table~\ref{tableone}), for which we know the spins remain essentially constant prior to merger, with changes less than $0.5\%$~\citep{Pollney:2007ss}, and that are very well captured by the fit. | Modelling the final spin in a generic binary BH merger is not trivial given the large space of parameters on which this quantity depends. We have shown that the results of recent simulations combined with basic but exact considerations derived from the EMRL allow us to model this quantity with a simple analytic expression in the case of BH binaries having unequal masses and unequal spins which are aligned with the orbital angular momentum. When compared with all other estimates coming either from numerical calculations or from approximation techniques, the estimates of the 2D fit show differences which are of few percent at most. \bigskip \noindent We % thank A. Buonanno, T. Damour, S. A. Hughes, L. Lehner, A. Nagar, and B. S. Sathyaprakash for discussions. We are grateful to D. Merritt for pointing out an error in the interpretation of our results. The computations were performed on the supercomputers at AEI, LITE, LSU, LONI and NCSA. Support comes also through the DFG grant SFB/TR% ~7. | 7 | 10 | 0710.3345 |
0710 | 0710.1340_arXiv.txt | Optically thin two-temperature accretion flows may be thermally and viscously stable, but acoustically unstable. Here we propose that the O-mode instability of a cooling-dominated optically thin two-temperature inner disk may explain the 23-day quasi-periodic oscillation (QPO) period observed in the TeV and X-ray light curves of Mkn~501 during its 1997 high state. In our model the relativistic jet electrons Compton upscatter the disk soft X-ray photons to TeV energies, so that the instability-driven X-ray periodicity will lead to a corresponding quasi-periodicity in the TeV light curve and produce correlated variability. We analyse the dependence of the instability-driven quasi-periodicity on the mass (M) of the central black hole, the accretion rate ($\rm{\dot{M}}$) and the viscous parameter ($\alpha$) of the inner disk. We show that in the case of Mkn~501 the first two parameters are constrained by various observational results, so that for the instability occurring within a two-temperature disk where $\alpha=0.05-1.0$, the quasi-period is expected to lie within the range of 8 to 100 days, as indeed the case. In particular, for the observed 23-day QPO period our model implies a viscosity coefficient $\alpha \leq 0.28$, a sub-Eddington accretion rate $\dot{M} \simeq 0.02\,\dot{M}_{\rm Edd}$ and a transition radius to the outer standard disk of $r_0 \sim 60 \,r_g$, and predicts a period variation $\delta P/P \sim 0.23$ due to the motion of the instability region. | Strong variability is one of the common observational properties of TeV emitting Active Galactic Nuclei (AGNs) \cite{cat99}. In many cases, the highly variable high-energy { gamma-rays and the X-rays appear to be correlated with no time delays evident on day-scales}, suggesting that the $\gamma$-rays { may} result from inverse Compton upscattering of lower energy photons. { The first TeV blazar, for example, that was observed simultaneously in multiple bands from radio to TeV gamma-rays is Mrk~421. The first campaign, conducted in 1994 \cite{mac95} on Mrk~421, revealed correlated variability between the keV X-ray and TeV gamma-ray emission. The gamma-ray flux varied by an order of magnitude on a timescale of 2 days and the X-ray flux was observed to double in 12 hr. On the other hand, the high-energy gamma-ray flux observed by EGRET, as well as the radio and UV fluxes, showed less variability than the keV or TeV bands. Another multiwavelength campaign on Mrk~421 performed in 1995 revealed another coincident keV/TeV flare \cite{buk96,taka96}. The UV and optical bands also showed correlation during the flares. With the detection of TeV emission from Mrk~501 \cite{qui96}, several multiband campaigns were organized on Mrk~501 to verify whether such a behavior is a general property of TeV-emitting blazars or whether it is unique to Mrk 421.} The gamma-ray blazar Mkn~501, detected as a strong TeV emitter in 1995 \cite{qui96}, is one of the closest ($z = 0.0337$) and brightest BL Lacertae objects. Historically (i.e., prior to 1997), its spectral energy distribution (SED) $\nu\,F_{\rm \nu}$ resembles that of X-ray selected BL Lac objects, having a peak in the extreme UV--soft X-ray energy band \cite{kat99,sam96}. { Earlier} optical observations of Mkn~501 have shown variations of up to $1.^{m}3$ and polarized emission up to $P_{\rm opt} \simeq 7 \%$ \cite{fan99}. During its active state in 1997, where Mkn~501 was monitored in the 2-10 keV X-ray band by the all sky monitor(ASM) on board RXTE and in the TeV energy band by several Cherenkov telescope groups \cite{aha99a,cat97,dja99,kraw00,pro97,rxte99}, both X-rays and TeV gamma-rays increased by more than 1 order of magnitude from quiescent level \cite{cat97,pia98}. Analysis of the X-ray and TeV data showed, that the variations were strongly correlated between both bands, { yet only marginally correlated with the optical UV band} \cite{cat97,dja99,kraw00,pet00}. { While the synchrotron emission peaked below 0.1 keV in the quiescent state, in 1997 it peaked at ~100 keV. This is the largest shift ever observed for a blazar \cite{pia98}. Earlier investigations of the 1997 April flare in Mkn~501 \cite{pia98}, showed that its origin may be related to a variation of $\gamma_{max}$ together with an increased luminosity and a flattening of the injected electron distribution.} During the 1997 high state, the X-ray and TeV light curves displayed a quasi-periodic signature \cite{hay98,kra99,pro97}. A detailed periodicity analysis, based on the TeV measurements obtained with all Cherenkov Telescopes of the HEGRA-Collaboration and the X-ray data with RXTE, was performed by Kranich et al. \cite{kra99,kra01} and more recently also by Osone \cite{oso06} including Utah TA data. The results indeed strongly suggest the existence of a 23 day periodicity, with a combined probability of $p \simeq 3 \times 10^{-4}$ in both the TeV and the X-ray light curves covering the same observational period \cite{kra99,kra01}. \footnote{{ Note that no QPOs have been seen by MAGIC during 1998-2000, when the source was not very active \cite{alb07}, suggesting that the QPOs only occur during a very active stage.}} { Rieger and Mannheim ~(2000) have shown that the origin of these QPOs may be related to the presence of a binary black hole in the center of Mkn~501. While this may well be possible, we explore here an alternative scenario, where the observed QPOs are related to accretion disk instabilities.} In a seminal paper, Shapiro et al.~(1976) (SLE) have pointed out that a hot (Compton cooling-dominated) optically thin two-temperature accretion disk may be present in the inner region of the standard optically thick disk \cite{sha73}, whenever the SLE inequality $\alpha^{1/4} \dot{M}_{*}^{2}M_{*}^{-7/4} \geq 0.6$ is satisfied, where $\dot{M}_{*} = \dot{M}/(10^{17} \rm{g} \rm{s}^{-1})$, $M_{*}=M/(3 M_{\odot})$, $\dot{M}$ is the accretion rate of the disk, $M$ the mass of the central black hole, and $\alpha$ the viscosity parameter, constrained by the model to lie within the range $ 0.05 < \alpha < 1$. Later work \cite{pir78,pri76} has shown that the SLE configuration might be thermally unstable (although less than the standard disk) if the simple standard viscosity prescription is employed. However, relatively small changes in the viscosity law can already ensure a stable configuration \cite{pir78}, in particular, if stabilizing effects of a disk wind are fully taken into account. A kind of SLE two-temperature disk structure may thus well exist in the inner region of BL Lac type objects \cite{wan91,wanu91,cao03} and provide a possible explanation for the X-ray and TeV variability phenomenon in Mkn~501. For, firstly, Compton processes in a inner two-temperature disk with electron temperature $T_e$ of about $10^{9}$ K (and ion temperature one or two orders of magnitude higher) can produce (steep) X-ray power-law spectra, in contrast to the standard disk model that can only produce emission up to the optical-UV band \cite{sha76,wan91}. Secondly, analysis of the linear stability of an optically thin two-temperature disk around a compact object shows that the disk is subject to a radial pulsational instability (inertial-acoustic mode instability) \cite{wu97}. This possible mode of pulsational overstability, in which radial disk oscillations with local Keplerian frequencies become unstable against axisymmetric perturbations, occurs if the viscosity coefficient increases sufficiently upon compression \cite{kat78}. In this case, thermal energy generation due to viscous dissipation increases during compression, leading to amplification of the oscillations analogous to the role played by nuclear energy generation in stellar pulsations. As we demonstrate below, the occurrence of such a type of disk oscillation may well account for some quasi-periodic time variability in AGNs in a way similar to Galactic black hole candidates \cite{blu84,hon92,man96,yan97}. | The X-ray to TeV spectra of TeV blazars have been often interpreted within a one-zone SSC model, e.g. \cite{blo96,der97,mas97,tav98,tav01}. However, if a two-temperature disk structure is present in some of these sources, the real situation may be more complex as the X-ray radiation from a two-temperature disk, for example, may represent an additional, non-negligible source of seed photons for the inverse Compton scattering to TeV energies, in particular during active source stages, likely to be associated with changes in accretion history. Hence our model assumes, that a pulsational instability occurring within a two-temperature disk leads to observable, periodic variations of its X-ray radiation field. Part of this periodically modulated X-ray emission will enter the jet and (in addition to direct synchrotron photons) serves as seed photons for Compton-upscattering to TeV energies { similar as in \cite{der92,der93}}. Our model thus takes it that both, a direct synchrotron self Compton and an external Compton contribution are relevant for modelling the SED of Mkn~501, the seed photons for external Compton-upscattering consisting of both, the infrared-optical seed photons from the (quasi-steady) SS disk component and the variable X-ray photons from the two-temperature disk component. A related, but more simpler scenario, assuming the seed photons for inverse Compton scattering to be provided by a (direct) synchrotron plus an quasi-steady flux component comparable to the observed infrared-optical flux, has been proposed by Pian et al.~(1998) in order to account for the different degrees of SED variations of Mkn~501 at X-ray and sub-X-ray energies during the April 1997 outburst (see also \cite{ghi98,kat99}). Detailed analysis indeed suggests that one-component SSC models cannot fit both the April 1997 SEDs and the lightcurves from X-ray to TeV, and that (at least) an additional, moderately variable low energy component contributing in the energy range between 3-25 keV is required \cite{kraw02,mas04}. A similar conclusion seems to hold for the strong X-ray outburst observed in July 1997 \cite{lam98}. Interestingly, observations in 1998 also provide evidence for an additional component in the optical regime, possible associated with the SS disk component in our hybrid disk model { (see also \cite{kat01} for an alternative interpretation)}: As shown by Massaro et al.~\cite{mas04} optical to X-ray data taken in June 1998 indicate that the optical spectrum is steep and does not match the low energy extrapolation of the X-ray spectrum, hence suggesting the presence of different emission components in the optical and in the X-ray regime as naturally expected in our model. Monitors in both the X-rays and the TeV emission show evidence for a 23-day periodicity during the 1997 high state. As demonstrated above the 23-day period in the X-ray light curve may be caused by a pulsational instability in two-temperature accretion disk, and via the inverse Compton process result in the same periodicity in the TeV light-curve. A pulsational instability occurring in a two-temperature disk with transition radius $r_0 \sim (48-65)\, r_{g}$ will result in a recurrence timescale of 8 to 100 days. Based on the observed TeV variability we have employed a characteristic black hole mass of $9 \times 10^7 M_{\odot}$ in our calculations. We note that quite different central mass estimates for Mkn~501 have been claimed in the literature, ranging from several times $10^7$ (mainly based on high energy emission properties) up to $10^9\,M_{\odot}$ (based on host galaxy observations), see e.g. \cite{cao02,dej99,fal02,fan05,rie03}. However, as shown by Rieger \& Mannheim~\cite{rie03} uncertainties associated with host galaxy observations may easily lead to an overestimate of the central black hole mass in Mkn~501 by a factor of three and thus reduce the implied central mass to $\simeq (2-3)\times 10^8\,M_{\odot}$, { a value in fact recently confirmed by an independent analysis of central mass constraints derived from host galaxy observations \cite{woo05}.} Moreover, as argued by the same authors some of the apparent disagreement in central mass estimates may possibly be resolved if a binary black hole system exists in the center of Mkn~501, see also \cite{rie00,rie03,vil99}, similar as in the case of OJ~287 \cite{sil88}. For example, if Mkn~501 harbours a binary system with a more massive primary black hole of $\lppr 10^{9} M_{\odot}$ and a less massive (jet-emitting) secondary black hole of $\sim 10^{8}\,M_{\odot}$, the mass ratio $\rho=m/M$ would be of order 0.1, which may compare well with the result $\rho<0.25$ estimated for OJ~287 \cite{liu02}. While our characteristic black hole mass employed falls well within the above noted range, we note that the SLE pulsational instability model may still work successfully, if a higher black hole mass is used. For example, if one adopts $M = 3 \times 10^{8} M_{\odot}$ and $P=23$ days, one obtains $\alpha \leq 0.07$, $r_{0*}=46$ and $\dot{M}\simeq 0.008 \dot{M}_{\rm Edd}$. Our analysis is based on a specific disk model (SLE) which is open to questions, in particular with respect to its possible stability properties. We note however, that a relatively small change in the usually employed viscosity description may already lead to a thermally stable configuration \cite{pir78}. On the other hand, it may as well be possible that the SLE configuration represents a quasi-transient phenomenon associated with those changes in accretion history that probably initiate the high states. An alternative (inner) disk configuration of interest may be represented by an optically thin two-temperature ADAF solution \cite{nar98,yi99}. Such a configuration can exist for accretion rates $\dot{m}= \dot{M}/\dot{M}_{\rm Edd}$ below a critical rate $\dot{m}_{\rm crit} \simeq 0.3\,\alpha^2 \simeq 0.019$, where the canonical ADAF value of $\alpha =0.25$ has been employed \cite{nar98}. ADAFs are generally less luminous than a standard disk, with the typical ADAF luminosity given by $L_A \simeq 0.02\,(\dot{m}/\alpha^2)\,\dot{M}\,c^2$ \cite{yi99}. Using the constraints above, the possible ADAF luminosity for Mkn~501 becomes $L_A \leq 1.3 \times 10^{43}$ erg/s, which is already about an order of magnitude smaller than required by the X-ray analysis. This suggests that -- at least during its high state -- an optically thin ADAF is not a viable option for Mkn~501. Based on Eq. (7) we can estimate the variation rate of the period due to the motion of the instability \begin{eqnarray} \delta P/P = {\frac{3}{2}} P{\frac{1}{r}} {\frac{\partial{r}}{\partial{t}}}\,. \end{eqnarray} Using $v_{r} = {\frac{\partial{r}}{\partial{t}}}$, $\dot{M} = 4 \pi \rho h r v_{r}$ and Eqs. (4) and (5), one finds \begin{eqnarray} \delta P/P =0.44 ~\alpha^{-7/6}~\dot{M_{24}}^{5/6}~M_{7}^{-5/6}~r_{*}^{-1/4} \zeta^{-1/6} \end{eqnarray} For $\alpha$ = 0.28 the relevant parameters result in $\delta P/P= 0.23$. From the period analysis performed by Kranich et al.~\cite{kra99} a $3\sigma$ deviation in period corresponds to 6.67 days. If we take this deviation as the intrinsic variation on the periodicity (P), then a $\delta P/P$ of ${\frac{6.67}{23}}=0.29$ can be estimated from the results by Kranich et al.~\cite{kra99}. As this estimate assumes that the deviation of the period is only affected by the motion of the instability, while it may in fact be caused by more than one effect, our theoretical $\delta P/P$ should not be greater than the observational results. We conclude that the observed 23-day QPOs in Mkn~501 might be caused by the instability of a two-temperature accretion disk. The model presented here may thus offers an alternative explanation to the binary-driven helical jet model of Rieger \& Mannheim~(2000). Comprehensive computational modelling of the pulsational instability in a two-temperature, cooling dominated disk will be essential to verify this in more detail. Our model predicts that a period correlation in the X-ray and $\gamma$-ray should always be present during an active source stage, while the period of the QPOs may vary as the instability region could change from one high state to the other. | 7 | 10 | 0710.1340 |
0710 | 0710.4086_arXiv.txt | { Based on the results of applying the extended ADC emission model to three Z-track sources: GX\th 340+0, GX\th 5-1 and Cyg\th X-2, we propose an explanation of the Z-track sources in which the Normal and Horizontal Branches are dominated by the increasing radiation pressure of the neutron star. The emitted flux becomes several times super-Eddington at the Hard Apex and Horizontal Branch and we suggest that the inner accretion disk is disrupted by this and that part of the accretion flow is diverted vertically. This position on the Z-track is exactly the position where radio emission is detected showing the presence of jets. We thus propose that high radiation pressure is a necessary condition for the launching of jets. We also show that flaring must consist of unstable nuclear burning and that the mass accretion rate per unit emitting area of the neutron star $\dot m$ at the onset of flaring agrees well with the critical theoretical value at which burning becomes unstable. | % \label{sect:intro} The Z-track sources are the brightest group of Galactic low mass X-ray binaries (LMXB) containing a neutron star persistently emitting at the Eddington luminosity or several times this. The sources trace out a Z-shape in hardness-intensity (Hasinger et al. 1989) clearly showing that strong physical changes take place, probably at the inner disk and neutron star, but a convincing explanation of the Z-track phenomenon does not exist. The majority of LMXB are of the Atoll class which show somewhat different shapes in hardness-intensity which are also not understood, and neither is the relation between the two classes making our understanding of LMXB very incomplete. Moreover, it is well-known that the Z-track sources are detected as radio emitters, but in one branch only, the horizontal branch. Not only is radio detected, but striking results from the VLA show the release of a massive radio condensation from the source Sco\th X-1 (Fomalont et al. 2001). Because radio is detected essentially in one branch only, the sources offer the possibility of determining the conditions found in this branch distinguishing it from the other two branches, and so finding the conditions necessary for jet formation. Possible ways of understanding the Z-track sources are by theoretical approaches, timing studies or spectral studies. A theoretical model for the Z-track sources was produced by Psaltis et al. (1995) based on a magnetosphere of the neutron star and the changing properties and geometry of this as the mass accretion rate changed. However, the model assumed that the Comptonized emission observed in the spectra (of all LMXB) originated in a small central region close to the neutron star, and this is inconsistent with our more recent measurements of Comptonizing region size (Church \& Ba\l uci\'nska-Church 2004, below). Extensive timing studies have been made to investigate QPO variations around the Z-track (e.g. van der Klis et al. 1987), but this has not revealed the nature of the Z-track. Previous spectral fitting has applied the Eastern model (Done et al. 2002, Agrawal \& Sreekumar 2003; di Salvo et al. 2002) which assumes the X-ray emission consists of disk blackbody emission plus non-thermal emission from a small central Comptonizing region. However, our work over a period of 10 years with the dipping class of LMXB provides strong evidence that the source of Comptonized emssion, the ADC, is very extended, typically having a radial extent that is 15\% of the accretion disk size, but increasing with source luminosity, and this is inconsistent with the Eastern model. As a result we have proposed the ``extended ADC'' emission model consisting of blackbody emission from the neutron star plus Comptonized emission from an extended ADC (Church Ba\l uci\'nska-Church 1995). Moreover, the pattern of parameter changes obtained by fitting the Eastern model to the Z-track sources is not very easy to interpret and does not immediately suggest a convincing physical explanation. Thus in the present work, we take the approach of applying the extended ADC model for the first time to the Z-track sources and we present the results of applying this model to the sources GX\th 340+0, GX\th 5-1 and Cygnus\th X-2. \begin{figure*}[!ht] % \begin{center} \includegraphics[width=60mm,height=140mm,angle=270]{church_2007_01_fig1a} % \includegraphics[width=60mm,height=80mm,angle=270]{church_2007_01_fig1b} % \caption{Top: Background-subtracted and deadtime-corrected PCA light curve of the 1997 September observation of GX\th 340+0 with 64 s binning. Bottom: the corresponding variation of hardness ratio \hbox{(7.3 -- 18.1 keV)/(4.1 -- 7.3 keV)} with intensity.} \label{} \end{center} \end{figure*} | We show that the radiation pressure of the enitting part of the neutron star is very strong at the hard apex and horizontal branch of the Z-track in three sources, exactly correlating with the parts of the Z-track where radio emission is observed showing the presence of jets, and we suggest that strong radiation pressure is a necessary condition for jet formation. | 7 | 10 | 0710.4086 |
0710 | 0710.1947_arXiv.txt | % {} {We study the variability of the Fe 6.4 KeV emission line from the Class I young stellar object Elias 29 in the $\rho$ Oph cloud.} {We analysed the data from Elias 29 collected by \xmm\ during a nine-day, nearly continuous observation of the $\rho$ Oph star-forming region (the Deep Rho-Oph X-ray Observation, named {{\sc Droxo}}). The data were subdivided into six homogeneous time intervals, and the six resulting spectra were individually analysed} {We detect significant variability in the equivalent width of the Fe 6.4 keV emission line from Elias 29. The 6.4~keV line is absent during the first time interval of observation and appears at its maximum strength during the second time interval (90 ks after Elias 29 undergoes a strong flare). The X-ray thermal emission is unchanged between the two observation segments, while line variability is present at a 99.9\% confidence level. Given the significant line variability in the absence of variations in the X-ray ionising continuum and the weakness of the photoionising continuum from the star's thermal X-ray emission, we suggest that the fluorescence may be induced by collisional ionisation from an (unseen) population of non-thermal electrons. We speculate on the possibility that the electrons are accelerated in a reconnection event of a magnetically confined accretion loop, connecting the young star to its circumstellar disk.} {} | \label{sec:intro} The X-ray emission from young stellar objects (YSOs) at CCD-resolution is usually modelled as thermal emission from a hot plasma in coronal equilibrium, with higher characteristic temperatures than observed in older and less active stars. An interesting deviation from a pure thermal X-ray spectrum is the presence of fluorescent emission from neutral (or weakly ionised) Fe as shown by the presence of the 6.4 keV line. This was first detected by \citet{ikt01} in the X-ray emission of the YSO YLW16A in $\rho$-Oph, during a large flare: in addition to the Fe\,{\sc xxv} complex at 6.7 keV, a 6.4 keV emission line was clearly visibe. Such fluorescence line is produced when energetic X-rays photoionise cold material close to the X-ray source, and it is therefore a useful diagnostic tool of the geometry of the X-ray emitting source and its surroundings. Since 2001, detections of the Fe K fluorescent emission line at 6.4-keV in the spectra of YSOs have been reported by a number of authors. \citet{tfg+05} has identified seven sources with an excess emission at 6.4 keV among 127 observations of YSOs within the COUP observation of Orion; \citet{fms+05} report 6.4 keV fluorescent emission in Elias 29 in $\rho$-Oph both during quiescent and flaring emission, unlike all other reported detection of Fe fluorescent emission in YSOs that were made during intense flaring; \citet{gfm+07} have detected Fe 6.4 keV emission from a low-mass young star in Serpens, during an intense, long-duration flare. Recently, \cite{sc2007} have reported intense Fe fluorescent emission in the spectrum of V 1486 Ori during a strong flare, when the plasma reached a temperature in excess of 10 keV. The 6.4 keV fluorescent line has been detected in different classes of X-ray emitters: X-ray binaries, active galactic nuclei (AGNs), massive stars, supernova remnants, and the Sun itself during flares. In the case of the Sun, the fluorescing material is the solar photosphere, in the YSOs, however, indications are that the material in the circumstellar disk and its related accretion structures could be responsible for the fluorescence. The typical equivalent width of the 6.4 keV emission line in the studies mentioned above is of the order of 150~eV and is too large to be explained with fluorescent emission in the stellar photosphere or in diffuse circumstellar material (e.g. \citealp{tfg+05}; \citealp{fms+05}). This scenario implies that the disk is ``bathed'' in high-energy X-rays emitted by the star, with significant astrophysical implications; for instance, X-rays, in addition to cosmic rays, would play an important role in photoionising the circumstellar material around young star and thus in coupling the gas to the ambient magnetic field (as suggested by e.g. \citealp{gfm00}). \citet{cbt+02} suggest that the ``hot'' component they observe in the disk, in the infrared, is heated by the stellar high-energy emission. \begin{figure*} \begin{center} \leavevmode \epsfig{file=7899fg01.ps, width=17.0cm, bbllx=0,bblly=190,bburx=842,bbury=1000, angle=270, clip=} \caption{The light curve of Elias 29 over the nine days of the \droxo\ observation from PN ({\em top}), MOS1 ({\em bottom-left}), and MOS2 ({\em bottom-right)}. The line with error bars gives the background-subtracted light curve of the source, while the thinner line without error bars gives the total counts (source plus background). The 6 time intervals that we selected for the spectral analysis, on the basis of the PN data, are indicated.} \label{fig:lc} \end{center} \end{figure*} Observations of the time variability of the Fe 6.4 keV emission in YSOs would provide useful constraints on the geometry and sizes of the star-accretion disk system and of other circumstellar structures such as funnel flows, jets, or wind columns. We present here the results of a time-resolved spectral study of the X-ray emission of Elias 29, during $\sim 9$ days of nearly continuous observation by \xmm\ in the context of the ultra-deep observation of $\rho$-Oph, named \droxo\ (from Deep Rho-Oph X-ray observation -- \citealp{psf+07}). We investigated the presence of variations in its strong Fe 6.4 keV emission over the 9-day time scale covered by the observations. This paper is structured as follows. After a summary, below, of the properties of Elias 29, the observations and data analysis are briefly presented in Sect.\,\ref{sec:obs}. Results are summarised in Sect.\,\ref{sec:res}; the simulations carried out to assess the reliability of the line detections and the significance of its variability are described in Sect.\,\ref{sec:sim}. The results are discussed in Sect.\,\ref{sec:disc}. \subsection{Elias 29 (GY214)} Elias 29 (16:27:09.4, $-$24:37:18.9) with a bolometric luminosity $L = 26-27.5~L_{\sun}$ (\citealp{bak+01}; \citealp{nts06}) is the most luminous Class I YSO in the $\rho$-Oph cloud. \cite{mhc98} used the luminosity in the Br$\gamma$ line to determine the object's accretion luminosity at $L_{\rm acc}$ = $15-18 L_{\sun}$, which makes it the source with the highest accretion luminosity in their sample. More recently, \citet{nts06} used the luminosity of the hydrogen recombination lines to derive an accretion luminosity of 28.8~$L_{\sun}$. Using millimeter interferometric observations, \citet{bhc+02} resolved the emission from the disk and the envelope surrounding Elias 29, showing that the disk is in a relatively face-on orientation ($i < 60^{\circ}$), which explains many of the remarkable observational features of this source, such as its flat spectral energy distribution, its brightness in the near-infrared, the extended components found in speckle interferometry observations, and its high-velocity molecular outflow. Their best-fitting disk model has an inner radius of 0.01 AU, outer radius of 500 AU, and a mass $M = 0.012 M_{\sun}$. Elias 29 was previously observed in X-rays with ASCA, \chandra\, and \xmm. In the \chandra\ observation (\citealp{ikt01}), the source quiescent phase is characterised by a temperature of 4.3~keV and luminosity of $2.0 \times 10^{30}$~\es, fully consistent with the values derived from the subsequent \xmm\ observations by \citet{ogm05}: $kT = (3.6-5.1)$ keV, $N({\rm H}) = (4.4-5.3) \times 10^{22}$ cm$^{-2}$, $Z=(0.8-1.3)~Z_{\sun}$, and $L_{\rm X} = 2.8 \times 10^{30}$ \es. The source was seen flaring during one of the ASCA observations (\citealp{kkt+97}) and during the \chandra\ observation. The two flares had similar intensity and duration with an $e$-folding time of $\sim 10$~ks (\citealp{tik+00}; \citealp{ikt01}). | So far, the Fe 6.4 keV emission in YSOs has been explained in terms of fluorescent emission from the photoionised (colder) material in the circumstellar disk. This scenario, however, cannot easily explain the observed variability of the Fe 6.4 keV emission in Elias 29, which occurs in the absence of signficant variations of the observed X-ray continuum. The equivalent width of the line at its maximum strength of $\sim 250$ eV is also not easily reconciled with a photoionisation scenario. An alternative line-formation mechanism is collisional excitation by a population of non-thermal electrons. We suggest that these electrons could be accelerated by magnetic reconnection events in the accretion tubes that connect the star to its circumstellar disk. The electrons are decelerated in situ by the accreting material, ionising it and causing the observed Fe 6.4 keV emission. | 7 | 10 | 0710.1947 |
0710 | 0710.0566_arXiv.txt | We present new observations of the fundamental ro-vibrational CO spectrum of V1647 Ori, the young star whose recent outburst illuminated McNeil's Nebula. Previous spectra, acquired during outburst in 2004 February and July, had shown the CO emission lines to be broad and centrally peaked---similar to the CO spectrum of a typical classical T Tauri star. In this paper, we present CO spectra acquired shortly after the luminosity of the source returned to its pre-outburst level (2006 February) and roughly one year later (2006 December and 2007 February). The spectrum taken in 2006 February revealed blue-shifted CO absorption lines superimposed on the previously observed CO emission lines. The projected velocity, column density, and temperature of this outflowing gas was 30 km s$^{-1}$, $3^{+2}_{-1}\times10^{18}$ cm$^{-2}$, and 700$^{+300}_{-100}$ K, respectively. The absorption lines were not observed in the 2006 December and 2007 February data, and so their strengths must have decreased in the interim by a factor of 9 or more. We discuss three mechanisms that could give rise to this unusual outflow. | McNeil's Nebula was recently illuminated by the outburst of V1647 Ori (McNeil 2004), a young star that is embedded in the Lynds 1630 dark cloud and coincides with the 850$\micron$ continuum source OriBsmm55 (Mitchell et al. 2001). V1647 Ori has a flat SED in the mid-infrared, and is thus a Class I young stellar object (YSO; Andrews et al. 2004). V1647 Ori underwent a similar outburst as recently as 1966 (Aspin et al. 2006), indicating that V1647 Ori is also an EXor. Such pre-main sequence stars undergo eruptive events that dramatically increase their luminosity for periods of months to years (Hartmann 1998). The outbursts are thought to be triggered by a rapid increase in the stellar accretion rate (Hartmann \& Kenyon 1996). The eruption in November 2003 of V1647 Ori lasted two years. During the outburst, the star brightened by a factor of 50 in X-rays (Kastner et al. 2006), a factor of 250 in the red (6 mag in the $R_C$ band; Fedele et al. 2007a; Brice\~{n}o et al. 2004), a factor of 15 in the near-IR($\sim$3 mag in the $J$, $H$, and $K$ bands; Reipurth \& Aspin 2004), and a factor of $\sim$15 at wavelengths from $3.6\micron$ to $70\micron$ (Muzerolle et al. 2005). From the overall brightening of the source, Muzerolle et al. (2005) concluded that the bolometric luminosity increased by a factor of 15 to 44$L_\sun$ (see also Andrews et al. 2004) and that the stellar accretion rate increased from $\sim10^{-7} M_{\sun}$ yr$^{-1}$ to $\sim10^{-5} M_{\sun}$ yr$^{-1}$. Similarly, Gibb et al. (2006) inferred a stellar accretion rate of $3-6\times10^{-6} M_{\sun}$ yr$^{-1}$ from the luminosity of the Br$\gamma$ emission one year later. This is somewhat larger than the typical accretion rate of a young low mass star ($10^{-8}-10^{-7} M_{\sun}$ yr$^{-1}$; Bouvier et al. 2007), yet lower than is expected for a star of the FUor type ($\sim10^{-4} M_{\sun}$ yr$^{-1}$; Hartman \& Kenyon 1996). During the onset of the outburst of V1647 Ori, observations of atomic lines with P Cygni profiles provided evidence for a hot ($T\sim$10,000 K), high velocity ($v = -400$ km s$^{-1}$) wind (Brice\~{n}o et al. 2004; Reipurth et al. 2004; Vacca et al. 2004; Walter et al. 2004; Ojha et al. 2006; Fedele et al. 2007a) with a mass-loss rate of $\dot{M}_{\rm wind}$=4$\times$10$^{-8}$ \Mdot (Vacca et al. 2004). This mass loss rate is much lower than that of the typical FUor (Hartmann \& Kenyon 1996) and comparable to that of a classical T Tauri star (cTTS; Hartigan et al. 1995). The absorption component of the P Cygni profile of several lines (e.g. Pa$\beta$) disappeared within a few months following the peak of the outburst in early 2004 (Gibb et al. 2006). However, P Cygni structure in the H$\alpha$ profile indicated that a weaker wind continued throughout the outburst phase (Ojha et al. 2006; Fedele et al. 2007a). In contrast to the hydrogen and helium lines, the fundamental near-infrared ro-vibrational emission lines of CO, observed 2004 February 27, were broad, centrally peaked, and compatible in their intensity with an excitation temperature of 2500 K (Rettig et al. 2005). The width of the lines was shown to be consistent with Keplerian orbital motion of the gas within the inner disk surrounding the central star, similar to the broad emission line profiles that are observed around cTTSs and Herbig Ae/Be stars (HAeBes; Najita et al. 2003; Blake \& Boogert 2004). A later observation, obtained on 2004 July 30, showed that the CO lines remained broad but the temperature of the gas decreased to 1700 K (Gibb et al. 2006). Neither observation showed any indication of CO in an outflow, as a blue shifted absorption component was not detected. We report followup observations of CO from V1647 Ori, which were acquired 2006 February, 2006 December, and 2007 February. By the time of our initial 2006 observation, V1647 Ori had returned to quiescence and the absorption component in the H$\alpha$ line profile had disappeared (Fedele et al. 2007a and references therein). Presumably the accretion rate had fallen by two orders of magnitude to its pre-outburst accretion rate as the star faded to its pre-outburst brightness (see Muzerolle et al. 2005). As suggested by the work of Najita et al. (2003), only a continued decrease in CO line intensity from the warm gas in the disk was therefore expected. However, the first post-outburst observation revealed the striking metamorphosis of these lines from centrally peaked emission features to emission lines with blue-shifted absorption. Subsequently, by late 2006 and early 2007, the CO emission lines returned to their original centrally peaked structure, indicating that the production of the outflow diminished within one year of the end of the outburst. In this paper we discuss three scenarios that can give rise to such a phenomenon. | During quiescence, V1647 Ori is a class I YSO (Andrews et al. 2004). However, its mass-loss rate during outburst was similar to a strongly accreting cTTS (Vacca et al. 2004). While outflows from cTTSs and HAeBes are common, fundamental ro-vibrational CO emission lines with a blue-shifted absorption component have not been observed around any of the more than 300 such sources observed to date (Najita et al. 2000, 2003; Blake \& Boogert 2004; Rettig et al. 2006; Brittain et al. 2007; J. Brown private communication). Further, Class I YSOs such as GSS 30 IRS 1, HL Tau and RNO 91 do not show ro-vibrational CO emission lines with blue-shifted absorption components (Pontoppidan et al. 2002, Brittain et al. 2005, and Rettig et al. 2006, respectively). In contrast to these systems, the FUor V1057 Cyg has shown blue-shifted overtone ro-vibrational CO absorption lines (Hartmann et al. 2004). While the similarity of the unusual CO outflows is intriguing, there are important differences between FUors such as V1057 Cyg and EXors such as V1647 Ori. First, FUors have greater accretion rates ($\sim$10$^{-4}$ \Mdot) and more extreme winds than EXors (Hartmann \& Kenyon 1996). Secondly, CO is always and {\it only} detected in absorption toward FUors, and is thought to originate in the accretion disk (Calvet et al. 1993; Calvet et al. 1991). Despite the significant differences between V1057 Cyg and V1647 Ori, both stars have undergone major eruptions that have generated outflows in CO, suggesting that a different mass-loss mechanism may be at work in these systems than in other young stars. Given the unusual nature of the outflowing CO from V1647 Ori, it is of interest to consider the mechanisms that could give rise to this phenomenon. One possibility is that the absorbing CO condensed out of the hot outflowing wind that was observed during the outburst. There were two outflow components noted in the H$\alpha$ absorption feature: a variable component at 400 km s$^{-1}$ and a steady component at 150 km s$^{-1}$ (Fedele et al. 2007a). If the lower-velocity gas decelerated at a constant rate to $30$ km s$^{-1}$ (the velocity of the outflowing CO), by the time of our second CO observation in 2004 July the wind would have expanded to 10 AU. There was no evidence of blue-shifted CO absorption on this date (Fig. 1). By 2006 February, when the absorption was observed, the moderate velocity wind would have expanded to nearly 40 AU. However, it seems unlikely that the wind could still retain a kinetic temperature of 700K at a distance of 40AU from the star. Furthermore, it is not clear why CO would condense out of this outflow to reveal warm, blue shifted absorption but not out of any of the other outflows with similar or even greater mass-loss rates. Thus we conclude that this scenario is unlikely. A second possibility is that the CO absorption formed in a shell of material that was swept up by the atomic wind. In this case the absorption did not appear until the column density of material was sufficient to produce measurable absorption, and the heating was the result of the interaction of the wind with the shell. When the mass loss decreased at the end of the outburst phase, this heating was eventually shut down. If the CO/H$_2$ ratio in such a shell was similar to that of a dense molecular cloud, 1.5 $\times$10$^{-4}$, then the column density of gas was 2$\times$10$^{22}$ cm$^{-2}$. Adopting a normal interstellar extinction--to--gas ratio, i.e., $A_V=5.6\times 10^{22} N_{\rm H}$ mag cm$^2$ atom$^{-1}$ (Bohlin et al. 1978), we find that the observed column density of CO corresponded to $\sim$20 mags of visible extinction. This is a lower limit, as it is possible that selective dissociation of gas in the nebula could drive down the relative abundance of CO. However, the extinction measured on the line of sight toward V1647 Ori appears to have remained relatively unchanged over the entire course of the outburst at $\sim$11 mags (Vacca et al. 2004; Gibb et al. 2006), of which 6.5 mags is due to the nebula (Fedele et al. 2007b). There is no evidence for an additional 10--20 mags of extinction toward V1647 Ori, and so we conclude that this scenario is also unlikely. A final scenario we consider is that the CO outflow was launched in response to the reorganization of the stellar magnetic field following the sharp drop in the accretion rate. Pre-main sequence stars tend to have kilogauss magnetic fields which mediate stellar accretion and outflows (e.g. Johns-Krull et al. 2000). This stellar magnetic field truncates the accretion disk where the ram pressure from accretion balances the magnetic pressure from the magnetosphere (Camenzind 1990). Consequently, the truncation radius of the disk, $R_T$, is inversely and nonlinearly proportional to the accretion rate, $\dot{M}$, and given by $ R_T \propto \dot{M}^{-2/7}$ (e.g. Bouvier et al. 2007). Thus the two order of magnitude change in the stellar accretion rate experienced by V1647 Ori would have resulted in the truncation radius being shifted by nearly a factor of four. Two years into the outburst, V1647 Ori rapidly faded by a factor of 40 in the R$_C$-band in just 180 days to return to its pre-outburst level (Fedele et al. 2007a). The mid-infrared flux also returned to its pre-outburst level, indicating that the drop in the optical/near-infrared lightcurve was due to the intrinsic fading of the source (Fedele et al. 2007a). This sharp drop in the luminosity of the source indicates that the accretion rate fell rapidly as the star returned to quiescence. In response to the drop in ram-pressure from the accretion flow, the truncation radius was pushed back by the magnetosphere. Evidence from simulations suggests that disk-magnetosphere systems tend to form outflows when the system is undergoing the greatest amount of dynamical rearrangement (e.g., Fig. 7 in Balsara 2004). It is possible that the realignment of the magnetic field as it pushed out against the circumstellar disk resulted in a warm, shortlived outflow, an outflow that is not observed toward any other cTTSs or HAeBes. While virtually all accreting low-mass stars drive outflows, the CO outflow from V1647 Ori is highly unusual. Indeed, the transformation of centrally peaked ro-vibrational CO emission lines to CO emission lines with blue-shifted absorption is unique. We suggest that the mechanism responsible for producing this outflow is distinct from the one that drives the outflow from typical cTTSs. The coincidence between the rapid fading of V1647 Ori and the subsequent observation of the CO outflow hints at a connection. Better sampling of the fundamental ro-vibrational CO spectrum of EXors as they brighten and fade is crucial for determining whether this coincidence is significant. The rapid and dramatic change in the accretion rate that characterizes the EXor phenomenon provides an important opportunity to study the interplay between stellar accretion, the inner disk, and outflows. This insight is key to reaching a satisfactory theoretical understanding of these events. | 7 | 10 | 0710.0566 |
0710 | 0710.2563_arXiv.txt | In this study we compile for the first time comprehensive data sets of solar and stellar flare parameters, including flare peak temperatures $T_p$, flare peak volume emission measures $EM_p$, and flare durations $\tau_f$ from both solar and stellar data, as well as flare length scales $L$ from solar data. Key results are that both the solar and stellar data are consistent with a common scaling law of $EM_p \propto T_p^{4.7}$, but the stellar flares exhibit $\approx 250$ times higher emission measures (at the same flare peak temperature). For solar flares we observe also systematic trends for the flare length scale $L(T_p) \propto T_p^{0.9}$ and the flare duration $\tau_F(T_p) \propto T_p^{0.9}$ as a function of the flare peak temperature. Using the theoretical RTV scaling law and the fractal volume scaling observed for solar flares, i.e., $V(L) \propto L^{2.4}$, we predict a scaling law of $EM_p \propto T_p^{4.3}$, which is consistent with observations, and a scaling law for electron densities in flare loops, $n_p \propto T_p^2/L \propto T_p^{1.1}$. The predicted ranges of electron densities are $n_p \approx 10^{9-10}$ cm$^{-3}$ for solar nanoflares at $T_p=1$ MK, $n_p \approx 10^{10-11}$ cm$^{-3}$ for typical solar flares at $T_p=10$ MK, and $n_p \approx 10^{11-12}$ cm$^{-3}$ for large stellar flares at $T_p=100$ MK. The RTV-predicted electron densities were also found to be consistent with densities inferred from total emission measures, $n_p=\sqrt{EM_p/q_V V}$, using volume filling factors of $q_V=0.03-0.08$ constrained by fractal dimensions measured in solar flares. Solar and stellar flares are expected to have similar electron densities for equal flare peak temperatures $T_p$, but the higher emission measures of detected stellar flares most likely represents a selection bias of larger flare volumes and higher volume filling factors, due to low detector sensitivity at higher temperatures. Our results affect also the determination of radiative and conductive cooling times, thermal energies, and frequency distributions of solar and stellar flare energies. | Scaling laws provide important diagnostics and predictions for specific physical models of nonlinear processes such as self-organized criticality, turbulence, diffusion, plasma heating, and particle accleration. These models have been widely applied in plasma physics, astrophysics, geophysics, and the biological sciences. Here we investigate scaling laws of physical parameters in solar and stellar flares, which should allow us to decide whether solar and stellar flare data are consistent with the same physical flare process. The scaling of solar and stellar flare data has been pioneered by Stern (1992), Feldman et al.~(1995b), and Shibata \& Yokoyama (1999; 2002), who showed evidence for a nonlinear scaling between the flare volume emission measure $EM_p$ and the flare peak temperature $T_p$. These parameters have been measured in solar flares with instruments like {\sl Skylab}, {\sl GOES}, {\sl Yohkoh/SXT} (Soft X-ray Telescope), and {\sl RHESSI (Ramaty High Energy Solar Spectroscopic Imager)}, and in stellar flares with {\sl ASCA}, {\sl BeppoSAX}, {\sl Einstein}, {\sl EUVE}, {\sl EXOSAT}, {\sl Ginga}, {\sl HEAO}, {\sl ROSAT}, {\sl Chandra}, and {\sl XMM-Newton}. Compilations of solar flare parameters have been presented in Aschwanden (1999), while stellar flare parameters were compiled in a recent review by G\"udel (2004). In this paper we present for the first time this host of mostly new measurements ``on the same page'' and investigate commonalities and differences between the scaling of solar and stellar flares. In Section 2 we present the statistical correlations found in stellar flare data, while the corresponding counterparts of solar flare data are shown in Section 3. In Section 4 we present theoretical modeling of the data, using the well-known RTV law, the generalization with gravitational stratification and spatially non-uniform heating, the fractal flare volume scaling, and volume filling factor. In Section 5 we discuss the differences between solar and stellar scaling laws, the consistency between two different electron density measurement methods, and a previously derived ``universal scaling law'' for solar and stellar flares. Section 6 summarizes our conclusions. | We compiled directly observed parameters from solar and stellar flares, such as the volume peak emission measure $EM_p$, flare peak temperature $T_p$, flare duration $\tau_f$, and flare length scale $L$ (the latter only for solar flares). A prominent statistical correlation is found between the volume emission measure $EM_p$ and flare peak temperature $T_p$, which scales as $EM_p \approx T_p^{4.7}$ for both solar and stellar flares. Another recent study demonstrated that the flare volume has a fractal scaling, $V(L)\propto L^{2.4}$, rather than the generally used Euclidian scaling of $V(L)\propto L^3$. Applying the RTV scaling law, combined with the fractal volume scaling and the statistical $L-T_p$ correlation $L(T_p) \propto T_p^{0.9}$, leads directly to a theoretically predicted scaling law of $EM_p \propto T_p^{4.3}$, which explains the observed correlations in both solar and stellar flares. A second result we find is an unexplained offset by a factor of about 250 between solar and stellar flares at the same temperature, which is likely due to a selection bias for stellar flare events with larger volume filling factors and larger spatial scales. Interestingly, however, this selection bias does not affect the overall $EM-T$ relationship and the lower threshold has a similar functional dependence of $EM_{min} \propto T_p^{4.7}$, probably because the detector sensitivities are dropping off with a similar function with higher temperatures. A third result is that our model of fractal flare volume scaling provides realistic estimates of volume filling factors, and thus of flare densities. We find that the electron densities in solar flare loops can be predicted based on our fractal scaling in close agreement to the predictions of the RTV law. The agreement of the predicted electron densities with both methods agrees always better than an order of magnitude (Fig.~7), although the absolute magnitude varies by three orders of magnitude between the smallest nanoflares and the largest solar flares, i.e., $n_e \approx 10^9-10^{12}$ cm$^{-3}$. Since the RTV scaling law (combined with the observed $L \propto T_p$ correlation) predicts about a linear relationship between the electron densities and flare peak temperatures, i.e., $n_p \propto T_p$, we expect up to an order of magnitude higher electron densities in the largest stellar flares due to the higher temperature than in solar flares. The determination of correct scaling laws allows us also to infer realistic estimates of the (conductive and radiative) flare cooling times, which can be tested from the e-folding decay time of individual peaks in (solar and stellar) flare light curves. The scaling laws allow us also to eliminate temperature biases in the statistics of the total thermal energy of flares, i.e., $E_T \propto n_p T_p V \propto EM_p T_p/n_p$. Since we find that the electron density (corrected for a fractal filling factor) scales approximately as $n_p \propto T_p^{1.1}$, the thermal energy scales approximately as $E_T \propto EM_p$, and thus the observed total emission measure $EM_p$ can be used as a good proxy for the thermal flare energy $E_T$. Such unbiased frequency distributions of flare energies $N(E_T)$ permit us then to determine whether there is more energy in large or small flares, an important test for nanoflare heating theories. | 7 | 10 | 0710.2563 |
0710 | 0710.0799_arXiv.txt | { Non thermal emission from galaxy clusters demonstrates the existence of relativistic particles and magnetic fields in the Intra Cluster Medium (ICM). Present instruments do not allow to firmly establish the energy associated to these components. In a few years gamma ray observations will put important constraints on the energy content of non thermal hadrons in clusters, while the combination of radio and hard X-ray data will be crucial to measure the energy content in the form of relativistic electrons and magnetic field. SIMBOL-X is expected to drive an important breakthrough in the field also because it is expected to operate in combination with the forthcoming low frequency radio telescopes (LOFAR, LWA). In this contribution we report first estimates of {\it statistical properties} of the hard X--ray emission in the framework of the {\it re-acceleration model}. This model allows to reproduce present radio data for Radio Halos and to derive expectations for future low frequency radio observations, and thus our calculations provide hints for observational strategies for future radio and hard--X-ray combined observations. | Clusters of galaxies represent the largest virialized structures in the present Universe. Rich clusters have typical total masses of $10^{15} M_{\odot}$, mostly in the form of dark matter, while $\sim 5\%$ of the mass is in the form of a hot ($T \sim 10^8 K$), tenuous ($n_{gas} \sim 10^{-3}-10^{-4} cm^{-3}$), X-ray emitting gas. In terms of energy density, the gas is typically heated to roughly the virial temperature, but there is also room to accomodate a non-negligible amount of non-thermal energy. Clusters are ideal astrophysical environments for particle acceleration and cosmic rays (CR) accelerated within the cluster volume are expected to be confined for cosmological times (e.g., Blasi, Gabici, Brunetti 2007, BGB07, for a review). The bulk of the energy of these CRs is expected in protons since they have radiative and collisional life--times much longer than those of the electrons. While present gamma ray observations can only provide upper limits to the average energy density of CR protons in the ICM (e.g. Reimer et al. 2004), evidence of a non-thermal component is in fact obtained from radio observations of a fraction of galaxy clusters showing synchrotron emission on Mpc scales : Radio Halos, fairly symmetric sources at the cluster center, and Radio Relics, elongated sources at the cluster periphery (e.g., Feretti 2005). Although the bulk of present data comes from radio observations, theoretically a substantial fraction of the non thermal radiation is expected from inverse Compton (IC) scattering of the photons of the cosmic microwave background (e.g., Sarazin 1999). Measuring IC emission from clusters in the hard X--rays is extremely important to derive the energy density of emitting electrons and the strength of the magnetic field when these measures are combined with radio data. Despite the poor sensitivity of present and past hard X--ray telescopes, several groups have claimed detection of hard X--ray emission (HXR) in a few massive clusters (e.g., Fusco-Femiano et al.~2004; Petrosian et al.~2006; Rephaeli et al.~2006; see also Rossetti \& Molendi 2004 and Fusco-Femiano et al.~2007 for a discussion on the strength of the HXR detection in the Coma cluster). Thanks to its sensitivity and capability to perform hard X--ray imaging SIMBOL--X will open a new era in the study of non thermal radiation from galaxy clusters. In this contribution we report first expectations on the Luminosity Functions (LFs) and number counts of HXR from clusters. We calculate only the contribution to the IC spectrum from electrons re-accelerated by turbulence in the ICM which are the responsible for the origin of Radio Halos in the context of the {\it re-acceleration scenario}. | In this contribution we report first expectations for HXRs from galaxy clusters, a more detailed study will be reported in a forthcoming paper. Calculations are performed in the framework of the {\it re-acceleration model} assuming physical parameters which allow the {\it re-acceleration model} to reproduce present data of the statistical behaviour of giant Radio Halos. The strength of the magnetic field in the ICM is a crucial parameter in our calculations and we have shown that SIMBOL-X will provide unique constraints. By assuming a value of the magnetic field averaged in Mpc$^3$ volume of $\approx$0.2$\mu$G we find that SIMBOL-X will discover HXRs in $\approx$30--100 clusters at z$\leq$0.2 . | 7 | 10 | 0710.0799 |
0710 | 0710.0367_arXiv.txt | { Combined X-ray synchrotron and inverse-Compton $\gamma$-ray observations of pulsar wind nebulae (PWN) may help to elucidate the processes of acceleration and energy loss in these systems. In particular, such observations provide constraints on the particle injection history and the magnetic field strength in these objects. The newly discovered TeV $\gamma$-ray source HESS\,J1718$-$385 has been proposed as the likely PWN of the high spin-down luminosity pulsar PSR\,J1718$-$3825. The absence of previous sensitive X-ray measurements of this pulsar, and the unusual energy spectrum of the TeV source, motivated observations of this region with \emph{XMM-Newton}. The data obtained reveal a hard spectrum X-ray source at the position of PSR\,1718$-$3825 and evidence for diffuse emission in the vicinity of the pulsar. We derive limits on the keV emission from the centroid of HESS\,J1718$-$385 and discuss the implications of these findings for the PWN nature of this object. } | Young pulsars drive relativistic winds into their environments, confinement of which leads to the production of extremely broadband emission via the synchrotron and inverse-Compton (IC) processes \citep[see][for a recent review]{PWN:review}. The most prominent PWN, the Crab Nebula, is detected in all wavebands from the radio to TeV $\gamma$-rays~\cite{Whipple:crab}, with the transition from synchrotron to IC emission at $\sim$1~GeV. The recent increase in sensitivity of ground-based TeV $\gamma$-ray instruments has led to a rapid increase in the number of putative PWN in this waveband. These objects are characterised by diffuse, typically offset, nebulae around high spin-down luminosity pulsars. The archetype of this new object class is the PWN G\,18.0---0.7/HESS\,J1825$-$137. G\,18.0---0.7 is a $\sim$5$'$ long asymmetric X-ray synchrotron nebula associated with the middle-aged (characteristic spin-down age $\tau$=21~kyr) pulsar PSR\,B1823$-$13~\cite{XMM:1825}. The IC nebula HESS\,J1825$-$137 is much larger ($\sim$100$'$ at 1~TeV) but exhibits energy-dependent morphology, shrinking towards the pulsar at high energies~\cite{HESS:1825p2,HESS:1825icrc}, suggestive of cooling of the highest-energy (X-ray synchrotron emitting) electrons away from the pulsar. The TeV $\gamma$-ray source HESS\,J1718$-$385 was discovered in deep observations of the supernova remnant RX\,J1713.7$-$3946 using H.E.S.S. in 2004-2005~\cite{HESS:twopwn}. The absence of other potential counterparts and the relatively compact nature of the source ($9'\times4'$ rms) make an association with PSR\,J1718$-$3825 ($8'$ from the centroid of the TeV source) plausible. The TeV source is unusual in its sharply peaked spectral energy distribution (SED), which is similar to that of the $\gamma$-ray nebula of the Vela pulsar~\cite{HESS:velax}. The $\gamma$-ray emission from these objects is commonly attributed to IC scattering of relativistic electrons \citep[see][for an alternative view]{Horns:VelaX}. In this scenario the spectral break seen at $\sim$10~TeV in these objects can be interpreted as a signature of electron cooling. However, PSR\,J1718$-$3825 (estimated distance 4.2~pc) has a characteristic spin-down age (90~kyr) almost an order of magnitude greater than that of the Vela pulsar, making such a high energy break very surprising. The search for a possible X-ray counterpart to HESS\,J1718$-$385 is important for two reasons: firstly, to verify the identification of the TeV source as the PWN of PSR\,J1718$-$3825 and secondly, to explore the physical conditions and electron energy distribution in the putative nebula. As no sensitive X-ray observations of the PSR\,J1718$-$3825/HESS\,J1718$-$385 region existed, \emph{XMM-Newton} was used to observe this region in September 2006. | The discovery of hard spectrum X-ray emission from the vicinity of PSR\,J1718$-$3825, and the evidence for a diffuse halo around the pulsar strongly suggest the existence of a synchrotron nebula around this pulsar. This discovery strengthens the association of the $\gamma$-ray source HESS\,J1718$-$385 to PSR\,J1718$-$3825, but the relationship of the X-ray emission to the $\gamma$-ray source is not straightforward. The overall asymmetry of the nebula with respect to the pulsar is consistent with the idea of SNR expansion into a non-uniform molecular environment \citep[see for example][]{blondin01:PWN}. The very different morphologies in the two wavebands suggest that either electrons of rather different energies are responsible for the two sources and/or that the magnetic field strength within the nebula is highly non-uniform. As the target for IC emission is the CMBR and other large-scale radiation fields, the IC flux $F_{\rm IC}$ is simply proportional to the number of radiating electrons, $n_{e}$, whereas the synchrotron flux goes as: $F_{\rm synch}\propto B^{2}n_{e}$. In either case the situation may be rather similar to that of HESS\,J1825$-$137 or indeed HESS\,J1813$-$178 \cite{Funk:1813}, with the lifetime of TeV $\gamma$-ray emitting electrons being longer than the age of the pulsar, and having time to propagate over distances of several parsecs. The SED of the source is presented in Figure~\ref{fig:sed}. Three features are of note: \begin{figure}[t] \centering \resizebox{0.9\hsize}{!}{\includegraphics{Figure3}} \vspace{-1mm} \caption{Spectral energy distribution for the pulsar wind nebula of PSR\,J1718$-$3825. The de-absorbed spectrum of emission from within $1'$ of the pulsar is shown together with limits for diffuse emission from the region covered by HESS\,J1718$-$385. Three sets of illustrative synchrotron and inverse Compton model curves are shown, based on assumptions of: {\bf A)} Mono-energetic 70~TeV electrons injected over a $10^{4}$ year period, $B = 5\mu$G, {\bf B)} 70 TeV electrons, $B=20\,\mu$G, $t=8$ years and {\bf C)} An electron energy distribution following a power-law ($\alpha=1.8$) with an abrupt cut-off at 100~TeV. These curves are calculated as described in \citet{GC:Hinton07}. } \label{fig:sed} \end{figure} 1) a hard IC spectrum at TeV energies, with a peak at $\sim$10~TeV. This suggests that the electrons responsible for this emission are uncooled. This is rather surprising given the spin-down age of the pulsar (90 kyr): cooling on the CMBR alone would result in a spectral break at 2~TeV after 90~kyr, inconsistent with the $\gamma$-ray data. Indeed ages $>40$ kyrs appear to be excluded by the data. A true age of $\sim$10~kyr could be explained by a birth period for the pulsar very close to its current period of 75~ms, or breaking deviating significantly from the pure magnetic dipole case (as it seems may commonly be the case, see \citet{Kramer:Spindown}) \footnote{We further note that the projected length of the $\gamma$-ray nebula ($\sim$10~pc) is roughly half that of HESS\,J1825$-$137, despite the fact that PSR\,J1718$-$3825 is apparently a factor 5 older. Another possibility is that HESS\,J1825$-$137 represents only the youngest part of a larger, softer spectrum, $\gamma$-ray nebula.}. The shape of the TeV spectrum appears to be consistent with a constant injection of (mono-energetic) $\sim$70~TeV electrons (curve {\bf A}), or with a hard power-law (index $\approx$1.8) with a sharp cut-off around 100 TeV ({\bf C}). 2) hard spectrum X-ray emission from the pulsar vicinity with a much lower energy flux. The spectrum from within $1'$ of PSR\,J1718$-$3825 shown in Figure~\ref{fig:sed} represents the combined flux of pulsar itself and the inner PWN. The pulsar contribution is certainly less than half of the total emission (as is clear from the flux in the annulus surrounding the pulsar) and for typical systems of this type represents only $\sim$20\% of the PWN emission in the $>2$keV range \citep{Kargaltsev:ChandraPWN}. Assuming the PWN is dominant, the hard spectrum suggests either higher electron energies or larger magnetic fields in this region ({\bf B}) in comparison to those found in the $\gamma$-ray nebula. The lower flux can be explained if only recently injected electrons are confined in the region around the pulsar. Whilst the two-zone scenario illustrated by curves {\bf A} and {\bf B} is clearly grossly oversimplified, is does appear that the data are consistent with the idea that electrons with a relatively narrow energy distribution rapidly escape from a high $B$-field region close to the pulsar into the extended nebula seen in $\gamma$-rays. Another plausible scenario is that the injection spectrum of electrons has changed significantly over the lifetime of the pulsar. 3) the energy flux level of diffuse X-ray emission from the HESS\,J1718$-$385 region exceeds the TeV flux by not more than a factor $\sim$2. The energy distribution of electrons is essentially fixed by the TeV data. The diffuse X-ray limits can therefore be used constrain the $B$-field in the extended nebula to be not much greater than $5\,\mu$G (curves {\bf C}+{\bf A}), close to the mean value of the ISM. We note that this constraint comes principally from the diffuse limit at 2--4.5 keV which is relatively independent of the assumed absorbing column. In conclusion, the diffuse X-ray emission around XMMU\,171813.8$-$382517 and HESS\,J1718$-$385 appear to represent different zones in the PWN of the middle-aged pulsar PSR\,J1718$-$3825. Future studies of this complex system are certainly well motivated. For example, with the superior angular resolution of Chandra, the contribution of the pulsar itself to the non-thermal emission could be separated from that of the PWN. | 7 | 10 | 0710.0367 |
0710 | 0710.5774_arXiv.txt | We review the properties of carbon-sequence ([WC]) Wolf-Rayet central stars of planetary nebulae (CSPNe). Differences between the subtype distribution of [WC] stars and their massive WC cousins are discussed. We conclude that [WO]-type differ from early-type [WC] stars as a result of weaker stellar winds due to high surface gravities, and that late- and early-type [WC] and [WO] stars generally span a similar range in abundances, X(He)$\sim$X(C)$\gg$X(O), consistent with a late thermal pulse, and likely progenitors to PG1159 stars. | % This review discusses the properties of the small fraction of central stars of Planetary Nebulae (CSPNe) which share a spectroscopic appearance with massive, carbon-sequence (WC-type) Wolf-Rayet stars. Massive WC stars are the chemically evolved descendents of initially very massive O stars ($M_{\rm init} \geq 25 M_{\odot}$) exhibiting the C and O products of core helium burning, plus a unique emission line spectral appearance due to fast, dense stellar wind outflows. Such stars are young, with ages of only a few Myr of which several hundred cases are known within the Milky Way, supplemented by thousands more known in external star-forming galaxies (Crowther 2007; Hamann these proc.). CSPNe possessing a similar spectral morphology to WC stars are denoted [WC] and are at a post-Asymptotic Giant Branch (AGB) phase in the late stages of evolution of low or intermediate mass stars ($M_{\rm init} \sim 1-5 M_{\odot}$?) with only a few dozen examples known in the Milky Way plus a handful in the Magellanic Clouds. With respect to normal H-rich CSPNe, the unusual surface chemical composition of [WC] stars apparently results from a late thermal pulse (LTP), causing a H-deficient surface, and likely connection with other H-deficient stars, most notably PG1159 stars (Werner et al. these proc.) | We discuss physical and wind properties of Wolf-Rayet CSPNe drawn from the recent literature plus new analyses for a range of subtypes. In general the higher ionization lines seen at earlier spectral type indicates an increased stellar temperature, although very early [WO] subtypes are favoured for hot CSPNe with weak winds, with [WC4--5] subtypes resulting for hot CSPNe with stronger winds. [WC8--9] subtypes correspond to lower temperature CSPNe with strong winds, with much lower temperatures indicated in [WC10] stars. We explain the differences in subtype distribution between Galactic disk CSPNe and massive WC stars as a result of decreased wind densities in the former, owing to increased surface gravities. Massive WC stars in the LMC differ from Galactic WC stars through reduced wind densities caused by lower metallicities, rather than increased surface gravities. There does {\it not} appear to be a systematic difference between carbon-to-helium mass fractions in late to early-type CSPNe, at least for subtypes for which we are able to employ common diagnostics ([WC9] to [WO1]). Consequently, evolution from late-type [WC] through early-type [WC] and [WO] to PG1159 stars appears to be consistent with most abundance patterns, and expectations for a late thermal pulse. Indeed, we note that similar analysis tools have been applied to the post He-flash system V605 Aql, which has now evolved through to a early-type [WC] spectral type over the past 80 years and shares a similar abundance pattern with X(He):X(C):X(O) = 54:40:5\% (Clayton et al. 2006). | 7 | 10 | 0710.5774 |
0710 | 0710.5542_arXiv.txt | There are periodic solutions to the equal-mass three-body (and $N$-body) problem in Newtonian gravity. The figure-eight solution is one of them. In this paper, we discuss its solution in the first and second post-Newtonian approximations to General Relativity. To do so we derive the canonical equations of motion in the ADM gauge from the three-body Hamiltonian. We then integrate those equations numerically, showing that quantities such as the energy, linear and angular momenta are conserved down to numerical error. We also study the scaling of the initial parameters with the physical size of the triple system. In this way we can assess when general relativistic results are important and we determine that this occur for distances of the order of $100M$, with $M$ the total mass of the system. For distances much closer than those, presumably the system would completely collapse due to gravitational radiation. This sets up a natural cut-off to Newtonian $N$-body simulations. The method can also be used to dynamically provide initial parameters for subsequent full nonlinear numerical simulations. | The closest star to the solar system, Alpha Centauri, is a triple system, so is Polaris and HD 188753. Triple stars and black holes are common in globular clusters~\cite{Gultekin:2003xd, Miller:2002pg}, and galactic disks. Triple black hole mergers can be formed in galaxy merger~\cite{Valtonen96} and a triple quasar, representing a triple supermassive black hole system has been recently discovered~\cite{Djorgovski:2007ka}. Full numerical simulations of black holes made possible only in the last couple of years have already produced numerous astrophysically interesting results, among them, the orbital hangup and respect of the cosmic censorship hypothesis for spinning black holes~\cite{Campanelli:2006uy,Campanelli:2006fg,Campanelli:2006fy}, precession and spin-flips~\cite{Campanelli:2006fy}, and the discovery~\cite{Campanelli:2007ew} of large recoil velocities in highly-spinning black hole mergers up to $4,000$ km/s~\cite{Campanelli:2007cga}. The 2005 breakthroughs in Numerical Relativity~\cite{Pretorius:2005gq,Campanelli:2005dd,Baker:2005vv}, not only provided a solution to the long standing two-body problem in General Relativity, but it also proved applicable to the black hole - neutron star binaries~\cite{Faber:2007dv} and recently to the three (and $N$) - black holes systems~\cite{Campanelli:2007ea}. In general, the solution of three-body problem in Newtonian gravity can be chaotic. There are however, periodic orbits in the problem of three equal masses on a plane. One of the most surprising solution is a figure-eight % orbit. The three bodies chase each other forever around a fixed eight-shaped curve. This was found first by Moore~\cite{Moore:1993} and discussed with the proof of the existence in Ref.~\cite{Chenciner:2000}. Heggie~\cite{Heggie:2000} also estimates the probability for such systems to occur in a galaxy. Because of effects such as the perihelion shift, it was unclear if the figure-eight orbits would exist in a low post-Newtonian expansion, even if it consist of only conservative terms. Imai, Chiba and Asada succeeded in obtaining the figure-eight solution in a first post-Newtonian order approximation by finding the general relativistic corrections to the Newtonian initial conditions. In Ref.~\cite{Chiba:2006ad} they also estimated the periodic gravitational waves from this system. In Ref.~\cite{Imai:2007gn} was used the Euler-Lagrange equations of motion in an approximation to first post-Newtonian order. In our paper we instead assume the Hamiltonian formulation to derive the equations of motion. We start from the Hamiltonian given in Ref.~\cite{Schaefer87} (with typos corrected in our Appendix). We derive the equations of motion in this formalism, which are different from those used in Ref.~\cite{Imai:2007gn} and have the virtue of explicitly satisfying the constants of motion of the problem, and thus being more amenable to numerical integration. The paper is organized as follows. In Section~\ref{sec:EOM}, we summarize the equations of motion to be solved numerically in order to obtain the figure-eight orbits. The starting point is the three-body Hamiltonian in the first post-Newtonian approximation. In Section~\ref{sec:1PN}, we discuss the initial conditions for the figure-eight solutions. We study the scaling relation between the orbital radius and the linear momenta. From this analysis, we can estimate when general relativistic effects are important. In Section~\ref{sec:2PN}, we extend our calculation to the second post-Newtonian order and in Section~\ref{sec:DIS}, we summarize the results of this paper and discuss some remaining problems. The 2PN three-body Hamiltonian is explicitly given in the Appendix. Throughout this paper, we use units in which $c=G=1$. | \label{sec:DIS} In this paper we have used the figure-eight orbits as a theoretical lab to test the properties of the low post-Newtonian expansions of General Relativity. We have found that those closed orbits exists for three (and presumably $N$) bodies. We have provided an improved first-post-Newtonian order formalism for deriving the equations of motion that satisfy the Hamiltonian (the linear and angular momenta) constraint to round-off error. The subsequent numerical evolution is well behaved during for more than $t\sim10,000m$. We have also extended this analysis to the $2PN$ corrections, still giving a conservative system of equations. In the process of finding the figure-eight solutions by trial of different initial momenta we also showed (numerically) the stability of the orbit against small perturbations. This method is particularly useful to determine, dynamically (as an alternative to determine them through families of initial data~\cite{Campanelli:2005kr}), initial orbital parameters for subsequent full numerical evolution~\cite{Campanelli:2007ea}, when the holes are close enough that general relativistic effects can no longer be ignored. Note that our method fully takes into account the three-body post-Newtonian interactions unlike other simulations that approximate the problem in successive two-body problems~\cite{Aarseth:2007wv}. It is interesting to note here that the scaling fits~(\ref{eq:2PNfit}) give a practical way to determine when relativistic or Newtonian approaches are appropriate. For $\lambda=1$ we have that the ratio of the first coefficient, $0.01617654493$ (Newtonian) to the second coefficient $0.002017242451$ first-post-Newtonian is nearly $0.12/\lambda$ and the second coefficient to the third one $0.0002463605227$ (dominated by second-post-Newtonian) is also approximately $0.12/\lambda$. This indicates that post-Newtonian corrections are important. For $\lambda=1$ the distance between the initial bodies is $200m$, what indicates that for nearly $67M$ with $M\approx3m$ the total mass of the system has strong post-Newtonian effects. For $\lambda \gg 1$ Newtonian gravity should describe the system accurately, while for $\lambda<1$ general relativistic effects should be very important, eventually leading to the total collapse of the system. It is interesting to remark here that most of the $N$-body codes use some sort of regularization of the Newtonian gravity for very close encounters \cite{aarseth-03}, instead the natural way to regularize these close encounters \cite{Campanelli:2007ea} is given by the General Theory of Relativity, and as we show here, the post-Newtonian corrections are already non-negligible at separations of the order of $100M$. In any case, for most of the astrophysical encounters this is way too short distance, but it can obviously be reached in systems involving black holes and neutron stars. | 7 | 10 | 0710.5542 |
0710 | 0710.5632_arXiv.txt | {Homogeneous anisotropic turbulence simulations are used to determine off-diagonal components of the Reynolds stress tensor and its parameterization in terms of turbulent viscosity and $\Lambda$-effect. The turbulence is forced in an anisotropic fashion by enhancing the strength of the forcing in the vertical direction. The Coriolis force is included with a rotation axis inclined relative to the vertical direction. The system studied here is significantly simpler than that of turbulent stratified convection which has often been used to study Reynolds stresses. Certain puzzling features of the results for convection, such as sign changes or highly concentrated latitude distributions, are not present in the simpler system considered here.} | The Reynolds stress, described by the correlation of fluctuating velocity components, $Q_{ij} = \overline{u_i u_j}$, is one of the most important generators of differential rotation in stars (R\"udiger \cite{Ruediger1989}). These stresses have been studied with the help of 3D convection simulations (e.g.\ Pulkkinen et al.\ \cite{Pulkkinenea1993}; Chan \cite{Chan2001}; K\"apyl\"a et al.\ \cite{Kaepylaeea2004}; R\"udiger et al.\ \cite{Ruedigerea2005}). These results have revealed some surprising features such as the peaking of the horizontal stress $Q_{xy}$ very close to the equator, and a positive (outward) flux for rapid rotation. Both of these results are at odds with theoretical considerations (Kitchatinov \& R\"udiger \cite{KitcRued1993}). Furthermore, disentangling of the diffusive (turbulent viscosity) and non-diffusive ($\Lambda$-effect) parts of the stress is difficult from convection simulations. Here, we present preliminary results from anisotropic homogeneous, isothermal, non-stratified turbulence simulations in which diffusive and non-diffusive effects can be studied separately. Imposing a linear shear flow on top of isotropically driven turbulence allows the study of turbulent viscosity without $\Lambda$-effect. On the other hand, using a special form of forcing, anisotropic homogeneous turbulence can be generated. Rotation is added to study the $\Lambda$-effect. A simple analytical closure model, based on the minimal tau-approximation (hereafter MTA, see e.g. Blackman \& Field \cite{BlackField2002}; Brandenburg et al.\ \cite{Brandea2004}), is used to compare with simulations in the cases with rotation. | \label{sec:conclusions} Shear flow turbulence simulations show that the ratio of turbulent to molecular viscosity increases linearly up to ${\rm Re} \approx 30$ with $\nu_{\rm t}/\nu\approx1.5\,\mbox{Re}$. For the largest Reynolds number the scaling seems somewhat shallower but the present data is not yet sufficient to substantiate this. The $\Lambda$-effect from homogeneous, anisotropic turbulence does not exhibit the puzzling features found in convection simulations. Further study is required in order to understand which of the neglected physics is responsible for the lack of these features. The MTA-closure is able to reproduce many of the qualitative aspects of the simulation results including a maximum of the horizontal stress at about $30^\circ$ latitude, with its largest value for $\mbox{Co}\approx0.5$. In the model, the vertical stress can have a maximum away from the equator for $\mbox{Co}\ga0.2$, which is not seen in the simulations. Nevertheless, both simulations and model have the largest vertical stress for $\mbox{Co}\approx0.3$. However, the model generally overestimates the magnitudes of the stresses. More detailed analysis of the simulation and closure results will be presented in a future publication. | 7 | 10 | 0710.5632 |
0710 | 0710.2538_arXiv.txt | We analyze the possibility of delensing Cosmic Microwave Background (CMB) polarization maps using foreground weak lensing (WL) information. We build an estimator of the CMB lensing potential out of optimally combined projected potential estimators to different source redshift bins. Our estimator is most sensitive to the redshift depth of the WL survey, less so to the shape noise level. Estimators built using galaxy surveys like LSST and SNAP recover up to 80-90\% of the potential fluctuations power at $l\leq 100$ but only \ax 10-20\% of the small-angular-scale power ($l\leq 1000$). This translates into a 30-50\% reduction in the lensing $B$-mode power. \par We illustrate the potential advantages of a 21-cm survey by considering a fiducial WL survey for which we take the redshift depth \zmax and the effective angular concentration of sources \nbar as free parameters. For a noise level of 1 $\mu$K arcmin in the polarization map itself, as projected for a CMBPol experiment, and a beam with $\theta_{\rm\scriptscriptstyle{FWHM}}$=10 arcmin, we find that going to \zmax=20 at \nbar=100 \gal yields a delensing performance similar to that of a quadratic lensing potential estimator applied to small-scale CMB maps: the lensing $B$-mode contamination is reduced by almost an order of magnitude. In this case, there is also a reduction by a factor of \ax4 in the detectability threshold of the tensor $B$-mode power. At this CMB noise level, the $B$-mode detection threshold is only $3\times$ lower even for perfect delensing, so there is little gain from sources with $z_{max}>20$. The delensing gains are lost if the CMB beam exceeds $\sim 20$ arcmin. The delensing gains and useful $z_{max}$ depend acutely on the CMB map noise level, but beam sizes below 10 arcmin do not help. Delensing via foreground sources does not require arcminute-resolution CMB observations, a substantial practical advantage over the use of CMB observables for delensing. | \label{I} The anisotropies in the CMB have been long recognized as a major probe for cosmology. The WMAP satellite has measured the temperature fluctuations up to $l_{max} <=$1000 and has confirmed the cosmological standard model of a power-law, flat, $\Lambda$CDM universe. Much hope lies with CMB polarization measurements because they have the potential to unveil some of the unknowns of inflation. While some predictions of inflation (such as nearly flat space curvature, nearly scale-invariant power spectrum, and Gaussianity of the primordial fluctuations) have been confirmed by CMB and large scale structure experiments, there is another prediction which is yet to be verified. This is the existence of an almost scale-invariant spectrum of gravitational waves, whose amplitude is directly related to the energy scale of inflation. The inflationary gravitational waves are tensor perturbations to the metric and we expect them to leave a curl-like signature in the CMB polarization field, i.e. a $B$-mode pattern. Density fluctuations, arising from scalar perturbations to the metric, create a gradient-like component in the polarization field, the $E$ mode. In the linear regime, density fluctuations do not create a $B$ mode, so a detection of the latter would confirm the existence of primordial gravitational waves; it would also provide the energy scale of inflation, which we could use to distinguish between different inflationary scenarios.\par CMB polarization measurements are difficult to carry out, primarily for two reasons. Firstly, the amplitude of the signal is very small: for example, the scalar $E$-mode power is 2-3 orders of magnitude smaller than the scalar temperature power, tensor $B$-mode much lower. Secondly, the polarization foregrounds are very poorly understood and they dominate the CMB signal at almost all relevant frequencies. Therefore, foreground removal and detector sensitivity are two of the most stifling limitations of a polarization experiment.\par $B$-mode measurements are additionally obstructed by WL contamination, especially of the recombination signal (l$\leq$100). Thus CMB polarization measurements are able to provide inflationary insights to the extent to which we can: 1. remove the foreground contribution to the overall signal and 2. delens the $B$ mode power and extract the tensor contribution to it. The delensing process consists of the reconstruction of the lensing potential from chosen observables. The estimated lensing potential is then used to evaluate the WL-created $B$-mode signal and to subtract it from the measured $B$-mode map. \par In this paper we probe the ability of weak lensing (WL) galaxy surveys to delens the CMB in the absence of foregrounds. To be specific, we try to answer two questions: what attributes should a galaxy or 21~cm WL survey and also a CMB polarization mission have in order to detect the tensor $B$ mode? What is the minimum amplitude of the $B$ mode, expressed in terms of the tensor-to-scalar ratio, $r$, that can be detected using galaxy or recombination observations for delensing? We take three examples of surveys to illustrate how our estimator works: the ground-based Large Synoptic Survey Telescope (LSST \footnote{\tt {http://www.lsst.org/}}), the space-based Supernova Acceleration Probe (SNAP \footnote{\tt {http://snap.lbl.gov/}}) and a toy model mimicking recombination-era 21-cm observations that we mention in more detail in section \S\ref{IV}. The outline of the paper is as follows. In \S\ref{II} we describe the WL contamination of the tensor $B$ mode. In \S\ref{III} we present our minimum-variance lensing estimator. In \S\ref{IV} we determine the minimum detectable $r$ when we use this estimator to delens. In \S\ref{V} we discuss our results and draw conclusions. Let us now briefly mention similar work existing in the literature. There has been a vast and impressive amount of work on the topic of CMB delensing and $r$-detection. The great majority of this work uses the CMB observables $\Theta, E, B$ to reconstruct the projected potential. Averaging over various quadratic combinations of the temperature field (e.g. see the work of \citet{1999PhRvD..59l3507Z}, \citet{1998A&A...338..767B}, \citet{2001PhRvD..64h3005H}, \citet{2001ApJ...557L..79H}) and the polarization field (e.g. \citet{2000PhRvD..62d3517G}, \citet{2002ApJ...574..566H}, \citet{2003PhRvD..67l3507K}) has been thoroughly considered for the projected potential reconstruction. To give a quick summary: one can build minimum-variance, unbiased estimators, using certain field statistics, as shown by \citet{2001ApJ...557L..79H}. For a post-Planck experiment (sensitivity of 0.3 $\mu$K arcmin and beam size of 3 arcmin), \citet{2002ApJ...574..566H} found that the most efficient of these estimators can map the potential up to l$\leq$1000. \citet{2002PhRvL..89a1304K} and \citet{2002PhRvL..89a1303K} used this last estimator to predict the minimum detectable $r$ as a function of CMB experimental characteristics. There is another, more promising method for lensing potential reconstruction, based on likelihood techniques. \citet{2003PhRvD..67d3001H} have built a maximum-likelihood estimator for the convergence field using temperature maps and have found its performance similar to that of the quadratic estimator introduced by \citet{2001ApJ...557L..79H}. \citet{2003PhRvD..68h3002H} found that the same maximum likelihood estimator built from polarization maps is even more effective: there is an order of magnitude reduction in the mean squared error in the lensing reconstruction compared to the quadratic estimator method, if the survey characteristics are adequate: sensitivity of 0.25 $\mu$K arcmin and a beam size of 2 arcmin. \citet{2005PhRvD..72l3006A} analyze the detectability of tensor $B$ modes in the presence of polarized dust emission, as a function of sky coverage; \citet{2006JCAP...01..019V} do a study of optimal surveys for $B$-mode detection, considering both dust and synchrotron emissions. One disadvantage of the reconstruction methods presented so far is that they require high-resolution CMB maps: they use arc-minute structures of the CMB fields, to reconstruct degree-scale maps of the deflection field, as explained by \citet{2001ApJ...557L..79H}. \citet{2005PhRvL..95u1303S} point out that one can use non-CMB observables to delens the CMB; in this case the requirement for high angular resolution of the CMB mission can be relaxed significantly. These authors determine lower limits for the detectable $r$ using the 21 cm radiation emitted by neutral hydrogen atoms to delens the CMB. We follow a similar approach here, but employ foreground galaxies instead of 21-cm emission as the source plane for delensing. | \label{V} In this paper we have examined the possibility of delensing $B$-mode polarization maps using galaxy WL surveys. We have proposed a weighted combination of projected potential estimators to different source redshift bins which optimally reconstructs the projected potential seen by the CMB. We have used three fiducial surveys to exemplify our estimator: LSST, SNAP and a generic survey relevant mostly to future 21-cm studies. These examples have different source redshift distribution, source density \nbar, and sky coverage $f_{\rm sky}$, and have enabled us to test the effect of each of these factors on the lensing potential estimator and also on its ability to reduce the WL contamination of $B$ mode maps. Using a $\Delta \chi^{2}$ test for the delensed $B$ mode field, we have determined the minimum value of the tensor-to-scalar ratio statistically distinguishable from 0 in optimally delensed data. Throughout this paper we have ignored the polarized foreground contamination and we have also made the assumption of Gaussianity for the delensed $B$ mode map. While this assumption may be inaccurate, especially for the cases when the efficiency of delensing is low, we expect the qualitative results of our work to hold. \par The lensing potential estimator is sensitive mostly to the redshift depth of the WL survey and reconstructs best the large angular scale multipoles. The performance of the estimator improves continually as higher redshift source galaxies are used. In the limit \zmax$\rightarrow z_{\scriptscriptstyle\rm CMB}$ and \nbar$\rightarrow\infty$, the reconstruction is perfect. An experiment like SNAP recovers \ax 90\% of the CMB projected potential power for $l\leq 100$ and \ax 20\% for $l\leq 1000$. The performance of LSST is a little worse, because it detects sources at an average redshift of $\langle z \rangle$=1, compared to $\langle z \rangle$=1.5 of SNAP. A box distribution (constant redshift distribution of source galaxies), provides better CMB potential estimator: if we go as far as \zmax=20, there is an order of magnitude improvement over SNAP, for the same angular concentration of galaxies. However, in this last case, the reduction in $r_{min}$, compared to the case where no delensing is done, is only of a factor of \ax 7 for $w^{-1/2}=0.1 \mu$K arcmin and $\theta_{\rm\scriptscriptstyle{FWHM}}$=10 arcmin. This happens because low-$l$ lensing $B$-modes are generated by beating of gravitational potential modes and $E$-modes on scales coresponding to $l$\ax1000, where the potential estimator is less faithful. In the case of LSST and SNAP, the reduction in $r_{min}$ relative to the no delensing case is only by \ax15\% and \ax50\% for the same CMB noise level and beam. We stress that delensing with foreground weak lensing offers significant gains in $r_{min}$ (factors of 5 or more) only when three conditions are met: \begin{enumerate} \item the noise in the CMB map is sufficiently low, $\sim 1\,\mu$K arcmin; \item the beam size of the CMB map is $<20$ arcmin; \item the lensing source distribution extends to $z\sim 20$ or greater. \end{enumerate} Delensing is not relevant to tensor mode detection if $w^{-1/2}>10 \mu$K arcmin or if $\theta_{\rm\scriptscriptstyle{FWHM}}$ \gt$2^{\,\circ}$. Also, there is no advantage in lowering the beam size beyond $\theta_{\rm\scriptscriptstyle{FWHM}}$=10 arcmin: for our delensing method, a 1 arcmin beam will do just as well as a 10 arcmin beam. This is perhaps the most important feature of our delensing technique, as both the quadratic estimator of \citet{2001ApJ...557L..79H} and the maximum likelihood estimator proposed by \citet{2003PhRvD..68h3002H} require beam sizes of 2-3 arcmin to yield their best performance. For the CMBPol mission, (which might reach a noise of 1 $\mu$K arcmin) delensing with a box of \zmax=20 results in $r_{min}\approx 2\times 10^{-5}$. If no delensing is applied, then $r_{min}\approx 8\times10^{-5}$, a factor of 4 worse. CMBPol detector noise keeps $r_{min}>7\times 10^{-6}$ even with perfect delensing, so there is less incentive to acquire delensing source planes above \zmax=20. Also for the CMBPol noise level, we have investigated the impact of $\bar n$ and have concluded that as long as we go to high redshift ($z_{max}>10$), even with a low density of source galaxies, we can still improve $r_{min}$ by a factor of a few. | 7 | 10 | 0710.2538 |
0710 | 0710.1011_arXiv.txt | We present a 40 minute time series of filtergrams from the red and the blue wing of the \halpha\ line in an active region near the solar disk center. From these filtergrams we construct both Dopplergrams and summed ``line center'' images. Several dynamic fibrils (DFs) are identified in the summed images. The data is used to simultaneously measure the proper motion and the Doppler signals in DFs. For calibration of the Doppler signals we use spatially resolved spectrograms of a similar active region. Significant variations in the calibration constant for different solar features are observed, and only regions containing DFs have been used in order to reduce calibration errors. We find a coherent behavior of the Doppler velocity and the proper motion which clearly demonstrates that the evolution of DFs involve plasma motion. The Doppler velocities are found to be a factor 2--3 smaller than velocities derived form proper motions in the image plane. The difference can be explained by the radiative processes involved, the Doppler velocity is a result of the local atmospheric velocity weighted with the response function. As a result the Doppler velocity originates from a wide range in heights in the atmosphere. This is contrasted by the proper motion velocity which is measured from the sharply defined bright tops of the DFs and is therefore a very local velocity measure. The Doppler signal originates from well below the top of the DF. Finally we discuss how this difference together with the lacking spatial resolution of older observations have contributed to some of the confusion about the identity of DFs, spicules and mottles. | \label{sec:intro} The solar chromosphere owes its name to the reddish rim that appears above the lunar limb during solar eclipses. This reddish color mostly stems from the Balmer {\halpha} spectral line which makes this line one of the most important chromospheric diagnostics. Due to the highly dynamic state of the chromosphere and strong NLTE effects, the line formation processes are still not yet fully understood \citep[e.g.][]{radyn,2006Leenaarts}. This is an important shortcoming in our interpretation tools which makes \halpha\ observations traditionally difficult to interpret. Due to the highly fibrilar structure of the chromosphere \citep{1908Hale}, a strong influence from magnetic fields on the chromosphere has been suspected for about a century. The most common of these fibrilar magnetic fine structures are the jet-like structures known as spicules, mottles, and dynamic fibrils (DFs). In short, spicules are traditionally observed at the limb, mottles on disk in the quiet Sun, and DFs in active regions. Whether or not these structures are manifestations of the same phenomenon viewed at different angles have been the subject of a long standing discussion \citep[e.g.][]{1968Beckers,1992Grossman,1994Tsiropoula,1995Suematsu, 2001Christo,2007Rouppe}. One important argument against these structures being caused by the same mechanism has been the difference in the measured absolute velocities \citep{1992Grossman}. Other authors have done direct measurements of mottles crossing the limb \citep{2001Christo}. They also state that since both proper motions and Doppler motions are used in the comparisons, systematic errors are probably introduced. Such errors might also be amplified by the rather limited spatial resolution of some of the data sets used. The detailed analysis of DFs has accelerated in recent years \citep[e.g.][]{2004dePontieu,2006deWijn,2006Hansteen,2007dePontieu, 2007Lars} due to major advances in both observational techniques and simulation efforts. One of the main conclusions from these studies is that DFs are driven by magneto-acoustic shocks caused by p-mode oscillations and convective flows leaking into the chromosphere. In a recent study, \citet[][from now on paper 1]{2007bLangangen} presented spectroscopic analysis of DFs as seen in one of the Ca~{\small{II}}~IR lines. Numerical analysis of the line formation process showed a much lower DF velocity derived from Doppler measurements as compared to the proper motion velocity. This was found to be due to both the low formation height and the extensive width of the contribution function of the Ca~{\small{II}}~IR line. Furthermore, the DFs analyzed in paper 1 showed mass motion, thus ruling out any ionization/temperature wave as explanation model for DFs \citep[e.g.][and references therein]{2000Sterling}. With the advantage of well sampled spectral line profiles, the number of analysed DFs in paper~1 was rather modest due to the limited spatial coverage of the spectrograph slit. With the current data we exploit the wide spatial coverage of a tunable filter instrument, at the expense of limited spectral resolution. \begin{figure*}[!ht] \includegraphics[width=\textwidth]{f1.eps} \caption{ Field of view (FOV) for one wideband image (left) and the corresponding sum of the two narrow band filtergrams (right). In the narrow band image several fibrilar structures are seen, and some DF axes are marked (solid white lines) for illustrative purposes. } \label{plotone} \end{figure*} In this paper we add to the understanding of these jet like structures by analysis of Dopplergrams obtained in an active region close to the disk center, hence the observed jet structures are commonly known as DFs. In \S\ref{sec:obs} we describe the observational program and the instrumentation. The data reduction and the calibration method is explained in \S\ref{sec:datareduction}. In \S\ref{sec:Obsres} we present the results of our measurements. We discuss our results in \S\ref{Disc} and finally we summarize the results in \S\ref{Sum}. | \label{Disc} The correlation between the maximum velocities and decelerations found from the proper motion measurements is similar to the correlation found by \citet{2006Hansteen} and \citet{2007dePontieu}. This is illustrated in Fig.~\ref{plotseven} by the grey-scaled cloud shown in the background. This correlation between the deceleration and the maximum velocity is known to be the signature of shock waves being the driving mechanism of DFs \citep{2006Hansteen,2007dePontieu,2007Lars}. Further support for this driving mechanism comes from the fact that we find a coherent behavior between the evolution of the Doppler signal and the proper motion for a large fraction of the DFs. This is a strong indication that there is actual plasma motion occurring during the life time of DFs. This supports the findings of paper~1, but based on a much larger sample. The Doppler measurements show a similar correlation as for the proper motion, but with much lower absolute values for both the decelerations and maximum velocities. One possible explanation for these lower values could be high inclination angles of the DF trajectories with the LOS. One could expect to be able to derive the full trajectory vector by combining the two measured deceleration components. This naive method would give very high inclination angles, typically $75^\circ$. We know, however, that this can not be the true inclination angle, since the Doppler velocity is a result of the local atmospheric velocity weighted with the response function to velocity over an extended height. In contrast, the measured proper motion is very local due to the high contrast boundary between the top of the fibril and the surroundings. Combining these two measurements leads to highly overestimated inclination angles. The difference in absolute values must be considered in the context of the results from Paper~1. The lower Doppler velocities found from the Ca~{\small{II}}~IR line was explained by a combination of lower formation height and extended formation range. This is probably also the case for the \halpha\ line, but the formation height usually extends over a larger height range as compared to the Ca~{\small{II}}~IR line. The lacking Doppler shifts in ${\sim}15$\% of the DFs can either be caused by very high inclination angles, or their driving mechanism is fundamentally different and the evolution of these DFs is not a result of mass motion. We believe that high inclination angles combined with the uncertainties in the measurements is a more plausible explanation for the lacking Doppler signals. The identification method of DFs introduces a bias toward the more inclined DFs. There are a number of suggestive cases where DFs are visible in the Dopplergrams but no clear signature can be seen in the corresponding intensity images. We refrain from measuring these DFs since this will complicate a comparison with other data sets. Furthermore, the identification of these DFs is not objective and we expect the measurement errors to be unacceptably high since low inclination angles would lead to potentially high Doppler velocities. Due to the lacking spectral sampling this could lead to strong saturation effects in the measured Doppler velocities. \subsection{Spicules, mottles and fibrils} The identification of the disk counterpart of spicules was already an important question forty years ago \citep[e.g][]{1968Beckers}. One of the main problems was to reconcile the velocities measured in spicules with those measured in mottles. This problem was also the main concern of \citet{1992Grossman}. They conclude that since the velocities are much larger in spicules than in mottles, the two could not be the same structure seen at different viewing angles. They, however, admit that the seeing might impair their results if the structures observed were smaller than $1\arcsec{}$. Later studies of mottles and spicule properties lead to the conclusion that spicules and mottles are in fact the same feature seen at different angles \citep{1993Tsiropoula,1994Tsiropoula}. \citet{1994Tsiropoula} showed, using cloud modelling, that the proper motion and the cloud velocities were consistent. In a more recent work, \citet{2001Christo} use a limb darkening correction method to directly observe mottles crossing the limb. They argue that the main reason for earlier confusion is caused by the lacking spatial resolution of the observations. In our study it is clear that the Doppler signal originates from spatially resolved structures. The excellent quality of the observations largely removes the errors due to lacking spatial resolution. Assuming reasonable inclination angles, i.e. the mean angles being not very large nor very small, we can conclude that the Doppler velocities are typically a factor of $\sim 2$--$3$ smaller than the corresponding proper motion. A similar, but larger difference was reported by \citet{1994Tsiropoula}, we believe that this difference can be attributed to worse spatial resolution. As discussed above, radiative transfer processes are the fundamental reason for the Doppler velocities being lower than the proper motion. We argue that the fundamental differences between Doppler and proper motion velocities that we find for DFs, also are valid for similar measurements on spicules and mottles. Besides the arguments of lacking spatial resolution, we believe that this difference was an important contributor to the earlier confusion about the unification of mottles and spicules. Similar work on spatially resolved limb spicules is needed to finally settle this discussion. | 7 | 10 | 0710.1011 |
0710 | 0710.4919_arXiv.txt | Quantum gravity may remove classical space-time singularities and thus reveal what a universe at and before the big bang could be like. In loop quantum cosmology, an exactly solvable model is available which allows one to address precise dynamical coherent states and their evolution in such a setting. It is shown here that quantum fluctuations before the big bang are generically unrelated to those after the big bang. A reliable determination of pre-big bang quantum fluctuations would require exceedingly precise observations. | Cosmology, as the physics of the universe as a whole, places special limitations on scientific statements to be reasonably inferred from observations. On a general basis, this was discussed in \cite{Uncertain} who noted that some general properties seem to arise automatically during cosmological evolution. Their verification for our universe then does not reveal anything about its possibly more complicated past. For instance, isotropization \cite{Isotropize} demonstrates that the observation of a nearly isotropic present universe does not present much useful information on an initial state at much earlier times. Similarly, one may add decoherence to the list which can be seen as doing the same for the observation of a nearly classical present universe: Most initial states would decohere and arrive at a semiclassical state; thus its observation does not rule out stronger quantum behavior at earlier stages. Given this, it is surprising to see recent discussions about a universe before the big bang in several approaches to quantum gravity. Not just statements about the possibility of the universe having existed before the big bang are being made, which could in principle be inferred from a general analysis of equations of motion, but even assumptions on its classicality or claims about the form of its state such as its fluctuations at those times are put forward. Even if this may be possible theoretically, it raises the question how much one can really learn about a pre-big bang universe from observations. A solvable model of quantum cosmology is used here to draw cosmological conclusions within this model about the possible nature of the big bang and a universe preceding it. The viewpoint followed has been spelled out in \cite{BeforeBB} where also some of the results have already been reported. This model does show a bounce at small volume instead of the classical singularity present in solutions of general relativity. We therefore use it to shed light on the general questions posed above. The promotion of this specific model is not intended as a statement that its bounce would be generic for quantum gravity even within the same framework. In fact, the conclusion of the bounce in this and related models available so far is based on several specific properties which prevent a generic statement about this form of singularity removal. We rather take the following viewpoint: Assume there is a theoretical description of a bouncing universe; what implications can be derived for its pre-bounce state? The specific model used is distinguished by a key fact which makes it suitable for addressing such a general question: As a solvable model, the system displays properties of its dynamical coherent states which, to some degree, are comparable to those of the harmonic oscillator. There is no back-reaction of quantum fluctuations or higher moments of a state on the time evolution of its expectation values. Their trajectory is thus independent of the spreading of a state if it occurs, a property which is well known for the harmonic oscillator (for instance, as a simple consequence of the Ehrenfest theorem). This behavior is certainly special compared to other quantum systems, but the precise derivation of properties of dynamical coherent states it allows are nevertheless important. For instance, a large part of theoretical quantum optics is devoted to a computation of fluctuations in squeezed coherent states of the harmonic oscillator. Similarly, the solvable model studied here allows detailed calculations of its dynamical coherent states which are of interest for quantum cosmology. In particular, the solvable model we will be using eliminates the classical big bang singularity by quantum effects. Dynamical coherent states highlight the behavior of fluctuations of the state of the universe before and after the big bang. Care is, however, needed for the physical interpretation of the results. While quantum optics allows one to prepare a desired state and perform measurements on it, quantum cosmology has to make use of the one universe state that is given to us. Unlike quantum optics, where states can be prepared to be close to harmonic oscillator states, we cannot realize states of the solvable quantum cosmological model. The real universe is certainly very different from anything solvable, and so the availability of a certain feature or numerical result in a solvable model is unlikely to be related to a realistic property of the universe. Thus, as spelled out in \cite{BeforeBB}, we focus on pessimism in our analysis of the solvable model: the {\em in}ability of making certain predictions even in a fully controlled model is likely to be a reliable statement much more generally; adding complications of a real universe would only make those predictions even more difficult. For solvable systems of this kind it is easier by far to solve equations of motion for expectation values and fluctuations directly, rather than taking the detour of a specifically represented wave function. Properties of coherent states are then determined by selecting solutions of fluctuations which saturate uncertainty relations. We will first illustrate this procedure for the harmonic oscillator with an emphasis on squeezed states. We present this brief review in section \ref{Squeezed} intended as an introduction of the methods then used in quantum cosmology. The main part of this paper is an application of those methods to the solvable system of quantum cosmology, as well as a physical discussion. Instead of different cycles of a harmonic oscillator, in quantum cosmology we will be dealing with the pre- and post-big bang phase of a universe. Just as fluctuations in a squeezed harmonic oscillator state can oscillate during the cycles, fluctuations in quantum cosmology may change from one phase to the next. Our main concern will be the reliability of predictions about the precise state of the universe before the big bang, based on knowledge we can achieve after the big bang. | 7 | 10 | 0710.4919 |
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0710 | 0710.3364_arXiv.txt | The focusing of the radiation generated by a polarization current with a superluminally rotating distribution pattern is of a higher order in the plane of rotation than in other directions. Consequently, our previously published asymptotic approximation to the value of this field outside the equatorial plane breaks down as the line of sight approaches a direction normal to the rotation axis, \ie is nonuniform with respect to the polar angle. Here we employ an alternative asymptotic expansion to show that, though having a rate of decay with frequency ($\mu$) that is by a factor of order $\mu^{2/3}$ slower, the equatorial radiation field has the same dependence on distance as the nonspherically decaying component of the generated field in other directions: it, too, diminishes as the inverse square root of the distance from its source. We also briefly discuss the relevance of these results to the giant pulses received from pulsars: the focused, nonspherically decaying pulses that arise from a superluminal polarization current in a highly magnetized plasma have a power-law spectrum (\ie a flux density $S\propto\mu^\alpha$) whose index ($\alpha$) is given by one of the values $-2/3$, $-2$, $-8/3$, or $-4$. | Radiation by polarization currents whose distribution patterns move faster than light {\it in vacuo} has been the subject of several theoretical and experimental studies in recent years \cite{BessarabAV:FasEsi,ArdavanA:Exponr,BessarabAV:Expser,BolotovskiiBM:Radssv,% BolotovskiiBM:Radbcm,ArdavanH:Genfnd,ArdavanH:Speapc,ArdavanH:Morph}. When the motion of its source is accelerated, this radiation exhibits features that are not shared by any other known emission. In particular, the radiation from a rotating superluminal source consists, in certain directions, of a collection of subbeams whose azimuthal and polar widths narrow (as $R\subP^{-3}$ and $R\subP^{-1}$, respectively) with distance $R\subP$ from the source \cite{ArdavanH:Morph}. Being composed of tightly focused wave packets that are constantly dispersed and reconstructed out of other waves, these subbeams neither diffract nor decay in the same way as conventional radiation beams. The field strength within each subbeam diminishes as $R\subP^{-1/2}$, instead of $R\subP^{-1}$, with increasing $R\subP$ \cite{ArdavanH:Genfnd,ArdavanH:Speapc,ArdavanH:Morph}. In earlier treatments \cite{ArdavanH:Speapc,ArdavanH:Morph}, we evaluated the field of a superluminally rotating extended source by superposing the fields of its constituent volume elements, \ie by convolving its density with the familiar Li\'enard-Wiechert field of a rotating point source. This Li\'enard-Wiechert field is described by an expression essentially identical to that which is encountered in the analysis of synchrotron radiation, except that its value at any given observation time receives contributions from more than one retarded time. The multivalued nature of the retarded time gives rise to the formation of caustics. The wave fronts emitted by each constituent volume element of a superluminally moving accelerated source possess a cusped envelope on which the field is infinitely strong (see Figs.~1 and 4 of Ref.~\cite{ArdavanH:Morph}). Correspondingly, the Green's function for the problem is nonintegrably singular for those source elements that approach the observer along the radiation direction with the speed of light and zero acceleration at the retarded time (see Fig.~3 of Ref.~\cite{ArdavanH:Morph}): the cusp of the envelope of wave fronts emanating from each such element is a spiraling curve extending into the far zone that passes through the position of the observer. When the source oscillates at the same time as rotating, the Hadamard finite part of the divergent integral that results from convolving the Green's function with the source density has a rapidly oscillating kernel for a far-field observation point. The stationary points of the phase of this kernel turn out to have different orders depending on whether the observer is located in or out of the equatorial plane. To reduce the complications posed by the higher-order stationary points of this phase, we restricted the asymptotic evaluation of the radiation integral thus obtained in Refs.~\cite{ArdavanH:Speapc,ArdavanH:Morph} to observation points outside the plane of rotation, \ie to spherical polar angles $\theta\subP$ that do not equal $\pi/2$. The purpose of this paper is to evaluate the field of a superluminally rotating extended source also for the smaller class of observers at polar coordinate $\theta\subP=\pi/2$ and to obtain, thereby, a more global description of the nonspherically decaying radiation that is generated by such a source. The asymptotic expansion presented in Refs.~\cite{ArdavanH:Speapc,ArdavanH:Morph} breaks down in the case of a subbeam that is perpendicular to the rotation axis because there is a higher-order focusing associated with the waves emitted by those source elements whose actual speeds (rather than the line-of-sight components of their speeds) equal the speed of light as they approach the observer with zero acceleration. Here, we present a brief account of the background material on the radiation field of a rotating superluminal source in Section 2, and the asymptotic evaluation of this field for an equatorial observer in Section 3. In Section 4, we give a description of the spectral properties of the nonspherically decaying component of this radiation in the light of the present results and those obtained in Refs.~\cite{ArdavanH:Speapc,ArdavanH:Morph}, and discuss the relevance of these properties to pulsar observations. | 7 | 10 | 0710.3364 |
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0710 | 0710.2542_arXiv.txt | {INTEGRAL, the European Space Agency's $\gamma$-ray observatory, tripled the number of super-giant high-mass X-ray binaries (sgHMXB) known in the Galaxy by revealing absorbed and fast transient (SFXT) systems.} {In these sources, quantitative constraints on the wind clumping of the massive stars could be obtained from the study of the hard X-ray variability of the compact accreting object.} {Hard X-ray flares and quiescent emission of SFXT systems have been characterized and used to derive wind clump parameters.} {A large fraction of the hard X-ray emission is emitted in the form of flares with a typical duration of 3 ks, frequency of 7 days and luminosity of $10^{36}$ erg/s. Such flares are most probably emitted by the interaction of a compact object orbiting at $\sim10~R_*$ with wind clumps ($10^{22-23}$ g) representing a large fraction of the stellar mass-loss rate. The density ratio between the clumps and the inter-clump medium is $10^{2-4}$ in SFXT systems. } {The parameters of the clumps and of the inter-clump medium, derived from the SFXT flaring behavior, are in good agreement with macro-clumping scenario and line driven instability simulations. SFXT have probably a larger orbital radius than classical sgHMXB.} | \label{sec:intro} Stellar winds have profound implications for the evolution of massive stars, on the chemical evolution of the Universe and as source of energy and momentum in the interstellar medium. Photons emitted by massive stars directly transfer momentum to the stellar wind through absorption in numerous Doppler shifted optically thick spectral lines \citep{cak1975} and drive the wind to highly supersonic speeds. Such line driven stellar winds are very unstable \citep{ocr1988,feldmeier1995}. The collisions between high speed and low density material with slower gas trigger strong shocks, form high density contrasts and wind clumps and lead to thermal X-ray emission. There are multiple observational lines of evidence for wind clumping in massive stars from wind line profiles \citep{Bouret2005,Fullerton2006}, discrete absorption components \citep{PrinjaHowarth1986}, variable line profiles \citep{Lepine1999,Markova2005}, polarimetry \citep{Lupie1987, Davies2007}, X-rays continuum \citep{CassinelliOlson1979} and line emission \citep{OskinovaFeldmeierHamann2006}. Wind clumping has an important effect on mass-loss rate diagnostics that depends on the square of the wind density. It leads to a reduction of the stellar mass-loss rate estimates by a factor $f_V^{-0.5}$ where $f_V$ is the clump volume filling factor \citep{Abbott1981,Fullerton2006}. A reduction by a factor $>3$ becomes problematic for massive star evolution \citep{Hirschi2007}. Optically thick clumps may however help to reconcile wind clumping and usual mass-loss rates by reducing spectral features and free-free emission \citep{OskinovaHamannFeldmeier2007}. Indirect measures of the structure of massive star winds are possible in binary systems through the analysis of the interaction between the stellar wind and the companion or its stellar wind. In colliding wind binaries, \cite{Schild2004} and \cite{Pollock2005} have pointed out that the observed X-ray column densities are much lower than expected from smooth winds and that clumped winds are the probable solution \citep[see also][]{pittard2007}. In wind-fed accreting High-Mass X-Ray Binaries (HMXB), clumping has been invoked to explain the orbital variability of X-ray line profiles emitted by the photo-ionized wind \citep{Sako2003,Vandermeer2005}. In this paper we aim at studying the clumping of stellar wind in HMXB using the hard X-ray variability observed by the IBIS/ISGRI instrument \citep{ubertini03AA,lebrun03AA} on board the INTErnational Gamma-Ray Astrophysics Laboratory \citep[INTEGRAL,][]{winkler03AA} which could be used to probe clump parameters and density contrasts in the stellar wind. This study follows a first paper \citep{Leyder2007} interpreting the flaring behavior of \object{IGR J08408$-$4503} in term of wind clumping. Classical wind-fed super-giant HMXB (sgHMXB) are made of a compact object orbiting within few stellar radii from a super-giant companion. Recently INTEGRAL almost tripled the number of sgHMXB systems known in the Galaxy and revealed a much more complex picture with two additional families of sources: the highly absorbed persistent systems \citep{walter04esasp,walter2006} and the super-giant fast X-ray transient (SFXT) systems \citep{Negueruela2006}. The highly absorbed systems have orbital and spin periods similar to those of classical Roche-lobe underflow sgHMXB, however the absorbing column densities are much higher than observed on average in classical systems \citep{walter2006}. The fast transient systems are characterized by fast outbursts, by a low quiescent luminosity and by super-giant OB companions \citep{Sguera2006,Negueruela2007}. The sample of sources and the INTEGRAL data analysis are described in Sect. 2. Their hard X-ray variability is discussed in the context of clumpy stellar winds in Sect. 3. Finally Sect. 4 summarizes our principal conclusions. | INTEGRAL tripled the number of super-giant HMXB systems known in the Galaxy and revealed two new populations: the absorbed and the fast transient (SFXT) systems. The typical hard X-ray variability factor is $\lesssim 20$ in classical and absorbed systems and $\gtrsim 100$ in SFXT. We have also identified some ``intermediate'' systems with smaller variability factors that could be either SFXT or classical systems. The SFXT behavior is best explained by the interaction between the accreting compact object and a clumpy stellar wind \citep{IntZand2005,Leyder2007}. Using the hard X-ray variability observed by INTEGRAL in a sample of SFXT we have derived typical wind clump parameters. The compact object orbital radius are probably relatively large ($10~R_*$) and the clumps which generate most of the hard X-ray emission have a size of a few tenth of $R_*$. The clump mass is of the order or $10^{22-23}~\rm{g}$ (for a column density of $10^{22-23}~\rm{cm}^{-2}$) and the corresponding mass-loss rate is $10^{-(5-6)}~\rm{M_{\odot}/y}$. At the orbital radius, the clump separation is of the order of $R_*$ and their volume filling factor is $0.02$. Depending how the clump density varies with radius, the average volume filling factor could be as large as 0.1. These parameters are in good agreement with the macro clumping scenario proposed by \cite{OskinovaHamannFeldmeier2007}. The observed ratio between the flare and quiescent count rates indicate density ratios between the clumps and the inter-clump medium which vary between 15 to 50 in ``Intermediate'' systems and $10^{2-4}$ in SFXT. Such ratios and the observed clump densities are in reasonable agreement with the predictions of line driven instabilities at large radii \citep{Runacres2005}. The main difference between classical sgHMXB and SFXT could be the orbital radius of the compact object. At small orbital radius ($R_{orb}\sim2~R_*$) the systems are persistent and luminous. At larger radius and if wind clumping takes place the fast transient SFXT behavior is observed. | 7 | 10 | 0710.2542 |
0710 | 0710.0547_arXiv.txt | We present new $H$-band echelle spectra, obtained with the NIRSPEC spectrograph at Keck II, for the massive star cluster ``B'' in the nearby dwarf irregular galaxy NGC~1569. From spectral synthesis and equivalent width measurements we obtain abundances and abundance patterns. We derive an Fe abundance of [Fe/H]=$-0.63\pm0.08$, a super-solar [$\alpha$/Fe] abundance ratio of $+0.31\pm0.09$, and an O abundance of [O/H]=$-0.29\pm0.07$. We also measure a low $\rm ^{12}C/^{13}C\approx 5\pm1$ isotopic ratio. Using archival imaging from the Advanced Camera for Surveys on board HST, we construct a colour-magnitude diagram (CMD) for the cluster in which we identify about 60 red supergiant (RSG) stars, consistent with the strong RSG features seen in the $H$-band spectrum. The mean effective temperature of these RSGs, derived from their observed colours and weighted by their estimated $H$-band luminosities, is 3790 K, in excellent agreement with our spectroscopic estimate of $T_{\rm eff} = 3800\pm200$ K. From the CMD we derive an age of 15--25 Myr, slightly older than previous estimates based on integrated broad-band colours. We derive a radial velocity of $\langle v_r \rangle=-78\pm 3$ km/s and a velocity dispersion of $9.6\pm0.3$ km/s. In combination with an estimate of the half-light radius of $0\farcs20\pm0\farcs05$ from the HST data, this leads to a dynamical mass of $(4.4\pm1.1)\times10^5$ M$_\odot$. The dynamical mass agrees very well with the mass predicted by simple stellar population models for a cluster of this age and luminosity, assuming a normal stellar IMF. The cluster core radius appears smaller at longer wavelengths, as has previously been found in other extragalactic young star clusters. | Massive star clusters are potentially useful test particles for constraining the evolutionary histories of their host galaxies. They can remain observable for the entire lifetime of a galaxy (as illustrated by the old \emph{globular clusters} which surround every major galaxy), and they are bright enough to be studied in detail well beyond the Local Group. The internal velocity broadening of their spectra is typically only a few km/s, making detailed abundance analysis at high spectral resolution feasible. In a previous paper we have taken a first step towards exploiting this potential by analysing $H$ and $K$-band spectra of a young massive star cluster (YMC) in the nearby spiral galaxy NGC~6946 \citep{lar06}. Here we apply a similar analysis to one of the young star clusters in the nearby (post-) starburst galaxy NGC~1569. NGC~1569 was one of the first galaxies in which the presence of exceptionally bright, young star clusters was suspected. A spectrum of one of the two bright ``stellar condensations'' in NGC~1569 was obtained already by \citet{mayall35}, although the true nature of these objects was probably first discussed in detail by \citet{as85}. They found that the spectra are of composite nature, and also noted that observations with the Hubble Space Telescope (HST) would definitively settle the issue whether or not these objects are really ``super star clusters''. Indeed, pre-refurbishment mission HST observations by \citet{oco94} settled the issue by showing that both objects are extended, with half-light radii of about two pc. Based on high-dispersion spectroscopy from the HIRES spectrograph on the Keck~I telescope, a dynamical mass of $\approx 10^6$ M$_\odot$ was soon after estimated for the brighter of the two objects, NGC~1569-A \citep{hf96,stern98}. The clusters in NGC~1569 are prime candidates for abundance analysis in the near-infrared where they are very bright. The IR holds a particular advantage over optical studies for NGC~1569 due to the large amount of foreground extinction. \citet{gg02} found the integrated $H$-band spectra of both clusters A and B to be well approximated by red supergiant templates. Although NGC~1569-A is the brighter of the two, it is actually a binary cluster itself with some evidence for a (small) age difference between the two components \citep{guido97,maoz01,origlia2001}. Whether or not the two components are physically connected or the result of a chance projection is unknown. In this paper we concentrate on NGC~1569-B. In addition to providing insight into the histories of their host galaxies, young star clusters with masses in excess of $10^5$ M$_\odot$ also offer an excellent opportunity to study large samples of coeval massive stars. Such clusters are rare in the Milky Way and even in the Local Group, so it is necessary to extend the search to a larger volume. The most massive known young star clusters in the Milky Way disk are Westerlund 1 \citep{clark05} and an object discovered in the 2MASS survey \citep{figer06}, both of which may have masses approaching $10^5$ M$_\odot$. Young clusters of similar masses are found in the Large Magellanic Cloud \citep{vdb99}, but these pale in comparison with objects like NGC~1569-A and NGC~1569-B. We have obtained new NIRSPEC $H$-band spectra of NGC~1569-B, optimised for abundance analysis, and we additionally present photometry for \emph{individual} stars in the cluster derived from archival observations with the high resolution channel (HRC) of the Advanced Camera for Surveys (ACS) on HST. The HST data provide an independent verification of the stellar parameters (notably $T_{\rm eff}$ and $\log g$) derived from the spectral analysis, and also allow us to compare the observed colour-magnitude diagram (CMD) with standard isochrones. Finally, we use new measurements of the structural parameters and velocity dispersion to derive the cluster mass and mass-to-light ratio and compare with predictions by simple stellar population (SSP) models. We begin by briefly describing the data in \S\ref{sec:data}. The analysis and our main results then follow in \S\ref{sec:results}, where we first discuss the near-infrared spectroscopy and the abundance analysis (\S\ref{sec:abundance}). We then proceed to construct a CMD from the HST imaging (\S\ref{sec:cmd}) from which we derive stellar parameters for the red supergiants (RSGs) that are compared against the spectroscopic results (\S\ref{sec:speccmp}) and theoretical isochrones (\S\ref{sec:interp}). In \S\ref{sec:virmass} we measure structural parameters for NGC~1569-B and combine these with velocity dispersion measurements to derive a dynamical mass estimate. Finally, some additional discussion and a summary are given in \S\ref{sec:summary}. | \label{sec:summary} We have presented a detailed investigation of the stellar content and other properties of the massive star cluster 'B' in NGC~1569. Using new high S/N $H$-band echelle spectroscopy from the NIRSPEC spectrograph on Keck~II, we have carried out abundance analysis of red supergiant stars in the cluster. We find an iron abundance of [Fe/H] = $-0.63\pm0.08$, close to that of SMC field stars. Our estimate of the oxygen abundance, [O/H] = $-0.29\pm 0.07$, is about a factor of two higher than that derived for H{\sc ii} regions \citep{sdk94,dev97,ks97}, and the resulting super-solar [O/Fe] abundance is unlike that observed in the Magellanic Clouds and young stellar populations in the Milky Way. However, according to our measurements NGC~1569-B also differs from these galaxies by having a higher [$\alpha$/Fe] ratio. The difference between our [O/H] measurement for NGC~1569-B and the O abundance derived for H{\sc ii} regions is greater than the formal uncertainties on either value, and our oxygen abundance is higher than the nebular abundances in most dwarf galaxies \citep[e.g.][]{vad07}. This raises the question whether one (or both) methods suffers from possible systematic errors, or if there might be a genuine difference between cluster and H{\sc ii} region abundances in NGC~1569 -- perhaps due to an ab initio oxygen enhancement of the gas out of which the cluster formed. Since oxygen is the only element in common between the cluster and H{\sc ii} region data, a detailed comparison of the abundance patterns is unfortunately not possible. However, it is of interest to note that the very strong Wolf-Rayet features in the spectrum of cluster A also suggest a high metallicity \citep{maoz01}. NGC~1569 is a popular target for testing models of chemical evolution, and it is generally recognised that the strong galactic wind observed in the galaxy must play an important role. \citet{martin02} observed a super-solar $[\alpha/{\rm Fe}]$ ratio for the wind, as we do for NGC~1569-B. However, it is unclear to what extent the two are related as material ejected in the outflow may not participate in star formation, and in any case not before it has had time to cool down and fall back into the galaxy. The chemical evolution of an NGC~1569-type galaxy has been modelled in detail by \citet{rec06}, assuming a variety of bursty star formation histories. Their chemo-dynamical simulations show that elements produced in the last burst of star formation generally do not get mixed with the ISM in the galaxy, but instead get injected into the hot gas phase. None of their models produce an O abundance as high as the one observed by us, and the model that comes closest ($12 + \log$ (O/H) $\sim 8.4$) severely underpredicts the N/O abundance of the H{\sc ii} regions observed by \citet{ks97}. Generally, the best fits to the H{\sc ii} region abundances are obtained for a ``gasping'' star formation history, but these models all predict a maximum $12 + \log$ (O/H) $\approx$ 8.1. Similarly, the O abundance predicted by the models of \citet{romano06} for NGC~1569 never exceeds $12 + \log$ (O/H) = 8.41. We speculate that the galactic wind which is responsible for removing metals may be less efficient in doing so near the bottom of the potential well where the massive clusters formed. Unfortunately, none of the models include predictions for the wide range of other elements we are observing here. The Small Magellanic Cloud presents an interesting comparison case. Early spectroscopic studies and Str{\"o}mgren photometry indicated that the young cluster NGC~330 is about 0.5 dex more metal-poor than the surrounding field \citep[and references therein]{gr92}. High-dispersion spectroscopy by \citet{hill99} showed a smaller and only marginally significant difference between the field stars and NGC~330, although still in the sense that the cluster is more metal-poor than the field ([Fe/H] = $-0.82\pm0.10$ vs.\ [Fe/H] = $-0.69\pm0.11$). The [O/Fe] abundance ratios in these stars were all found to be sub-solar ([O/Fe] $\approx$ $-$0.15 to $-$0.3 dex), as in the LMC and in young Galactic supergiants \citep{hbs97}. \citet{hill99} also found the $\alpha$-elements abundances (Mg, Ca, Ti) relative to Fe to be around Solar for NGC~330. \citet{gw99} found slightly lower iron abundances of [Fe/H] = $-0.94\pm0.02$ for K supergiants in NGC~330 than \citet{hill99} and an [O/Fe] ratio closer to Solar. They derived somewhat enhanced [Ca/Fe], [Si/Fe] and [Mg/Fe] ratios ($+0.18$, $+0.32$ and $+0.11$), in contrast to sub-solar ratios derived for B stars. While there is some evidence for a difference in metallicity between NGC~330 and the SMC field, the situation there seems to contrast with the case of NGC~1569-B where we find a \emph{higher} oxygen abundance of the cluster compared to the H{\sc ii} regions and an \emph{enhanced} [O/Fe] ratio. Our mean [$\alpha$/Fe] = $+0.31\pm0.09$ is also higher than that derived for the stars in NGC~330. One significant difference may be the mass of NGC~1569-B, which is probably an order of magnitude greater than that of NGC~330 \citep{fb80}. The spectral analysis returns a mean $T_{\rm eff} = 3800\pm200$ K and $\log g \approx 0.0$ for the RSGs in NGC~1569-B. We have checked these values using resolved photometry of the RSGs in the cluster, derived from archival HST/ACS images. About 60 RSGs are easily identifiable in the colour-magnitude diagram, and from their $m_{\rm F555W} - m_{\rm F814W}$ colours we derive a mean effective temperature of $\langle T_{\rm eff}\rangle = 3850$ K and $\log g = 0.1$. An even closer match to the spectroscopic $T_{\rm eff}$ is obtained if we weigh the $T_{\rm eff}$ estimates for the individual stars by their $H$-band luminosities, in which case we get $\langle T_{\rm eff} \rangle = 3790$ K. We have compared the CMD with $Z=0.004$ and $Z=0.008$ isochrones from the Padua and Geneva groups. Both sets of models provide the best match to the observed colours and magnitudes of the RSGs for $Z=0.008$ and ages of 15--25 Myrs, but no isochrone can reproduce the observed CMD in detail. For these sub-solar metallicities, the models predict higher (by up to a few 100 K) mean effective temperatures $\langle T_{\rm eff} \rangle$ for the RSGs than observed. Since the RSGs generally become cooler with increasing metallicity, models of higher metallicity are favoured, but models of Solar metallicity would be required to match the observed colours. For such models, however, the \emph{blue} supergiants become much too cool to match the observations. The observed ratio of blue to red supergiants (BSG/RSG = $0.39\pm0.10$) is also significantly lower than most of the model predictions. Although we are detecting an impressive number of RSGs, it should be noted that since we are only resolving stars located well outside the half-light radius, the actual number of RSGs in NGC~1569-B is likely to be more than a factor of two greater than the number we have detected. We derive a velocity dispersion of $9.6\pm0.3$ km/s from the integrated spectrum of NGC~1569-B, somewhat higher than the value of 7.5 km/s of \citet{gg02}. Combining this with an estimate of the cluster half-light radius of $0\farcs20\pm0\farcs05$ measured on the ACS images, we obtain a dynamical mass of $(4.4\pm1.1)\times10^5 \left(\frac{D}{2.2 {\rm Mpc}}\right) M_\odot$. This is in excellent agreement with the mass predicted by simple stellar population models for a cluster of this age and luminosity and a standard IMF. A correction for mass segregation would bring the two estimates into even closer agreement. Our analysis of structural parameters reveals a decrease in core radius with wavelength for NGC~1569-B. This trend is consistent with the colour gradients observed by \citet{hunter00} and by us in the HRC data when variations in the PSF are taken into account. A similar colour gradient was observed in an YMC in NGC~6946 \citep{laretal01}. The correlation between core radius and wavelength is also in the same sense as that found for M82-F by \citet{mgv05} and may be an indication that mass segregation (primordial or dynamical) is present in NGC~1569-B. However, this needs to be confirmed by a more detailed analysis, including a modelling of how mass segregation would translate into observables such as core- and half-light radius. As a final note, we find it fascinating to contemplate that only a couple of decades ago, it was still uncertain whether NGC~1569-A and NGC~1569-B were in fact star clusters in NGC~1569 or mere foreground stars. Today, the Advanced Camera for Surveys on HST has made it possible to not only settle this question definitively, but to study the individual stars that make up these clusters. | 7 | 10 | 0710.0547 |
0710 | 0710.0637_arXiv.txt | We combine high-resolution images in four optical/infra-red bands, obtained with the laser guide star adaptive optics system on the Keck Telescope and with the Hubble Space Telescope, to study the gravitational lens system \lensname (lens redshift~\zdvalue, source redshift~\zsvalue). We show that (under favorable observing conditions) ground-based images are comparable to those obtained with HST in terms of precision in the determination of the parameters of both the lens mass distribution and the background source. We also quantify the systematic errors associated with both the incomplete knowledge of the PSF, and the uncertain process of lens galaxy light removal, and find that similar accuracy can be achieved with Keck LGSAO as with HST. We then exploit this well-calibrated combination of optical and gravitational telescopes to perform a multi-wavelength study of the source galaxy at $0\farcs01$ effective resolution. We find the \sersic index to be indicative of a disk-like object, but the measured half-light radius (\re=\revalue) and stellar mass (\stellarmass=\stellarmassvalue) place it more than three sigma away from the local disk size-mass relation. The \lensname source has the characteristics of the most compact faint blue galaxies studied, and has comparable size and mass to dwarf early-type galaxies in the local universe. With the aid of gravitational telescopes to measure individual objects' brightness profiles to 10\% accuracy, the study of the high-redshift size-mass relation may be extended by an order of magnitude or more beyond existing surveys at the low-mass end, thus providing a new observational test of galaxy formation models. | \label{sect:intro} Galaxies do not appear in arbitrary combinations of luminosity, mass and shape, but instead obey empirical scaling relations (such as the Fundamental Plane for early-type galaxies). Explaining the origin, and cosmic evolution, of the scaling relations is a fundamental goal of galaxy formation theories. As far as disk galaxies are concerned, the hierarchical structure formation scenario predicts a correlation between size and stellar mass, with width depending on the distribution of the initial spin of the dark halos \citep{F+E80}. At any given mass, the expected distribution of sizes is well-approximated by a log-normal distribution. Qualitatively, this prediction is quite robust, although the exact forms of the correlation and the distribution depend on the details of baryonic processes such as energy feedback from star formation and bulge instability \citep{MMW98,She++03,Ton++06,Dut++07,S+B07}. Therefore, measuring the shape and width of the correlation provides not only a test of the standard paradigm, but also valuable information on the poorly-understood baryonic processes happening at sub-galactic scales. From an empirical point of view, the relation between size, luminosity (or equivalently surface brightness) and stellar mass is well established for disk galaxies in the local Universe \citep[\eg][]{She++03,Dri++05}. Analysis of suitable objects in the Sloan Digital Sky Survey shows that at any given mass (luminosity) the distribution of galaxies is indeed well-approximated as log-normal, although the scaling with mass of the characteristic size and the width of the distribution are non-trivial. Defining disk galaxies as those being well-fit by a single \sersic component with index $n<2.5$, \citet{She++03} find that above a characteristic stellar mass ($\log M_{*,0}/\msun\sim 10.6$ corresponding to approximately M$_{r,0}=-20.5$), size scales rapidly with stellar mass ($R\sim M_*^{0.39}$) and the scatter is relatively small ($\sigma_{\ln R}\sim 0.34$). Below the characteristic stellar mass the correlation flattens ($R\sim M_*^{0.14}$) and the scatter increases significantly ($\sigma_{\ln R}\sim 0.47$). At intermediate redshift ($0.1\lsim z \lsim 1$) the nature and interpretation of the size-luminosity or size-mass relation is more uncertain. Several authors \citep[\eg][]{Fer++04,Bar++05,Tru++06,Mel++06} have used Hubble Space Telescope images to determine the sizes of intermediate and high ($z \gsim 1$) redshift galaxies, down to the resolution and completeness limits of HST (roughly equivalent to 1~kpc and 10$^{10} \msun$). Recent studies conclude, taking selection effects into account, that there is significant evolution in the size-luminosity relation \citep{Bar++05,Tru++06,Mel++06}. However, it is hard to disentangle luminosity evolution from size evolution, to ensure that samples at different redshifts are directly comparable, and to compare results from different studies, as the selection criteria are often similar but not identical (\eg color vs. morphology; morphology determined via \sersic index vs.\ bulge to disk decomposition vs.\ concentration parameter vs. visual classification). Overall, it appears that disk galaxy evolution cannot be explained by pure luminosity or pure size evolution, but requires a combination of both. In contrast, the relation between size and stellar mass appears to have changed very little since $z\sim1$ \citep{Bar++05}, much less than would be expected in the naive model where stellar mass and size are proportional to the virial mass and radius (and hence size is expected to scale as $H(z)^{-\frac{2}{3}}$, where $H(z)$ is the Hubble parameter). Rather, galaxies appear to be growing ``inside-out'' in scale radius as their stellar mass increases such that the size-mass relation is preserved over cosmic time \citep{Bar++05}. It has been suggested that galaxy evolution models that take into account the ever-increasing concentration of dark matter halos, and the further effect of baryons via adiabatic contraction could provide the physics required to reproduce the observed trend \citep{Som++06}, although this may make it more difficult to reproduce simultaneously other scaling laws, for example the Tully-Fisher \citep{T+F77} relation \citep{Dut++07}. Lower mass ($M_* \lsim 10^{10} \msun$) galaxies are even less well understood. While the local size-mass relations of \citet{She++03} for low ($n<2.5$) and high ($n>2.5$) \sersic index objects diverge, the interpretation of \sersic index as a morphological galaxy classifier becomes more uncertain at lower masses \citep[e.g.\ ][]{CCD92,TBB04}. At the same time, the measurement of the structural parameters themselves becomes harder as the galaxy size decreases. Nevertheless, such small galaxies are important objects to understand: the luminous compact blue galaxies first noted by \citet{K+K88} appear in large numbers at intermediate redshifts in deep HST images \citep[\eg][]{Noe++06,Raw++07}, but evolve very rapidly to vanishing abundance in the local universe. What becomes of these objects, which represent sites of small-scale but vigorous star formation, is a topic of some debate, with dwarf spheroids \citep[\eg][]{Koo++94,Noe++06} and the bulges of disk galaxies \citep[\eg][]{Ham++01,Raw++07} the principle candidates. Gravitational lensing is a powerful tool with which to extend the investigation of scaling laws over cosmic time \citep[\eg][]{Tre07}. On the one hand, the lensing geometry provides a precise and almost model-independent measure of total mass of the lens galaxy. Since the lens galaxies are mostly early-type galaxies (or the bulges of spirals), this gives a new handle on the mass profile of these systems \citep{T+K04,Koo++06} and hence, for example, on the relationship between stellar and total mass \citep{Bol++07}. On the other hand, the background source is typically magnified by a factor of $\sim$10, mostly in the form of a stretch along the azimuthal direction. While lensing preserves surface brightness, the increase in apparent size of the lensed source means that the number of pixels at any one surface brightness also increases, such that the isophotes are observed at higher signal-to-noise. Thus, gravitational lenses act as natural telescopes, allowing one to gain a factor of $\sim10$ in sensitivity and spatial resolution, and thus improve markedly our ability to study the size and dynamical mass (through rotation curves) of intermediate and high redshift galaxies. For example, studies of the internal structure of faint blue galaxies~\citep{Ell97}, and in particular the most compact of these \citep{Koo++94}, are currently limited by the resolution of HST \citep{Phi++97}. When magnified by a gravitational lens, such objects become well-resolved. Thanks to the dedicated efforts of several groups, the number of known gravitational lenses is increasing dramatically: it is now possible to envision statistical studies of relatively large sample of lensing or lensed galaxies in the near future. In this paper we present multi-color high-resolution images of the gravitational lens system \lensname \citep{Bol++06}, obtained with both the Hubble Space Telescope and with the Laser Guide Star Adaptive Optics (LGSAO) System on Keck II. The scientific goal of the analysis of this case study is two-fold. First, we perform a detailed comparison of the results of the lens modeling across bands, showing that -- when a bright nearby star is available for tip-tilt correction and conditions are favorable -- the most important parameters can be measured with comparable accuracy with HST and Keck-LGSAO. Second, we exploit this particular cosmic telescope to achieve super-resolution of the source galaxy. \citep[See][ for Keck LGSAO observations of a lens with a point-like source.]{McK++07} With a lens magnification of $\mu \gsim 10$, the resolution of the HST and Keck images ($\sim 0\farcs1$ FWHM) corresponds to a physical scale of ($0.66{\rm kpc} / \mu \approx 0.05{\rm kpc}$) at the redshift of the source {\zs~=~\zsvalue}, comparable to the resolution attainable from the ground when studying galaxies in the Virgo Cluster in $1$~arcsec seeing. We derive the \sersic index, size, and stellar mass of the source, and show that using gravitational telescopes the size-mass relation may be extended by an order of magnitude in size with respect to current studies, thus allowing one to probe, for example, whether the change in slope and intrinsic scatter below the characteristic mass persists to higher redshifts. This paper is organized as follows. After describing the observations in section~\ref{sect:obs}, we outline in sections~\ref{sect:psf} and~\ref{sect:subtraction} two sources of systematic error and our strategies for dealing with them, before explaining our modeling methodology in section~\ref{sect:modelmethod}. In sections~\ref{sect:modelresults} and ~\ref{sect:source} we present our results, which are then discussed (section~\ref{sect:discuss}) before we draw conclusions in section~\ref{sect:conclude}. Throughout this paper magnitudes are given in the AB system. We assume a concordance cosmology with matter and dark energy density $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$, and Hubble constant H$_0$=70 kms$^{-1}$Mpc$^{-1}$. | \label{sect:conclude} We find that high quality images from \nirc are capable of providing very similar precision on simple lens and source model parameters to typical datasets from \acs and \nicmos. The data themselves contain information about the most appropriate PSF model to use, to the extent that a set of nearby unsaturated stars can be fruitfully compared using suitable statistics that are sensitive to the goodness-of-fit. We estimate that even for the LGSAO imaging this way of modeling the PSF allows a photometric precision of 0.05 mag. However, the calibration of isothermal However, the calibration of \emph{isothermal} galaxy-scale gravitational lenses as cosmic telescopes is very likely limited by the subtraction of the lens galaxy light. We estimate that this procedure introduces up to 0.1 magnitudes of systematic error into the source galaxy photometry. However, this is still smaller than the error introduced by the assumption of an isothermal density profile for the lens itself. With this in mind we draw the following conclusions about the source behind \lensname: \begin{itemize} \item Our photometry is robust enough to permit a reconstruction of the SED, and we find a stellar mass of (\stellarmassvalue). This is a factor of 5 smaller than the completeness limit of the GEMS disk galaxy analysis of \citet{Bar++05}, and also smaller than the least massive spheroid at this redshift studied by~\citep{McI++05}. \item The \sersic profile parameters of the source can be measured to high accuracy. We find an effective radius of (\revalue) ($\approx 0.09$ arcsec with $\sim10$\% accuracy), and a \sersic index of (\bluesindexvalue) in the \Iband ($\sim$ rest-frame~B), and that these values change little over the rest-frame optical range. \item This very small galaxy lies approximately 3-sigma below the local size-mass relation for disks. However, it shares the properties of the smallest of the compact narrow emission line galaxies of \citet{Koo++94}, and, despite its low \sersic index, is more typical of the dwarf early-type galaxies observed in the Virgo cluster~\citep{Fer++06} and the ``elliptical'' galaxies studied by~\citet{McI++05} at high redshift. \end{itemize} While the planned statistical analysis of a large sample of lensed galaxies will rely on the detailed understanding of the selection function, it is clear that the magnifying effect of gravitational lenses allows us to extend current size-mass studies to smaller sizes and lower masses than would otherwise be available, posing fresh challenges to models of galaxy formation and evolution. | 7 | 10 | 0710.0637 |
0710 | 0710.2318_arXiv.txt | % Proposed mechanisms for the formation of km-sized solid planetesimals face long-standing difficulties. Robust sticking mechanisms that would produce planetesimals by coagulation alone remain elusive. The gravitational collapse of smaller solids into planetesimals is opposed by stirring from turbulent gas. This proceeding describes recent works showing that ``particle feedback," the back-reaction of drag forces on the gas in protoplanetary disks, promotes particle clumping as seeds for gravitational collapse. The idealized streaming instability demonstrates the basic ability of feedback to generate particle overdensities. More detailed numerical simulations show that the particle overdensities produced in turbulent flows trigger gravitational collapse to planetesimals. We discuss surprising aspects of this work, including the large (super-Ceres) mass of the collapsing bound cluster, and the finding that MHD turbulence aids gravitational collapse. | Coagulation is the dominant mechanism for the growth of dust grains via van der Waals forces \citep{dt97} and for the growth of solid protoplanets by gravitational binding \citep{gls04}. But sticking is difficult in the ~mm--km size range, as confirmed by an extensive body of experimental work \citep{wb06}. Moreover, incremental growth of planetesimals leads to the rapid ($\sim10^2$ yr) inspiral of particles near a meter in size. An alternative hypothesis \citep[hereafter GW]{saf69, gw73} proposes that km-sized planetesimals formed from the gravitational collapse of smaller solids. This mechanism overcomes the sticking and radial drift obstacles in one fell swoop. However stirring by turbulent gas opposes gravitational collapse. GW noted that their dense particle midplane would trigger a turbulent boundary layer via Kelvin-Helmholtz instabilities. \citet{sw80} argued that this particle-driven turbulence (an example of drag force feedback) would stir up the particle midplane enough to prevent gravitational collapse. This appeared to be a fatal flaw of the GI hypothesis, since particle settling was a self-limiting process. \citet{sek98} and \citet{ys02} showed that particles could actually help their cause of becoming planetesimals. When the surface density of solids relative to gas is larger (by factors of few) than solar abundances, then vertical shear is no longer strong enough to overcome the anti-buoyancy of the dense midplane layer. The ability of particles to stir themselves is limited. However these works did not model the back reaction of drag forces on the gas in detail. Indeed most analyses of midplane Kelvin-Helmholtz instabilities \citep[with the exception of][]{jhk06} reduce the particle and gas dynamics to a simplified set of equations which omit relative motion and prevent a detailed examination of the effects of drag forces. Section \ref{sec:SI} describes the streaming instability \citep[hereafter YG]{yg05} which takes the opposite approach of neglecting stratification and vertical shear to consider two-way drag forces in a self-consistent, if incomplete, model. This idealized system shows that particle feedback triggers spontaneous particle clumping. The detailed 3D simulations of \citet[hereafter JOMKHY]{nature07} include vertical stratification, self-gravity, MHD turbulence and multiple particle sizes, see \S\ref{sec:KS} These models show that clumps of 15 -- 60 cm boulders, augmented by feedback effects, collapse gravitationally into bound clusters, which should continue to contract into planetesimals. | The long-standing mystery of planetesimal formation suddenly appears much less daunting, if no less fascinating. For decades, the obstacles seemed insurmountable: low sticking efficiencies, rapid radial migration and the inevitability of turbulent stirring. Now it is clear that turbulent stirring does not necessarily prevent gravitational collapse of solid particles into planetesimals. As discussed in \S\ref{sec:surprise}, more turbulence can promote gravitational collapse in some cases! The breakthoughs in particle-gas dynamics are related to the role of particle feedback. Once thought only to hinder particle settling by triggering Kelvin-Helmholz instabilities, we now realize that feedback triggers and augments particle clumping. The link between feedback and clumping has many pillars of support: the idealized streaming instability (YG, YJ, JY), Kelvin-Helmholz instabilities with uniform rotation \citep{jhk06}, and now 3D stratified models of Keplerian disk midplanes (JOMKHY and its supplement). Much work and many uncertainties remain. The large initial sizes assumed by JOMKHY might be unrealistic, either because coagulation stalls at smaller sizes, or if less violent gravitational collapse occurs first (see work on dissipative gravitational collapse by \citealp{war76,y05a}). The fact that most primitive undifferentiated meteorites betray no structures larger than mm-sized chondrules suggest this concern may be valid. However, particles which are small and tightly coupled to the gas disk are less effected by the dynamical clumping mechanisms discussed here, at least at the resolutions studied to date. Perhaps more surprises are needed. | 7 | 10 | 0710.2318 |
0710 | 0710.1631_arXiv.txt | We present a comprehensive study of accretion activity in the most underdense environments in the universe, the voids, based on the SDSS DR2 data. Based on investigations of multiple void regions, we show that Active Galactic Nuclei (AGN) are definitely common in voids, but that their occurrence rate and properties differ from those in walls. AGN are more common in voids than in walls, but only among moderately luminous and massive galaxies ($M_r < -20$, log $M_*/M_{\sun} < 10.5$), and this enhancement is more pronounced for the relatively weak accreting systems (i.e., $L_{\rm [O III]} < 10^{39}$ erg s$^{-1}$). Void AGN hosted by moderately massive and luminous galaxies are accreting at equal or lower rates than their wall counterparts, show lower levels of obscuration than in walls, and similarly aged stellar populations. The very few void AGN in massive bright hosts accrete more strongly, are more obscured, and are associated with younger stellar emission than wall AGN. These trends suggest that the accretion strength is connected to the availability of fuel supply, and that accretion and star-formation co-evolve and rely on the same source of fuel. Nearest neighbor statistics indicate that the weak accretion activity (LINER-like) usually detected in massive systems is not influenced by the local environment. However, H {\sc ii}s, Seyferts, and Transition objects are preferentially found among more grouped small scale structures, indicating that their activity is influenced by the rate at which galaxies interact with each other. These trends support a potential H{\sc~ii}$\rightarrow$Seyfert/Transition Object$\rightarrow$LINER evolutionary sequence that we show is apparent in many properties of actively line-emitting galaxies, in both voids and walls. The subtle differences between void and wall AGN might be explained by a longer, less disturbed duty cycle of these systems in voids. | The regions that are apparently devoid of galaxies \citep{kir81} and clusters \citep{ein80}, the voids, are arguably the best probes of the effect of the environment and cosmology on galaxy formation and evolution. If, as suggested by the standard cosmological paradigm, structure in the present-day universe formed through hierarchical clustering, with small structures merging to form progressively larger ones, galaxies in the currently most underdense regions must be the least ``evolved'' ones, as they must have formed at later times than those in the dense regions. Therefore, void and cluster galaxies must follow different evolutionary paths. Disturbing processes like stripping and harassment, that operate preferentially in crowded environments, should occur rarely in voids. Studies of the properties of the void galaxies, in contrast to those in relatively crowded regions, or walls, should provide some of the strongest constraints for distinguishing the intrinsic properties, which characterize a galaxy when it is first assembled, from properties that have been externally induced, over the whole history the universe: the ``nature versus nurture'' problem. Statistically significant conclusions regarding the distinctness of the void galaxies relative to those in denser regions, the ``walls'' hereafter, emerged only recently, with the advent of large surveys such as SDSS and 2dF. Such data, and in particular SDSS, offered for the first time the possibility to find and analyse both photometrically and spectroscopically, large samples of extremely low density regions (i.e., $\delta\rho/\rho < -0.6$ measured on a scale of $7 h^{-1}$ Mpc, Rojas et al. 2004, 2005), and allowed for accurate estimates of the void galaxy luminosity and mass functions \citep{hoy05, gol05}. These studies show that void galaxies are fainter, bluer, have surface brightness profiles more similar to those of late-type systems, and that their specific star formation rates are higher than those in denser regions. The mass and luminosity functions are found to be clearly shifted towards lower characteristic mass and fainter magnitudes ($M^*$). Moreover, the faint end slopes of the wall and void luminosity functions are very similar which suggests that voids are not dominated by an excess population of low-luminosity galaxies. Consistently, no significant excess in the amount of dark matter is apparent. This means that, although largely devoid of light, the most underdense regions conform to a galaxy formation picture which is clearly not strongly biased. All these peculiarities demonstrate that the cosmological evolution of void systems is different from that of those living in environments of average cosmic densities. Given the tight correlations between the mass of black holes ($M_{\rm BH}$) and the dispersions and the masses of the galactic bulges within which they reside \citep{mag98, fer00, geb00, mar03}, one would then expect that the growth of massive BHs in galaxy centers (and therefore the accretion process within active galactic nuclei, AGN), also differs among distinct environs. Extension of environmental studies of AGN properties to extreme regions like cosmic voids is thus crucial to understand the co-evolution of galaxies and their central BHs. Moreover, while there is general agreement that the growth of black holes must be closely related to galaxy assembly \citep{silk98, kau00, beg05}, there is no consensus as to how exactly accretion and star formation are coupled. The void galaxies could be, arguably, the best test-bed for understanding whether these processes are synchronized, or precede one another, and whether feedback from the actively growing BHs facilitates star formation (e.g., by dynamically compressing gas clouds through radio jets), or suppresses it (e.g., by blowing away the gas). To date, studies of the spectral properties of the void AGN remain limited to individual voids, e.g., the Bootes void \citep{kir81}, permitting the identification of only a few AGN among only a few dozen void galaxies \citep{cru02}. Quite surprisingly, such investigations find that the AGN fraction and their emission-line properties are similar in voids and in their field counterparts. Moreover, their associated stellar populations appear to share similar characteristics in the two extreme environs. The conclusions of these studies are based on small number statistics and do not however exclude the hypothesis that the void emission-line activity, whether originating in star-formation or accretion, could be connected with, e.g., filaments within voids; such structures would provide local environs similar to those in the field. The present SDSS samples of voids and void galaxies offer us for the first time the possibility to test and observationally constrain such ideas. It is important to note that previous investigations of the environmental dependence of nuclear activity in the relatively nearby universe do not reach the extreme spatial densities representative of voids. For example, in \citet{kau04}, the lowest density regions include over 25\% of the galaxies, which is more than 3 times more galaxies than the true void regions encompass. Their conclusions are interesting, and an extension of such an investigation at truly low densities is clearly desirable. In particular, it is important to quantify the degree to which the finding that, at fixed stellar mass, twice as many galaxies host strong-lined AGN in low-density regions than in high, extends to cosmic voids. Our work provides such an analysis. We employ in this work the most accurately classified samples of voids identified within SDSS to date, that yield $\sim 10^3$ void galaxies. Motivated by our recent study on the AGN clustering phenomena \citep{con06a}, which shows that there are differences in the large scale structure of active galaxies, and that their clustering amplitude correlates with their strength or rate of accretion, and possibly with the availability of fuel, we compare void and wall active galaxies of different types as classified based on their emission-line properties. Through such a comparison we aim to understand: 1) how the large and small-scale structures influence accretion onto their central black holes, and 2) to what degree AGN activity is triggered by interactions or mergers between galaxies. To answer these questions, we investigate the occurrence rate of different types of spectrally defined AGN, and how their accretion activity relates to their associated black hole mass, the mass and the age of their associated stellar populations, host morphology, brightness, and nearest-neighbor distance. We organize the paper as follows. In Section ~\ref{data} we present the void and wall sample selection, and the spectral classification we use in defining various types of actively emitting galaxies. We compare and discuss the AGN occurrence rate in voids and walls, both globally and at fixed host properties in Section ~\ref{fractions}. We examine in Section ~\ref{voidaccretion} the accretion rates, the fuel supply, and the properties of the associated star-formation in void and wall actively line emitting systems, while in Section ~\ref{nnprop} we discuss potential differences in their small scale environments. We summarize our findings in ~\ref{discussion}, and discuss the possible implications on the nature of the power sources in the low luminosity AGN, the AGN--host connection, and current models of galaxy formation. In particular, we show empirical evidence for a possible evolutionary sequence that links different types of strong line-emitting galaxies defined based on their spectral characteristics. Throughout this work, unless otherwise noted, we assume $\Omega_m = 0.3$, $\Omega_{\Lambda} = 0.7$, and $H_0 = 100h$ km s$^{-1}$ Mpc$^{-1}$. | \label{discussion} \subsection{\sc Summary of Results} \label{results} We use the largest sample of voids and void galaxies yet defined to investigate the galactic nuclear accretion phenomenon in the most underdense regions, in relation to their more populous counterparts. By employing spectroscopic and photometric data based on the SDSS DR2 catalog, and in particular measurements available in the Garching catalog, we conduct a comparative analysis of void and wall systems of different radiative signature. We find that all types of Low Luminosity AGN exist in void regions. However, their occurrence rate and intrinsic properties show variation from their wall counterparts. The differences between the wall and void AGN seem to be driven by the properties of their hosts, which are correlated with (or governed by) their small scale environment. Following is a summary of our main results: (i) Among moderately bright or fainter galaxies ($10 < log (M_*/M_{\sun}) < 10.5$, $M_r \ga -20$), the rate of occurrence of AGN is higher in voids than in walls. The most common accretion activity in voids is of medium power, with $L_{\rm [O III]} \sim 10^{38}$ erg s$^{-1}$. For the more luminous massive hosts, due to small number statistics, the relative prevalence of accretion activity in voids versus walls remains poorly constrained. (ii) The majority of void AGN, which are hosted by $M_r \ga -20$ galaxies, have the tendency to accrete at lower rates than in those in walls. This behavior seems related to the fact that the void systems show less obscuration and, perhaps, less dense emitting gas. That the stellar populations associated with void and wall AGN are similarly aged suggests that fuel might be equally available for accretion in void and wall galaxies of similar properties (i.e., $M_r$), but that fuel is less efficiently driven towards the nucleus in void galaxies. (iii) The few void AGN hosted by bright, massive galaxies ($M_r \la -20$, log $M_*/M_{\sun} > 10.5$) are LINERs that show peculiarly higher accretion rates, larger amounts of obscuring matter and more recent star-formation than in their wall counterparts. These particular systems reinforce the general trends other objects show: higher accretion rates are invariably associated with younger stellar populations and higher obscuration. These trends suggest that the amount of obscuration could be a measure of the available fuel for both star formation and accretion. (iv) The radio activity of line-emitting galaxies appears both less frequent and weaker in voids than in walls. Were we able to support these differences with statistically significant measurements, they would imply that central radio activity in wall systems, including H {\sc ii}s, is more pronounced because it builds on contributions from accretion that remains optically obscured and therefore undetected. (v) Nearest neighbor statistics show that the type of emission-line activity is correlated with the small-scale local environment. The star-forming regions (the H{\sc~ii}s), populate the most crowded sub-regions of voids while populating relatively sparse regions in walls; both are environments where low-mass galaxies recently formed. The weakly active galaxies (LINERs) live within the clusters in walls but the most rarefied regions in voids. This finding is puzzling and suggests that these systems, which are generally old, were probably not aware of their environments when they formed. Actively accreting systems (Seyferts and possibly the Ts) inhabit intermediate regions, which are relatively dense galaxy neighborhoods in voids but are of average density in walls. (vi) These correlations among the type and strength of galactic nuclear activity, incidence rates of different types, and their small and large scale environments, suggest an H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER evolutionary scenario in which interaction is responsible for propelling gas towards the galaxy centers, triggering star formation and feeding the active galactic nucleus. \subsection{\sc The H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER Evolutionary Sequence} \label{sequence} Figure ~\ref{h2stl} illustrates how various intrinsic and host properties of actively line-emitting galaxies follow this H {\sc ii}$\rightarrow$S/T$\rightarrow$LINER sequence. The early stages of such objects manifest themselves as H{\sc~ii}, as the accreting source remains heavily embedded in dust. As the star-burst fades in time, the dominance of the Seyfert-like excitation in systems of generally small but actively accreting black holes becomes more evident. Successive evolution reveals aging stellar populations associated with objects spectrally classified as Transition objects, that are still showing signs of accretion, followed by LINERs, whose stars are predominantly old and whose accretion onto already grown-up BHs is close to minimal. Note that this H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER progression is very similar in walls and voids. The lower accretion rates and the higher frequency of actively accreting systems in void versus wall galaxies of similar properties indicate a potential delay in the AGN dominance phase within voids. Thus, void AGN progress through the H {\sc~ii}$\rightarrow$S/T$\rightarrow$LINER sequence more slowly, while the sequence is similar for both void and wall galaxies. This picture fits well the observed properties of each type of galaxy nucleus: \begin{itemize} \item That H{\sc~ii}-type of activity is significantly more frequent in voids than in walls suggests that their void-like environments, in which they have closer (both 1st and 3rd) neighbors than other types of objects have, are essential in triggering their activity. In other words, close encounters that produce either harassment, or major and/or minor mergers may be an important cause for igniting both accretion and star formation. \item Seyferts' environments in both voids and walls are intermediate between those of H{\sc~ii}s and LINERs, regardless of their brightness. If the Seyfert-like activity is triggered by interactions, probably the same ones that turn on the H{\sc~ii}s, there must be a time lag between the onset of the star-burst and when accretion becomes dominant, or simply observable. Such a time interval corresponds to a period of aging of the stellar population (as seen in the differences in $D4000$ and H$\delta_{\rm A}$ between H{\sc~ii}s and Seyferts), when the post starburst fuel becomes increasingly available for accretion. Moreover, this progression develops relatively uninterrupted in voids, and therefore, possibly, at a slower pace than it would in walls where close encounters or other types of interactions can either accelerate or terminate it. \item Void and wall Ts are barely distinguishable in their physical characteristics. In both voids and walls, their nn-statistics and intrinsic and host properties are intermediate between those of LINERs and H{\sc~ii}s, for any given range of $M_r$. Their BHs are apparently growing (their accretion activity is stronger than that of Ls but weaker than that of Seyferts), their fuel supply seems plentiful (they are found in some of the most obscured systems), and they are associated with (quite massive) stellar populations that are generally younger than those of LINERs, but older than the majority of star-forming systems. It is among Ts however that the most massive void accreting BHs are observed; this might suggest that, in the proposed H {\sc~ii}$\rightarrow$S/T$\rightarrow$LINER sequence, massive void galaxies reach the low accretion rate (i.e., LINER) phase later than in walls. \item Whether Seyferts or Transition objects are first in this sequence remains ambiguous. Both their intrinsic properties and nn-statistics are very similar, and remain intermediate between those of H {\sc~ii}s and LINERs. The H$\alpha$/H$\beta$ Balmer decrements and the $d_{\rm 1nn}$ are the only parameters that show a ``jump'' in the otherwise smooth H{\sc~ii}$\rightarrow$S$\rightarrow$T$\rightarrow$LINER sequence manifested by other properties of these systems. Further investigations of these objects should address the differences between them in terms of a possible evolutionary progression. \item Although we do not provide any quantitative estimates of the time spent in or during the various phases we propose here, this this analysis shows that both void and wall galaxies follow the same cycle. The AGN evolution does not affect the gravitational environment. To the contrary, it is the environment that sets the time scale for evolution along such a sequence; it seems to take longer to march through the different phases in voids than in walls, but the physics is the same. The large scale clustering is consistent with this picture: LINERs are {\it now} more clustered because objects in dense regions underwent the H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER evolution more quickly; the higher rate of galaxy-galaxy interactions speeds up the way AGN proceed through the sequence. Hosts whose central regions are now in the H {\sc~ii} phase will always be less clustered than current LINERs. \end{itemize} Although far from being complete, this proposed evolutionary sequence is engaging and offers a comprehensive picture for the co-evolution of AGN and their host galaxies. The broad idea that mergers trigger star formation and that the AGN appears afterwards, in fact shutting off the star formation because of feedback, has been discussed previously in the literature. For example, N-body simulations by \citet{byr86} and \citet{her95} showed that interactions drive gas toward the nucleus and can produce intense star formation followed by an AGN. More recent state-of-the-art hydrodynamical models \citep{dimat05, spr05, hop05, hop06} show that during mergers, the BH accretion peaks considerably {\it after} the merger started, and {\it after} the star-formation rate has peaked. However, whether early bright quasars and later, dimmer AGN obey similar physics needs still to be addressed. The H{\sc~ii}$\rightarrow$S/T$\rightarrow$L sequence that this study reveals, based on the smooth alignment of several of their spectral properties, may be the first empirical evidence for an analogous duty cycle in high redshift bright systems and in nearby galaxies hosting weak quasar-like activity. This scenario can also accommodate the rather inconclusive findings regarding the role of mergers in activating AGN: their hosts do not show evidence for bars \citep{mul97, lai02} or disturbances caused by galaxy-galaxy interactions \citep{mal98}, exhibit morphologies very similar to those of field galaxies \citep{derob98a, derob98b}, and pair counting in both optical and IR remain inconclusive as possible excess of companions are sometimes found \citep{dah84, raf95}, but not always \citep{fue88, lau95}. Moreover, \citet{sch01, sch04} show that claims of evidence of interactions in the literature could be attributed to selection effects. The H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER cycle suggests that the majority of AGN might be detected only a certain period after the interaction, allowing time for the starburst to fade and for the BH accretion to gain strength. One might argue that that other forms of evolution could also exist for these emission-line galaxies, as opposed to this simple progression. We note that this proposed sequence does not imply that every HII galaxy at the present epoch must necessarily become a LINER; it is certainly possible that some systems go through H{\sc~ii}-only phases, or L-only phases, which might not be part of the larger progression. We would however like to emphasize that the timescales necessary to transform from one galaxy type to the next are quite reasonable. At a first glance, this is somewhat surprising given the relatively large range in BH masses ($2 \times 10^7 M_{\sun}$ for H{\sc~ii}s, to $2 \times 10^8 M_{\sun}$ for LINERs, inferred from $\sigma_*$'s assuming, e.g., Tremaine et al. 2002), and their significantly low accretion rates ($L/L_{\rm Edd} \la 0.005$; see Figure ~\ref{h2stl}, and note that log $L[{\rm O III}]/\sigma_*^4 = -1$ corresponds to approximately $L/L_{\rm Edd} = 0.05$, according to Kewley et al. 2006); apparently, for a canonical value for the accretion efficiency, i.e., 10\%, it takes approximately a Hubble time to e-fold in BH mass. The key is in the fact that the low luminosity AGN accrete inefficiently, as very little energy generated by accretion is radiated away (the optically thin cooling time of the gas is longer than the inflow time). Studies show that, in these cases, the radiative efficiency can be as low as $10^{-6}$ \citep{ree82, nar94, qua03}. With a (not a necessarily extreme) efficiency value of $10^{-5}$ then, the e-folding time in the BH mass is $\approx 4.5 \times 10^2/l$ years, and thus, as little as few Myrs for, e.g., $l = L/L_{\rm Edd} = 0.0005$. This (not necessarily extreme) example shows then that such an evolutionary scenario is actually physically possible. The peculiarity of LINER environments is a new and intriguing result and clearly shows that this type of activity is not controlled by its surrounding environment. Galaxies hosting LINERs could be associated with high initial density peaks in the dark matter distribution, which evolved subject to their large scale environment. Being generally massive systems to start with, LINERs' hosts would be prone to accreting material around them. In voids, this ``cleaning'' enterprise would contribute to emptying the already rarefied neighboring space, leaving little or insufficient material for future formation of (massive, bright) galaxies; they would thus end up in the most underdense void neighborhoods. In walls, and in particular within clusters, accretion of surrounding material would make a small difference as the matter density is higher. Simulations of dark matter halos and correlations with the properties of their inhabiting galaxies should be able to address these ideas. Further tests of this H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER evolutionary scenario are clearly needed. These tests require larger samples that allow separation into different morphological subsamples, and observables that parametrize the galaxy morphology better than the concentration index. Moreover, we need better constraints on the BH masses and consequently the Eddington rates for the H{\sc~ii}s in particular, or the late type galaxies in general. When available, analysis of such parameters would shed light on the assumed coevality of star-formation and Seyfert-like BH accretion in centers of galaxies, removing ambiguities regarding the initial stages of the H{\sc~ii}$\rightarrow$S/T$\rightarrow$LINER progression. \vspace{1cm} SPECIAL NOTE: During the review process of this paper, we became aware of another piece of work which introduces the same idea of an evolutionary sequence, ``from star formation via nuclear activity to quiescence`` \citep{sch07}. Interestingly, although their general approach, analysis, and samples used, are quite different from ours, the measurements on which the evidence for a ``star-forming, transition object, Seyfert, LINER'' sequence is built shares a common set of parameters with those used in our analysis: stellar velocity dispersions, ages of starbursts, and reddening. The global frame in which this evolutionary sequence in presented is however definitely different. While Schawinski et al. take a step forward toward understanding this possible time sequence among early-type galaxies and provide an in-depth analysis of the possible time scales involved in this picture, our paper offers a broader perspective on how such an evolutionary sequence constrains the galaxy evolution models as we provide important links to the environment. \vspace{1cm} | 7 | 10 | 0710.1631 |
0710 | 0710.3634_arXiv.txt | It is believed that quark matter can exist in neutron star interior if the baryon density is high enough. When there is a large isospin density, quark matter could be in a pion condensed phase. We compute neutrino emission from direct Urca processes in such a phase, particularly in the inhomogeneous Larkin-Ovchinnikov-Fulde-Ferrell (LOFF) states. The neutrino emissivity and specific heat are obtained, from which the cooling rate is estimated. | \label{introduction} The phase structure of quantum chromodynamics (QCD) is one of the most challenging problem in particle and nuclear physics. We schematically illustrate in Fig.\ \ref{phase} the phase diagrams in $T-\m_B$ and $T-\m_I$ plots, where $T$, $\m_B$ and $\m_I$ are the temperature, baryon and isospin (or equivalently electron) chemical potentials respectively. In the $T-\m_B$ diagram, the left panel of Fig.\ \ref{phase}, the hadronic phase locates at low $T$ and low $\m_B$ region and undergoes a phase transition or a crossover to the deconfined quark phase at certain critical temperature $T_c$ or baryon chemical potential $\m_{Bc}$ of the orders of $T_c\sim 200$ MeV or $\m_{Bc}\sim 1$ GeV. At very high temperature, the quark-gluon-plasma (QGP), made of free quarks and gluons, forms. At asymptotically high $\m_B$ but low $T$, the ground state of QCD is the color-flavor-locked (CFL) superconductor \cite{alford1999} where the condensation of quark pairs spontaneously breaks color and chiral symmetries. At intermediate $T$ and $\m_B$, although quarks and gluons are deconfined they are still strongly coupled. In this regime, many QCD phases are proposed in recent years, such as, at low $T$, two-flavor color superconductivity (2SC) \cite{Ruester:2004eg}, gapless 2SC (g2SC) \cite{shovkovy2003}, gapless CFL (gCFL)\cite{alford2004}, spin-1 color superconductor \cite{iwasaki1995,schafer2000,schmitt2002}, kaon condensation in the CFL phase \cite{kaon-condensation}, et al.. For reviews of color superconductivity, see, e.g. Ref. \cite{csc-reviews}. There may exist resonance states at intermediate $T$, such as strongly coupled QGP (sQGP) \cite{shuryak2006} at low $\m_B$ or the pseudo-gap phase at moderate $\m_B$ \cite{kitazawa2005}. In the $T-\m_I$ diagram, the right panel of Fig.\ \ref{phase}, the hadron phase is in the region with low $T$ and low $\m_I$ while the QGP phase locates at very high $T$. At low $T$, when $\m_I$ grows above the value of the pion mass $m_\p$, the ground state turns out to be a Bose-Einstein condensation (BEC) of pions, as $\m_I$ increases further and is larger than about 230 MeV, the pion BEC crossover smoothly into the BCS superfluid of quark-anti-quark pairs with the condensate $\langle \ub i\g_5 d\rangle$ or $\langle \db i\g_5u\rangle$ \cite{son2001}. At intermediate $T$, resonance states such as sQGP may occur at low $\m_I$ while the states with strong fluctuations of thermally excited mesons or Cooper pairs are possible at moderate $\m_I$. \begin{figure}[!htb] \begin{center} \includegraphics[width=6cm]{phase1.eps} \includegraphics[width=6cm]{phase2.eps} \caption{The schematic phase diagrams of QCD on $T-\m_B$ and $T-\m_I$ planes.} \label{phase} \end{center} \end{figure} In this paper we consider the quark matter cores in neutron stars. The neutrino emission from direct Urca processes $d\ra u+e^-+\nb, u+e^-\ra d+\n$ is the most efficient way of cooling in quark matter. Sch\"afer and Schwenzer summarized the neutrino emissivities and specific heats for a variety of color superconducting phases of quark matter \cite{schafer2004}. Due to beta equilibrium the isospin chemical potential $\m _I$ is nonzero, quark matter could be a pion BEC ($\m_I<230$ MeV) or a BCS superfluid ($\m_I > 230$ MeV) when $\m _I > m_\p$ \cite{son2001}. On the other hand, the baryon chemical potential is still large (in the order of 1 GeV) and makes a big mismatch between the fermi surfaces of the pairing quarks $u(\bar u)$ and $\bar d(d)$, thus the BEC or BCS state is gapless. Such a gapless phase is stable in the BEC region but unstable in the BCS one with respect to the formation of nonzero LOFF momentum. In a previous work \cite{huang2007}, we have studied the neutrino emissivity and cooling rate due to Urca processes for the gapless pion condensed quark matter in the BEC region. In this paper we study the neutrino emission in the LOFF phase. We work in two-flavor case in the moderate baryon density, where the role of strange quarks is not important. Our units are $\hbar=k_B=c=1$ except particular specifications. As a convention, we denote a 4-momentum as $K^\mu =(k_0,\mathbf{k})$, and its 3-momentum magnitude as $k=|\mathbf{k}|$. \section {Quark Propagator\label{quark}} Our starting point is the two flavor Nambu-Jona-Lasinio Lagrangian of QCD \begin{eqnarray} \label{NJL} \cl=\jb(i\g^\m\pt_\m+\m\g_0-m_0)\j+g[(\jb\j)^2+(\jb i\vec{\t}\g_5\j)^2], \end{eqnarray} where $\j=(u,d)^{\rm T}$ is the quark fields, $g$ is the coupling constant and $\vec{\t}$ is the Pauli matrices. We have introduced the chemical potential matrix in flavor space, $\m=\mathrm{diag}(\m_u, \m_d)=(\m+\d\m, \m-\d\m)=(\m_B/3+\m_I/2, \m_B/3-\m_I/2)$ with $\m_B, \m_I$ the baryon and isospin chemical potential respectively. We assume that the $\b$-equilibrium is reached $\m_d=\m_u+\m_e$, which gives $\m_I=-\m_e$. In the chiral limit and without the chemical potentials, the Lagrangian (\ref{NJL}) respects the symmetry $U_B(1)\otimes SU_V(2)\otimes SU_A(2)$ corresponding to the baryon number, isospin vector and pseudovector conservation respectively. However, the presence of the chemical potentials explicitly break the isospin symmetry down to $U_V(1)$ with the conserved quantum number of $\t_3$, and chiral symmetry down to $U_A(1)$ with the conserved quantum number of $i\g_5\t_3$. By introducing the chiral condensate $\s=-2g\lan\jb\j\ran$ and pion condensates $\p^-=\D e^{-2i\bl\cdot\bx}=4g\langle\ub i\g_5 d\ran$, $\pi^+=\D e^{2i\bl\cdot\bx}=4g\langle\db i\g_5 u\ran$, we arrive at the mean field Lagrangian \begin{eqnarray} \cl_\mf=\jb\lb\begin{array}{cc}i\g^\m \pt_\m+\m_u\g_0-m & \p^+i\g_5 \\ \p^-i\g_5 & i\g^\m \pt_\m+\m_d\g_0-m \end{array}\rb\j -\frac{\s^2+\D^2}{4g}, \end{eqnarray} with the effective quark mass $m=m_0+\s$. There are mismatches in Fermi surfaces of anti-u and d quarks or anti-d and u quarks by the baryon chemical potential, hence pion condensates with nonzero total momentum or the LOFF states may be favored. The formation of the condensate $\s\sim\lan\jb\j\ran$ breaks the $U_A(1)$ chiral symmetry spontaneously with the Goldstone boson $\p_0\sim\jb i\g_5\t_3\j$, and that of pion condensates $\D\sim \langle\db i\g_5u\ran\sim \langle\ub i\g_5d\ran$ break the $U_V(1)$ isospin symmetry spontaneously. The translational and rotational symmetries are spontaneously broken by nonzero LOFF momentum $\bl$. The partition function of the system can be written as a functional integral \begin{equation} Z=\int [d\jb ] [d\j ] \exp{\left( \int_0^\b d\t\int d^3\bx\cl_\mf\right) } . \end{equation} Rewriting the quark fields via a gauge transformation which leaves the partition function unchanged, $\c_u(x)=u(x)e^{-i\bl\cdot\bx}$, $\c_d(x)= d(x)e^{i\bl\cdot\bx}$ (we still call $\c _{u,d}$ quark fields), the inverse quark propagator in flavor and momentum space reads \begin{eqnarray} \label{s-inverse} S^{-1}(K)=\lb\begin{array}{cc}\g^\m K_\m-\bl\cdot\g+\m_u\g_0-m & \D i\g_5 \\ \D i\g_5 & \g^\m K_\m+\bl\cdot\g+\m_d\g_0-m \end{array}\rb. \end{eqnarray} Note that $K^\m$ represents the 4-momentum of the $\c$ fields instead of the $\j$ fields. The propagator is written as \begin{eqnarray} \label{propagator} S(K)=\lb\begin{array}{cc} S_{uu}(K) & S_{ud}(K) \\ S_{du}(K) & S_{dd}(K) \end{array}\rb. \end{eqnarray} A straightforward calculation from Eq. (\ref{s-inverse}) gives the four elements \begin{eqnarray} S_{uu}(K)&=&\frac{(\g^\m K_{+\m}+m)(K_-^2-m^2-\D^2)+2\D^2\g^\m l_\m}{(K_+^2-m^2+\D^2)(K_-^2-m^2+\D^2)-\D^2[(K_++K_-)^2-4m^2]},\non S_{dd}(K)&=&\frac{(\g^\m K_{-\m}+m)(K_+^2-m^2-\D^2)-2\D^2\g^\m l_\m}{(K_+^2-m^2+\D^2)(K_-^2-m^2+\D^2)-\D^2[(K_++K_-)^2-4m^2]},\non S_{ud}(K)&=&\frac{(\g^\m K_{+\m}+m)(K_-^2-m^2-\D^2)+2\D^2\g^\m l_\m}{(K_+^2-m^2+\D^2)(K_-^2-m^2+\D^2)-\D^2[(K_++K_-)^2-4m^2]}\non &&\times \frac{\g^\m K_{-\m}-m}{K_-^2-m^2}i\g^5\D,\non S_{du}(K)&=&\frac{(\g^\m K_{-\m}+m)(K_+^2-m^2-\D^2)-2\D^2\g^\m l_\m}{(K_+^2-m^2+\D^2)(K_-^2-m^2+\D^2)-\D^2[(K_++K_-)^2-4m^2]}\non &&\times \frac{\g^\m K_{+\m}-m}{K_+^2-m^2}i\g^5\D,\non \end{eqnarray} where $K_\pm^\m=(k_0+\m\pm\d\m,\bk\pm\bl)$ and $l^\m=(\d\m, \bl)$. The excitation spectra of the quasi-particles can be obtained by solving the equation $\det S^{-1}(k_0,\bk)=0$ or equivalently the roots of the denominator of the propagator for $k_0$, \begin{eqnarray} 0&=&(K_+^2-m^2+\D^2)(K_-^2-m^2+\D^2)-\D^2[(K_++K_-)^2-4m^2]\non &\approx&[(k_0+\m+\ve^-_{\bk,\bl})^2-(\ve_{\bk,\bl}^+ +\d\m)^2 -\D^2]\non &&\times [(k_0+\m-\ve_{\bk,\bl}^-)^2-(\ve_{\bk,\bl}^+-\d\m)^2-\D^2], \end{eqnarray} where $\ve_{\bk,\bl}^\pm=(E_{\bk+\bl}\pm E_{\bk-\bl})/2$ with $E_\bk\equiv\sqrt{\bk^2+m^2}$. To arrive at the last line, we have taken the assumption that both $\D$ and $\bl$ are small comparing to the quark Fermi momenta in the LOFF phase. We have four excitation branches, $E_r^a(\bk,\bl)=-r\sqrt{(\ve_{\bk,\bl}^+ - a\d\m)^2+\D^2}-(\m - a\ve_{\bk,\bl}^-)$, with $a,r=\pm$. Taking the same approximation to the numerators in the elements of the quark propagator, we neglect the terms proportional to $\D^2\bl$. Thus we rewrite the elements of the quark propagator as, \begin{eqnarray} \label{suu-sdd} S_{uu}(K)&\simeq&\sum_{a,r=\pm}\frac{B_{r}^{a}(\bk ,\bl )\L^{a}_{\bk+\bl }\g_0}{k_0-E_r^a(\bk,\bl)},\non S_{dd}(K)&\simeq&\sum_{a,r=\pm}\frac{B_{-r}^{a}(\bk ,\bl )\L^{-a}_{\bk -\bl }\g_0}{k_0-E_r^a(\bk,\bl)}, \end{eqnarray} where we have introduced the energy projectors, \begin{equation} \L^a_{\bk }=\frac{1}{2}\left[1+a\frac{\g_0(\g\cdot\bk+m)}{E_\bk}\right], \end{equation} and the Bogoliubov coefficients, \begin{equation} \label{bog-coeff} B^{a}_r(\bk,\bl)=\frac 12 \left[1-ar\frac{\varepsilon ^+_{\bk,\bl}-a\delta\mu} {\sqrt{(\varepsilon ^+_{\bk,\bl}-a\delta\mu)^2+\Delta^2 }}\right], \end{equation} | 7 | 10 | 0710.3634 |
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0710 | 0710.3896_arXiv.txt | We analyze the mean rest-frame ultraviolet (UV) spectrum of Type Ia Supernovae (SNe) and its dispersion using high signal-to-noise Keck-I/LRIS-B spectroscopy for a sample of 36 events at intermediate redshift ($\overline{z}$=0.5) discovered by the Canada-France-Hawaii Telescope Supernova Legacy Survey (SNLS). We introduce a new method for removing host galaxy contamination in our spectra, exploiting the comprehensive photometric coverage of the SNLS SNe and their host galaxies, thereby providing the first quantitative view of the UV spectral properties of a large sample of distant SNe~Ia. Although the mean SN~Ia spectrum has not evolved significantly over the past 40\% of cosmic history, precise evolutionary constraints are limited by the absence of a comparable sample of high quality local spectra. The mean UV spectrum of our $z\simeq$0.5 SNe~Ia and its dispersion is tabulated for use in future applications. Within the high-redshift sample, we discover significant UV spectral variations and exclude dust extinction as the primary cause by examining trends with the optical SN color. Although progenitor metallicity may drive some of these trends, the variations we see are much larger than predicted in recent models and do not follow expected patterns. An interesting new result is a variation seen in the wavelength of selected UV features with phase. We also demonstrate systematic differences in the SN~Ia spectral features with SN lightcurve width in both the UV and the optical. We show that these intrinsic variations could represent a statistical limitation in the future use of high-redshift SNe~Ia for precision cosmology. We conclude that further detailed studies are needed, both locally and at moderate redshift where the rest-frame UV can be studied precisely, in order that future missions can confidently be planned to fully exploit SNe~Ia as cosmological probes. | \label{sec:introduction} Supernovae of Type Ia (SNe~Ia) are now well-established as cosmological distance indicators. In addition to the original surveys by the Supernova Cosmology Project \citep[SCP;][]{1997ApJ...483..565P,1999ApJ...517..565P} and the High-Z Supernova Search Team \citep{1998ApJ...507...46S,1998AJ....116.1009R}, a new generation of SN~Ia surveys is underway both locally \citep{2002SPIE.4836...61A,2005coex.conf..525L,2006PASP..118....2H} and at higher redshifts \citep{2006A&A...447...31A,2007ApJ...659...98R,2007ApJ...666..694W}. Despite the availability of independent probes of the presence and properties of dark energy from studies of the cosmic microwave background \citep{2007ApJS..170..377S} and galaxy redshift surveys \citep{2002MNRAS.330L..29E,2005MNRAS.362..505C,2005ApJ...633..560E}, the luminosity distance--redshift relation for SNe~Ia provides the only {\it direct} evidence for a cosmic acceleration. The detection and characterization of dark energy, via measurements of the average cosmic equation of state parameter $<$$w$$>$, requires the precision measurement of SNe~Ia to redshifts $z\simeq$0.5--1, sampling the epoch of cosmic acceleration \citep{2006A&A...447...31A}. However, more precise constraints on the nature of dark energy, for example evidence for any variation in $w$ with redshift, requires extending these studies to redshift $z$$>$1 \citep{2004ApJ...607..665R,2007ApJ...659...98R} where the early effects of deceleration may be detectable. As projects are developed which plan to probe SNe~Ia beyond $z$=1 for this purpose \citep[e.g][]{2005NewAR..49..346A,2006SPIE.6265E..67B}, it becomes important to understand the possible limitations of using SNe~Ia as distance probes. Key issues relating to the diversity of SNe~Ia and their possible evolution with redshift as a population, together with the limiting effects of dust and/or color corrections to their photometric properties, are particularly crucial to understand. Several local studies \citep{1995AJ....109....1H,1999AJ....117..707R, 2000AJ....120.1479H,2001ApJ...554L.193H, 2005A&A...433..807M,2005ApJ...634..210G} have already indicated correlations between SN~Ia properties and host galaxy morphologies. More recently, \citet{2006ApJ...648..868S} have shown that the properties of distant SNe~Ia appear to be a direct function of their local stellar population, with the distribution of light curve widths and hence peak luminosities correlating with the host galaxy specific star-formation rate. This work also determined that the rate of SNe~Ia per unit stellar mass of their host galaxies is larger in actively star-forming galaxies, suggesting many must be produced quite rapidly in recently formed stellar populations, perhaps suggestive of more than one progenitor mechanism. The authors conclude that SNe~Ia may well be a bimodal or a more complex population of events \citep[see also][]{2005ApJ...629L..85S,2006MNRAS.370..773M}. Such diversity in the properties of SNe~Ia could have far-reaching implications, particularly if the {\it mix} of mechanisms or delay times within the broad population gradually changes with look-back time \citep[e.g.][]{2006MNRAS.370..773M,2006ApJ...648..868S,2007astro.ph..1912H}. These recent developments, which illustrate how improved precision reveals new physical correlations in the SN~Ia population, raise the broader question of whether future SN~Ia experiments might be limited in precision by variations of a systematic nature within the population, for example with redshift, which cannot be removed via empirical correlations. Detailed local surveys such as the LOSS/KAIT \citep{2005coex.conf..525L} and CfA surveys \citep{1999AJ....117..707R,2006AJ....131..527J} have presented valuable data on the homogeneity and trends in the SNe~Ia population. Further promising work is being undertaken via the Supernova Factory \citep{2002SPIE.4836...61A} and the Carnegie Supernova Project \citep{2006PASP..118....2H}. Important though these continued programs will be, they are insufficient to address all possible concerns about the use of SNe~Ia as precision tools in cosmology. Comparable studies at intermediate redshift\footnote{Defined here to represent the range 0.2$<$$z$$<$0.7.} will be particularly important in order to address questions relating to possible evolutionary effects and environmental dependencies. In addition, it is not always practical at low redshift to cover the full wavelength range necessary to test for systematic trends. In this paper we analyze high signal-to-noise ratio rest-frame ultraviolet (UV) spectra of a large sample of intermediate redshift SNe~Ia drawn from the Canada-France-Hawaii Telescope Supernova Legacy Survey \citep[SNLS;][]{2006A&A...447...31A}, a rolling search which is particularly effective for locating and studying events prior to their maximum light. Our aim is obtain a substantially higher signal-to-noise in the spectra than that typically obtained during spectroscopic programs to type SNe and measure redshifts. We target the UV wavelength region because in this wavelength region the SN spectrum is thought to provide the most sensitive probe of {\it progenitor metallicity} \citep[e.g.][]{1998ApJ...495..617H,2000ApJ...530..966L}, a variable which may shed light on the possibility of progenitor evolution. The time-dependent UV spectrum is also needed in estimating ``cross-band'' $k$-corrections, particularly at redshifts $z>1$ where optical bandpasses probe the rest-frame near-UV \citep{2004ApJ...607..665R,2007ApJ...659...98R}. Little is known about the properties and homogeneity of the UV spectra of SNe~Ia, largely because of the absence of suitable instruments for studying this wavelength range in local events. Although some local SN~Ia UV spectra are available from International Ultraviolet Explorer (IUE) or Hubble Space Telescope (HST) satellite data \citep[e.g.][]{1991ApJ...371L..23L,1993ApJ...415..589K,1995ESASP1189.....C}, the bulk of the progress now possible in this area can be provided from optical studies of intermediate-redshift events with large ground-based telescopes. The goals of this paper are thus to address the question of both the diversity and possible physical evolution in the intermediate redshift SN~Ia family. We compare the rest-frame UV behavior of local SNe~Ia with that derived for intermediate-redshift ($z\simeq$0.5) events where the rest-frame UV enters the region of high efficiency of the Keck LRIS-B spectrograph. We also study the degree to which the UV spectra at intermediate redshift represent a homogeneous population, independent of other variables such as the physical environment and light curve stretch. A plan of the paper follows. In $\S$~\ref{sec:selection-cfhtls-sne} we introduce the salient features of the SNLS and our method for selecting SNe~Ia for detailed study. In $\S$~\ref{sec:keck-observations} we discuss the Keck spectroscopic observations and their reduction, including the treatment of host galaxy subtraction and flux calibration. In $\S$~\ref{sec:analyses} we consider our sample with respect to the broader set of SNe found by SNLS, ensuring it is a representative subset in terms of various observables, and discuss existing local UV spectra. In $\S$~\ref{sec:results}, we undertake the detailed analysis. First we compare the UV spectra found in our sample with those found locally. We then examine the diversity of intermediate-redshift SNe~Ia in various ways and correlate the UV variations with the light curves of the SNe and the properties of the host galaxies. We discuss these trends in terms of progenitor mechanisms in $\S$~\ref{sec:discussion} and examine the implications in terms of possible long term limitations of SNe~Ia as probes of dark energy. We also present the mean phase-dependent SN Ia spectrum and its uncertainties for use in future work. | \label{sec: conclusions} We summarize our findings as follows: \begin{enumerate} \item{} We have secured high signal-to-noise ratio Keck spectra for a sample of 36 intermediate redshift SNe~Ia, observed at various phases, spanning the redshift range 0.15$<z<$0.7, and drawn from the Supernova Legacy Survey (SNLS). We demonstrate via inspection of the SN properties that our Keck sample is a reasonably fair subset of the larger sample of distant SNe~Ia being studied by the SNLS. \item{} We develop a new method for removing host galaxy contamination from our spectra based on measures of the galaxy and SN photometry. These refinements to traditional spectral reduction techniques allow us to achieve host-galaxy subtracted and flux-calibrated rest-frame spectra of high quality, extending down to rest-frame wavelengths of 2900\AA. \item{} Although no strong evidence is found for spectral evolution in the mean early-phase and maximum light spectra, when compared to local data, such evolutionary tests are hampered by the paucity of quality data at low redshift and a significant scatter in the spectra shortward of 4000\AA. We find no evidence for evolution internal to our data. We argue that the well-used local UV spectral template \citep{2002PASP..114..803N} is likely to be less representative than the mean spectrum compiled from the Keck data which we tabulate with the measured dispersion for use in future cosmological applications. \item{} Our principal finding is a large scatter from one SN to the next in the rest-frame UV spectrum even after standard dust corrections are made. By constructing various photometric bandpasses that avoid uncertainties arising from differential $k$-corrections associated with the range of redshifts in our sample, we show that while we can reproduce the stretch-dependent trends seen locally at 3500-4000\AA, the scatter at 3000-3400\AA\ is 3-5 times larger. \item{} Although progenitor metallicity may drive some of the trends seen in the Keck data, the UV variations are much larger than in contemporary models which span the expected metallicity range. Moreover, the UV spectrum also changes with phase in a manner which is not consistent with models. We conclude there are significant variations in the UV properties of SNe~Ia which are not accounted for by either the presently-employed empirical trends or the available SN~Ia models. \item{} As an illustration of the importance of understanding these new results, we calculate the error arising from the use of a single UV spectral template for calculating the cross-color $k$ correction, a correction essential for constructing the SN~Ia Hubble diagram as a probe of the expansion history. The dispersion arising from our UV spectra, if not randomly distributed along the Hubble diagram, presents an uncertainty 2-3 times larger than than would be necessary for recovering the equation of state parameter $w$ to 5\% using SNe~Ia at $z\simeq$1. We conclude that further detailed studies are essential if SNe~Ia are to be useful for precision measures of dark energy. \end{enumerate} | 7 | 10 | 0710.3896 |
0710 | 0710.1541_arXiv.txt | By considering simple, but representative, models of brane inflation from a single brane-antibrane pair in the slow roll regime, we provide constraints on the parameters of the theory imposed by measurements of the CMB anisotropies by WMAP including a cosmic string component. We find that inclusion of the string component is critical in constraining parameters. In the most general model studied, which includes an inflaton mass term, as well as the brane-antibrane attraction, values $n_{\rm s}<1.02$ are compatible with the data at $95\%$ confidence level. We are also able to constrain the volume of the warped throat region (modulo factors dependent on the warp factor) and the value of the inflaton field to be $<0.66M_{\rm P}$ at horizon exit. We also investigate models with a mass term. These observational considerations suggest that such models have $r< 2\times 10^{-5}$, which can only be circumvented in the fast roll regime, or by increasing the number of antibranes. Such a value of $r$ would not be detectable in any CMB polarization experiment likely in the near future, but the B-mode signal from the cosmic strings could be detectable. We present forecasts of what a similar analysis using PLANCK data would yield and find that it should be possible to rule out $G\mu > 6.5\times 10^{-8}$ using just the TT, TE and EE power spectra. | The inflationary paradigm is strongly supported by observations of the cosmic microwave background (CMB) made by the COBE and WMAP satellites~\cite{Smoot:1992td,Spergel:2006hy}. However, it is still a paradigm in search of a specific model based on fundamental physics. Brane inflation~\cite{Dvali:1998pa} has emerged as one of the most popular ways of embedding inflation within string theory. It uses a natural candidate for the inflaton: the field which describes the brane-antibrane separation. This field has a non-trivial potential due to the attractive brane-antibrane interaction which is flattened by the effect of the compactification of the extra dimensions \cite{Burgess:2001fx}, and the geometry of the branes \cite{GarciaBellido:2001ky}. Cosmic strings are also a natural occurrence within this model and are result of inhomogeneities in the tachyon field of the brane-antibrane pair \cite{C_Strings1,C_Strings2}. These are represented by a complex field with a non-trivial potential which supports the formation of codimension 2 defects which have been shown to exhibit the correct properties of lower dimension branes \cite{Tachyon_Strings}. Since reheating of the universe in these model proceeds via the annihilation of the brane-antibrane pair, the formation of cosmic strings is expected at the end of the inflation and the strings will be naturally located at the bottom of the throat as their tension is also warp-factor dependent. The most complete model of brane inflation incorporating moduli stabilization has been proposed in ref.~\cite{Kachru:2003sx}. In this model inflation happens naturally as a mobile brane falls down the warped throat, being attracted by an antibrane stuck at the bottom of the throat. Antibranes have a warp-factor dependent potential, and therefore minimize their energy by moving to the bottom of the throat, the region of strongest warping. The mobile brane is not affected by the warping. One can understand this by considering the warped geometry as being generated by a large stack of branes. An antibrane is attracted to the stack and will move towards it, that is towards largest warping. The brane feels no force from the stack of branes (it is BPS with respect to them), only from the antibrane located at the bottom of the throat. This situation can be described by a simple scalar field, since in this model all moduli are stabilized and only the brane-antibrane separation is evolving. Constraints on the cosmic string tension, $G\mu$, come from a variety of observations. Of most interest here are those which are most robust. In particular, we will concentrate on the constraints which result from their inclusion as a sub-dominant component in the angular power spectrum of anisotropies of the cosmic microwave background (CMB)~\cite{CHMb,WB,Wyman:2005tu,Seljak:2006bg,Bevis:2006mj,Battye:2006pk,Bevis:2007gh}. As pointed out recently~\cite{Battye:2006pk,Bevis:2007gh} in the context of the third year WMAP data, larger values of the spectral index of density fluctuations, $n_{\rm s}$, are compatible with observations if a sub-dominant string component with around $5-10\%$ of the large-scale amplitude is included. This is a generic feature of any inflation model which produce strings. In ref.~\cite{Battye:2006pk} accurate constraints on the coupling constant, $\kappa$ and mass scale, $M$, relevant to the simplest models of supersymmetric F- and D-term hybrid inflation were computed, taking into account the fact that the observed power spectrum is described by 3 parameters, $G\mu$, $n_{\rm s}$ and the power spectrum amplitude, $P_{\cal R}$, each of which can be derived from $\kappa$ and $M$. In this model dependent approach more powerful constraints are possible. We will adapt the same approach, where relevant to the case of brane inflation in this paper. | \label{sec:discussion} In this paper we have presented limits and constraints on the parameters (and derived parameters) of brane inflation models in the slow-roll regime when a cosmic string component is included in the fitting. Important aspects of the results are an increased range of acceptable values of $n_{\rm s}$, limits on $\gamma$, which is related to the volume of the internal space, $\beta$ the inflaton mass parameter and a constraint on the value of $\phi_{\rm e}<M_{\rm P}$ irrespective of any theoretical considerations with the consequent implications for $r$. We note that one way out of the bounds which we have discussed here is to consider the large $\beta$ regime, where there are again viable inflationary models~\cite{Bean:2007hc}. These models fall in the category of inflationary models with general speed of sound studied in~\cite{Peiris:2007gz}. In this case fast-roll inflation is possible due to the kinetic term actually being of the Dirac-Born-Infeld type (DBI)~\cite{Silverstein:2003hf,Alishahiha:2004eh,Kecskemeti:2006cg}. This modifies some of the preceding discussion. The Lagrangian for DBI inflation is\be {\mathcal L} = \frac{1}{f\left(\phi\right)}\left[1-\sqrt{1- f\left(\phi\right)g^{\mu\nu}\partial_{\mu}\phi\partial_{\nu}\phi}\right] - V\left(\phi\right)\,, \ee where the function $f\left(\phi\right)$ depends on the throat geometry and the potential $V(\phi)$ would be given by (\ref{masspot}). In this case, the expression for the scalar-to-tensor ratio is modified to $r=16c_{\rm s}\epsilon$ where the parameter $\epsilon$ is defined as a generalized slow-roll parameter $\epsilon = 2c_{\rm s}M_{\rm P}^2(H^{\prime}/H)^2$ and the speed of sound is \be c_{\rm s}^{2} = 1- f\left(\phi\right)g^{\mu\nu}\partial_{\mu}\phi\partial_{\nu}\phi\,. \ee Since $c_{\rm s}^2\le 1$ one might think that $r$ would be small, however $\epsilon$ can be much larger in fast roll models. This was emphasized in ref.~\cite{Bean:2007hc} who also pointed out in the case where $r$ is small the non-Gaussianity of the density fluctuations, quantified by $f_{\rm NL}$, would be particularly high due to consistency relation~\cite{Lidsey:2006ia} \be 1-n_{\rm s}=0.4r\sqrt{f_{\rm NL}}\,, \ee which was derived in the equilateral triangle limit in momentum space. As $\beta$ increases in the slow-roll $\text{regime}^{\footnotemark[1]}$ \footnotetext[1]{Simultaneously with our work, Ref.\cite{Bean:2007eh} appeared where an alternative model, the so-called infrared DBI brane inflation, is compared with observations. For this class of models an even stronger bound on the cosmic string tension is found, $G\mu < 10^{-14}$.} the value of $n_{\rm s}$ increases beyond that which is compatible with the data. However, at some value of $\beta$ the effects of the DBI kinetic term kick in and the potential (\ref{masspot}) becomes dominated by the term $\propto\beta$. Hence $P_{\cal R}$ and $n_{\rm s}$ are those of simple Klein-Gordon field, albeit modified by $c_{\rm s}$. This breaks the link between $G\mu$ and $r$, and therefore the constraints discussed in the previous sections only apply in the slow roll regime. We note that all the brane inflation models constructed so far require some amount of fine-tuning to work. In the present case, the tuning corresponds to the value of $\beta$ being low. The effects of moduli stabilization have led to a large value of the $\eta$ parameter, making fine-tuning necessary in models of F-term inflation~\cite{Kachru:2003sx,McAllister:2005mq}. D-term inflation models do not suffer form the same $\eta$-problem but threshold corrections to the superpotential \cite{Berg:2004ek,Berg:2005ja} generate a large mass for the inflaton, making fine-tuning necessary in these models as well. A possible way to alleviate the fine-tuning problem is made possible by a remarkable property of a class of multi-brane models~\cite{Cline:2005ty}. Usually one balances the effect of the volume stabilization mechanism against the Coulombic brane-antibrane interaction to obtain a potential with a flat enough region to support inflation~\cite{Burgess:2004kv}. The model requires a fine balancing of the two effects. For generic values of the parameters the potential will either be too steep, or it will feature a local de-Sitter minimum where the mobile branes being separated from the antibranes by a potential barrier. However, if one or more branes tunnel out of the local minimum and annihilates with the antibranes, the height of the barrier decreases, and for a critical number of branes it disappears, resulting in a monotonic but almost flat potential. Inflation then proceeds via slow-roll as the remaining branes roll towards the antibranes stuck at the bottom of the warped throat. We will investigate this type of model using the techniques applied here in a future study. A similar, but complementary study, to ours for the case $\beta=0$ but not including the string component has been performed in ref.~\cite{Lorenz:2007ze}. In order to get meaningful constraints they applied the bound on the exit scale as suggested in ref.\cite{Baumann:2006cd} which is discussed in section~\ref{sec:strat}. This study is relevant to the domain where the string tension is weakened by a large number of antibranes. We estimate that if $M>36$ these constraints will apply. Finally we comment that there may be limits on $G\mu$ which come from pulsar timing if the cosmic string network achieves scaling by the creation of loops and the subsequent emission of radiation (see ref.~\cite{Caldwell:1996en} and references therein). We caution that these should be considered to be less robust since they are more strongly effected by the small-scale dynamics, such as loop formation, of the cosmic string network. This is not completely understood in the case of standard cosmic strings in 3+1 dimensions; the situation with respect to the higher dimension cosmic strings is even less clear. Nonetheless, limits on the energy density in gravitational waves, $\Omega_{\rm g}h^2$, have substantially improved in recent times. Probably the most reliable limit is $\Omega_{\rm g}h^2<2\times 10^{-8}$ at frequencies $f=2\times 10^{-9}{\rm Hz}$~\cite{Jenet:2006sv}. Limits on $G\mu$ from such a bound were considered in ref.~\cite{Battye:2006pk} and they are dependent on the loop production size relative to the horizon $\alpha$. If $\alpha<10^{-4}$ then $G\mu>10^{-6}$ is excluded and hence the limit from CMB anisotropy is strongest, whereas for $\alpha>10^{-4}$ one finds that $G\mu> 10^{-10}/\alpha$ is excluded, which would be tighter than the CMB limit for $\alpha>10^{-3}$. It is clear that an improved understanding of the loop production mechanism coupled with expected improvements in the bound on $\Omega_{\rm g}h^2$ could lead to a more powerful constraint than is expected from PLANCK. | 7 | 10 | 0710.1541 |
0710 | 0710.3772_arXiv.txt | Procyon A is a bright F5IV star in a binary system. Although the distance, mass and angular diameter of this star are all known with high precision, the exact evolutionary state is still unclear. Evolutionary tracks with different ages and different mass fractions of hydrogen in the core pass, within the errors, through the observed position of Procyon A in the Hertzsprung-Russell diagram. For more than 15 years several different groups have studied the solar-like oscillations in Procyon A to determine its evolutionary state. Although several studies independently detected power excess in the periodogram, there is no agreement on the actual oscillation frequencies yet. This is probably due to either insufficient high-quality data (i.e., aliasing) or due to intrinsic properties of the star (i.e., short mode lifetimes). Now a spectroscopic multi-site campaign using 10 telescopes world-wide (minimizing aliasing effects) with a total time span of nearly 4 weeks (increase the frequency resolution) is performed to identify frequencies in this star and finally determine its properties and evolutionary state. | The bright F5 subgiant Procyon A is the primary of an astrometric binary system with a white dwarf in a 40 year orbit. Procyon A is the brightest northern-hemisphere asteroseismology candidate with well-determined characteristics, such as distance, mass and angular diameter. Brown {\etal} \cite{brown1991} were among the first to observe an excess power between 0.5 and 1.5 mHz in radial velocity observations confirmed by several other radial velocity studies, e.g. \cite{martic1999}, \cite{bouchy2002}, \cite{kambe2003}, \cite{martic2004}, \cite{eggenberger2004}, \cite{claudi2005}, \cite{leccia2007}. So far these studies have independently revealed detections of power excess, but there is no agreement yet on the actual oscillation frequencies. This may be due to aliases present in the spectral window, short mode lifetimes, shifts from the asymptotic relation due to avoided crossings, or any combination of these factors. Although the frequencies are not yet known in detail, most studies obtain a large frequency spacing of $55 \pm 1$ $\mu$Hz. \begin{table}[h!] \caption{\label{stelpar} Stellar parameters of Procyon A from \cite{allende2002}.} \begin{center} \begin{tabular}{lr@{$\pm$}l} \br Mass [M$_{\odot}$]& 1.42 & 0.06\\ T$_{\rm eff}$ [K] & 6512 & 49 \\ Radius [R$_{\odot}$] & 2.071 & 0.020\\ $\rm[Fe/H]$ [dex] & $-$0.09 & 0.03 \\ $v_{\rm rot}\sin i$ [km\,s$^{-1}$] & 3.16 & 0.50\\ \br \end{tabular} \end{center} \end{table} Another point of discussion is the fact that \cite{matthews2004} did not detect any power excess in their MOST photometry and published Procyon A as a flat liner. Other photometric studies such as WIRE \cite{bruntt2005} and a reanalysis of the MOST 2004 data \cite{regulo2005} claim to detect power excess in the same region as the radial velocity studies. A recent reanalysis of the MOST 2004 data and the analysis of new MOST data taken in 2005 reinforce the null detection of p-modes \cite{guenther2007}. This issue is discussed more extensively by \cite{bedding2005}, who claim that the non-detection of oscillations in Procyon by the MOST satellite is fully consistent with the ground based radial-velocity studies due to a combination of several noise sources and the low photometric amplitude of the oscillations. Procyon A is in a very interesting evolutionary state near the end of its main sequence life. The stellar parameters (see Table~\ref{stelpar}) are all known to high precision, but several evolutionary tracks with different ages and different hydrogen core mass fractions overlap, within the errors, with its position in the HR-diagram. The exact evolutionary state of a star can be revealed by means of asteroseismology and therefore the oscillation frequencies are needed. To determine the actual frequencies of Procyon A a ground-based multi-site campaign using 10 telescopes with a total time span of nearly 4 weeks was performed from December 28, 2006 till January 22, 2007. Here we present first results of this campaign. In Section 2 the campaign is described, while Section 3 discusses the way the data of the different telescopes are combined, and a final power spectrum is obtained. Some concluding remarks and future prospects are provided in Section 4. | The Procyon campaign presented here is the largest spectroscopic campaign, so far, aimed at the detection and identification of solar-like oscillations. Using 10 telescopes over a time span of nearly 4 weeks provided us with a unique data set with high time coverage and frequency resolution. The details of the data processing methods will be fully described by \cite{arentoft2008}, while the oscillation frequencies extracted from the full data set acquired during the spectroscopic Procyon campaign will be presented by \cite{bedding2008}. \begin{figure} \begin{center} \includegraphics[width=\linewidth]{powersmall.eps} \caption{\label{power}Final power spectrum of Procyon A with window-optimised weights. The data are high-pass filtered to remove low-frequency drifts.} \end{center} \end{figure} | 7 | 10 | 0710.3772 |
0710 | 0710.1777_arXiv.txt | Studies of stellar magnetism at the pre-main sequence phase can provide important new insights into the detailed physics of the late stages of star formation, and into the observed properties of main sequence stars. This is especially true at intermediate stellar masses, where magnetic fields are strong and globally organised, and therefore most amenable to direct study. This talk reviews recent high-precision ESPaDOnS observations of pre-main sequence Herbig Ae-Be stars, which are yielding qualitatively new information about intermediate-mass stars: the origin and evolution of their magnetic fields, the role of magnetic fields in generating their spectroscopic activity and in mediating accretion in their late formative stages, and the factors influencing their rotational angular momentum. | \subsection{Magnetism and rotation in the main sequence A and B stars} Between about 1.5 and 10 M$_{\odot}$, at spectral types A and B, about 5 \% of main sequence (MS) stars have magnetic fields with characteristic strengths of about 1kG. Such stars also show important chemical peculiarities and are thus usually called the magnetic chemically peculiar Ap/Bp stars. The strength of the magnetic fields of these stars cannot be explained by an envelope dynamo as in the sun. Until now, the most reliable hypothesis has been to assume a fossil origin for these magnetic fields. This hypothesis implies that the stellar magnetic fields are relics from the field present in the parental interstellar cloud. Its also implies that magnetic fields can (at least partially) survive the violent phenomena accompanying the birth of stars, and can also remain throughout their evolution and until at least the end of the MS, without regeneration. According to the fossil field model, we should observe magnetic fields in some pre-main sequence (PMS) stars of intermediate mass, the so-called Herbig Ae/Be stars. However no magnetic field was observed up to recently in these stars \cite[except HD~104237,][]{donati97}. Can we obtain some observational evidence of the presence of magnetic fields during the PMS phase of evolution, as predicted by the fossil field hypothesis? If some Herbig Ae/Be stars are discovered to have magnetic fields, is the fraction of magnetic to non-magnetic Herbig stars the same as the fraction for main sequence stars? Is the magnetic field in Herbig stars strong enough to explain the strength of that of Ap/Bp stars? Chemical peculiarities and magnetism are not the only characteristic properties observed in the Ap/Bp stars. Most magnetic MS stars have rotation periods (typically of a few days) that are several times longer than the rotation periods of non-magnetic MS stars (a few hours to one day). It is usually believed that magnetic braking, in particular during PMS evolution, when the star can exchange angular momentum with its massive accretion disk, is responsible for this low rotation \cite[][]{stepien00,stepien02}. An alternative involves a rapid dissipation of the magnetic field during the early stages of PMS evolution for the fastest rotators, due to strong turbulence induced by rotational shear developed under the surface of the stars, as the convection do in the solar-type stars \cite[see e.g.][]{lignieres96}. In this scenario, only slow rotators could retain their initial magnetic fields, and evolve as magnetic stars to the main sequence. So the question to be addressed is the following: does the magnetic field control the rotation of the star, or else does the rotation of the star control the magnetic field? We propose that this question can be answered by studying rotation and magnetic fields in Herbig Ae/Be stars. \subsection{The Herbig Ae/Be stars} The Herbig Ae/Be stars are intermediate-mass pre-main sequence stars, and therefore the evolutionary progenitors of the MS A and B stars. They are distinguished from the classical Be stars by their IR emission and the association with nebulae, characteristics which are due to their young age. They display many observational phenomena often associated with magnetic activity. First, high ionised lines are observed in the spectra of some stars \cite[e.g.][]{bouret97,roberge01}, and X-ray emission have been detected, coming from some Herbig stars \cite[e.g.][]{hamaguchi05}. In active cool stars, many of these phenomena are produced in hot chromospheres or coronae. Some authors mentioned rotational modulation of resonance lines which they speculate may be due to rotation modulation of winds structured by magnetic field \cite[][]{praderie86,catala89,catala91,catala99}. In the literature we find many clues of the presence of circumstellar disks around these stars, from spectroscopic data showing strong emission, and also from photometric data \cite[e.g.][]{mannings97,mannings00}. Recently, using coronagraphic data and interferometric data, some authors have also found direct evidence of circumstellar disks around these stars \cite[][]{grady99,grady00,eisner03}. A careful study of these disks shows that they have similar properties to the disk of their low mass counterpart \cite[][]{natta01}, the T Tauri stars, whose the emission lines are explained by magnetospheric accretion models \cite[][]{konigl91,muzerolle98,muzerolle01}. Finally \citet{muzerolle04} have sucessfully applied their magnetospheric accretion model to Herbig stars to explain the emission lines in their spectra. For all these reasons we suspect that the Herbig stars may host large-scale magnetic fields that should be detectable with current instrumentation. However, many authors tried to detect such fields without much success \cite[][]{catala93,catala99,donati97,hubrig04,wade07}. But in 2005, a new high-resolution spectropolarimeter, ESPaDOnS, has been installed at the canada-france-hawaii telescope (CFHT). We therefore decided to proceed to survey many Herbig stars in order to investigate rotation and magnetism in the pre-main sequence stars of intermediate mass. | We used the new spectropolarimeter ESPaDOnS installed at the CFHT to proceed in a survey of the Herbig stars, in order to investigate their rotation and magnetic field. We discovered four magnetic stars whose field topology is similar to the MS magnetic A-B stars. We also show that the magnetic intensities of these fields and the proportion of magnetic Herbig stars can explain the magnetic intensity and the proportion of magnetic fields among the MS stars, in the context of fossil field model. We therefore bring fundamental arguments in favour of this hypothesis. The four magnetic Herbig stars are slow rotators ($v\sin i < 41$ km.s$^{-1}$) which supports that magnetic Herbig Ae/Be stars are the progenitors of the magnetic Ap/Bp stars. Among these magnetic stars two have very low $v\sin i$ ($< 10$ km.s$^{-1}$) and are very young (age$ < 2.8$ Myr). Assuming these stars are true slow rotators, this implies that there exists a braking mechanism which acts very early during the PMS evolution of the intermediate mass stars. We could think that this braking mechanism has a magnetic origin, although among the undetected stars we also observe same stars with small $v\sin i$ ($\sim 15$ km.s$^{-1}$). The nature of the braking mechanism requires addition study. Finally it has been proposed that all Herbig stars should undergo magnetospheric accretion, as the spectra of all stars show similar emission. However, we calculated the minimum polar magnetic intensity of the magnetic Herbig stars and of two well-constrained undetected stars to get magnetospheric accretion, using three different models which have been developed for the T Tauri stars \cite[][]{konigl91,cameron93,shu94}. Considering the intensity of the magnetic stars, as well as the maximum magnetic intensity of the two well-constrained undetected stars, if they host a magnetic field, our first conclusion is that magnetospheric accretion cannot occur in all the Herbig stars (a statistic study taking into account all the undetected stars is in progress in order to confirm this result). Therefore, either the models that we used are not well adapted to the Herbig stars, or the emission lines are not only produced by magnetospheric accretion. A thorough observational and theoretical study of the emission in the spectra, as well as the surroundings of the Herbig Ae/Be stars is necessary to better understand the interaction of these stars with their surroundings. \vspace{1cm} | 7 | 10 | 0710.1777 |
0710 | 0710.4370_arXiv.txt | Based on a simulation of galaxy formation in the standard cosmological model, we suggest that a consistent picture for Gamma-Ray Bursts and star formation may be found that is in broad agreement with observations: {\it GRBs preferentially form in low metallicity environments and in galaxies substantially less luminous that $L_*$.} We find that the computed formation rate of stars with metallicity less than $0.1\zsun$ agrees remarkably well with the rate evolution of Gamma-Ray Bursts observed by Swift from $z=0$ to $z=4$, whereas the evolution of total star formation rate is weaker by a factor of about $4$. Given this finding, we caution that any inference of star formation rate based on observed GRB rate may require a more involved exercise than a simple proportionality. | The intriguing observational linkage between long duration Gamma-Ray Bursts (GRBs) and core-collapse supernovae (e.g., Stanek \etal 2003; Hjorth \etal 2003) suggests that the progenitors of GRBs may be very massive stars. This possible connection was predated by a proposed unified picture (Cen 1998). As such, it has been hoped that GRBs may be a good tracer of cosmic star formation (Wijers \etal 1998; Totani 1999; Lamb \& Reichart 2000; Blain \& Natarajan 2000; Porciani \& Madau 2001; Daigne \etal 2006; Coward 2006; Le \& Dermer 2007; Li 2007). However, recent observations indicate that typical GRBs may prefer relatively low metallicity environments ($\sim 0.1\zsun$) (Fynbo \etal 2003; Le Floch \etal 2003; Christensen \etal 2004; Fruchter \etal 2006; Stanek \etal 2006) and host galaxies significantly less luminous than $L_*$ (Fruchter \etal 1999,2006), although there is evidence that the actual spread in metallicity may be wide (Berger \etal 2006; Prochaska 2006; Wolf \& Podsiadlowski 2007). The aim of this {\it Letter} is to first address the issue of consistency of GRB environment with respect to metallicity and galaxy luminosity, i.e., the galaxy luminosity-metallicity relation, in the context of detailed simulation of galaxy formation in the standard cosmological model. Then, we make predictions on the evolution of GRB rate with redshift and highlight a possible dramatic difference between overall star formation history and GRB rate history, if GRBs are not an unbiased tracer of star formation. In particular, if GRBs are predominantly produced by stars with metallicity $\le 0.1\zsun$, the GRB rate is expected in our model to rise obstinately from $z=0$ to $z\sim 5$ by a factor of $\sim 100$, when it flattens out towards higher redshift, whereas the overall star formation rate rises rapidly only from $z=0$ to $z\sim 3$ and is roughly flat from $z\ge 3$ until $z\sim 7$. The evolution of GRB rate with redshift is thus expected to be stronger than that of star formation. | We utilize a simulation of galaxy formation in the standard cosmological model that has been shown to produce results consistent with extant observations of galaxy formation (e.g., Nagamine \etal 2006) to shed light on the relation between GRB rate and star formation rate. We find that a consistent picture for Gamma-Ray Bursts and star formation that is in broad agreement with observations would emerge, {\it if GRBs preferentially form in low metallicity environments and in galaxies substantially less luminous that $L_*$.} Because of the increase of metallicity with cosmic time, GRB rate consequently evolves more strongly with redshift than the overall star formation rate. We find that the observed evolution of GRB rate from $z=0$ to $z=4$ can be explained, {\it if GRBs are primarily produced by massive stars with metallicity less than $0.1\zsun$}, whereas an inclusion of stars with metallicity as high as $0.3\zsun$ yields GRB rate evolution from $z=0$ to $z=4$ inconsistent with observations. Therefore, we reach the conclusion that GRBs may not be a good tracer of cosmic star formation, especially over a long timeline. As a result, a simple inference of star formation rate or its derived quantities such as the ionizing photon production rate at high redshifts, based on the observed GRB rate, should be done with caution and may require careful calibrations. \smallskip | 7 | 10 | 0710.4370 |
0710 | 0710.4146_arXiv.txt | Motivated by the increasing use of the Kennicutt-Schmidt (K-S) star formation law to interpret observations of high redshift galaxies, the importance of gas accretion to galaxy formation, and the recent observations of chemical abundances in galaxies at \ztwo--3, I use simple analytical models to assess the consistency of these processes of galaxy evolution with observations and with each other. I derive the time dependence of star formation implied by the K-S law, and show that the sustained high star formation rates observed in galaxies at \ztwo--3 require the accretion of additional gas. A model in which the gas accretion rate is approximately equal to the combined star formation and outflow rates broadly reproduces the observed trends of star formation rate with galaxy age. Using an analytical description of chemical evolution, I also show that this model, further constrained to have an outflow rate roughly equal to the star formation rate, reproduces the observed mass-metallicity relation at \ztwo. | The motivations of this paper are several. First, the empirical Kennicutt-Schmidt (K-S) law \citep{s63,k98schmidt}, which states that the surface density of star formation is proportional to the surface density of gas, is widely used to interpret and describe star formation in galaxies, though its origins are not fully understood and it is just beginning to be tested at high redshift \citep{btg+04,csn+07,bcd+07}. The K-S law is a valuable tool when gas masses are not directly measurable; it has been used to estimate the gas masses of both high redshift galaxies \citep{ess+06mass} and local galaxies in the distant past \citep{cjp+07}. The consequences of the K-S law for the evolution of star formation at high redshift are therefore worth considering in more detail, as is the consistency of these consequences with observations. Second, the fueling of star formation by gas accretion is an essential element of models of galaxy formation, but has largely been neglected by observers of galaxies at high redshifts, mostly because of the lack of evidence for inflow in the spectra of these galaxies. In contrast, the evidence for strong outflows in galaxies at \ztwo--3 is well-known, most notably in the form of offsets between the redshifts of nebular emission lines, rest-frame UV absorption lines, and \lya\ emission (\citealt{pss+01}, Steidel et al.\ 2007, in prep). In spite of the lack of observations of inflow, however, its effects should be considered in the context of other observations, since a significant inflow rate would affect other, measurable properties. Finally, in recent years metallicity measurements of increasingly large samples of galaxies at $z>1$ have become possible (e.g.\ \citealt{kk00,pss+01,sga+04,sep+04,mlc+06}), including the detection of a mass-metallicity relation at $z>2$ \citep{esp+06}. These measurements still suffer considerably from limitations on the methods that can be used and from calibration uncertainties, but nevertheless they offer a unique opportunity to place constraints on star formation histories and gas flows at high redshift, provided that the effects of inflow, outflow and star formation can be disentangled. Many recent studies have addressed this issue in some detail, including \citet{ke99,kh05,d07}; and \citet{fd07}. The goal of this paper is to formulate simple, analytical models including star formation according to the K-S law, gas inflows and outflows, and chemical evolution. We would like to assess whether or not these models are consistent with each other and with our current observational knowledge of high redshift galaxies, and to see if they might give a general picture of how gas flows, star formation and metal enrichment may proceed at high redshift. In \S\ref{sec:ks} I derive the explicit time dependence of star formation implied by the K-S law and test its consistency with observations of galaxies at \ztwo. In \S\ref{sec:metals} I consider simple models of chemical evolution which incorporate both inflows and outflows of gas, and again test their consistency with observations of high redshift galaxies and with the results of the previous section. Some implications of the results are discussed in \S\ref{sec:discuss}. I adopt a cosmology with $H_0=70\;{\rm km}\;{\rm s}^{-1}\;{\rm Mpc}^{-1}$, $\Omega_m=0.3$, and $\Omega_{\Lambda}=0.7$, and use the \citet{c03} initial mass function (IMF). | \label{sec:discuss} The above results provide a coherent picture in which strong star formation is sustained by the accretion of gas at approximately the gas processing rate, the outflow rate is roughly equal to the SFR, and metal enrichment is modulated by both outflows and inflows. This is not a new result; \citet{fd07} reached many of the same conclusions using cosmological hydrodynamic simulations to reproduce the \ztwo\ mass-metallicity relation, and the idea of a system in which star formation is balanced by inflow dates to work by \citet{l72} and \citet{tl78}. Whatever the methods used to reach these conclusions, however, many questions remain about the mechanisms of gas flows and chemical enrichment at high redshift. The only quantity not yet tied to observations is the gas accretion rate, which the models require to be approximately equal to the gas processing rate. If the outflow rate is roughly equal to the SFR, the required accretion rate is $\sim60$ \msunyr\ (assuming the average SFR of the UV-selected sample; \citealt{ess+06stars}), and as much as several hundred \msunyr\ or higher for the most rapidly star forming galaxies. These values are in general agreement with theoretical expectations. For example, the predicted average gas accretion rates for galaxies in $10^{12}$ \msun\ haloes\footnote{Clustering results indicate that the \ztwo\ UV-selected galaxies are typically associated with $\sim10^{12}$ \msun\ haloes \citep{asp+05}.} given by \citet{kkwd05} are $\sim50$ \msunyr\ at the relevant redshifts, rising to several hundred \msunyr\ for the $10^{13}$ \msun\ haloes expected to host the most massive galaxies (though more recent simulations indicate rates a factor of $\sim2$ or more lower; D. Kere{\v s}, private communication). Gas accretion and cooling rates from the semi-analytic models of \citet{csw+06} are also of the right order of magnitude. The question remains as to why such high accretion rates have not yet been observed. One difficulty is that the velocity range of inflowing gas is likely to be much narrower than the several hundred \kms\ spread observed in the outflows. Another suggestion is that cold, filamentary accretion may dominate at high redshifts \citep{kkwd05,db06}, in which case detection would depend strongly on projection effects; alternatively, the accretion could occur largely in the form of minor mergers. Another possibility is that the accreting gas may be too hot to produce signatures in the observed wavebands, or such signatures may simply be too weak to detect. Even if the hot accreting gas does produce \ion{C}{4} emission, for example, the line would likely be weak because of the low metallicity of the gas, and it would be superposed on the already complicated \ion{C}{4} profile. Detailed modeling of the likely line strengths would help to place limits on detectable accretion rates. Gas heated to the virial temperature of $\sim10^6$ K must also be sufficiently cooled in order to fuel star formation; work by \citet{yssw02} and \citet{csw+06}, among others, discusses the mechanisms by which this might proceed. The strong star formation and gas accretion discussed herein will not continue indefinitely. Observations suggest that most of the star-forming galaxies currently detected at \ztwo--3 will become largely passively evolving by $z\sim1$ \citep{asp+05,pmd+06}. Because the high observed star formation rates require accretion of new gas to sustain, a decline in the SFR is a natural consequence of the drop in accretion rates at lower redshifts predicted by theoretical models. Many theorists and observers have also proposed that AGN feedback may be responsible for the termination of star formation (e.g.\ \citealt{hhc+06} and references therein). If an additional mechanism to shut off star formation is required, this is a plausible candidate, as the ubiquity of outflows suggests that starburst-driven winds may regulate star formation but do not usually terminate it, and this work implies that strong accretion and outflows may operate simultaneously, or at least alternate in relatively quick succession. Finally, these results underscore the importance of metallicity measurements for understanding gas flows. Until the flows can be observed and quantified directly, measurements of gas phase abundances offer the best hope for constraining the outflow and inflow rates of galaxies at high redshift. There are still considerable difficulties associated with the measurements of these metallicities, but we look forward to improved constraints from new IR spectra and photoionization modeling, and to more direct estimates of outflow rates from detailed spectra (e.g.\ \citealt{psa+00}). Such measurements will give a far more robust picture of star formation, gas flows and metallicity at high redshift than these simple models can provide. | 7 | 10 | 0710.4146 |
0710 | 0710.3416_arXiv.txt | The \wire\ satellite was launched in March 1999 and was the first space mission to do asteroseismology from space on a large number of stars. \wire\ has produced very high-precision photometry of a few hundred bright stars ($V<6$) with temporal coverage of several weeks, including K~giants, solar-like stars, \dss\ stars, and $\beta$~Cepheids. In the current work we will describe the status of science done on seven detached eclipsing binary systems. Our results emphasize some of the challenges and exciting results expected from coming satellite missions like COROT and Kepler. Unfortunately, on 23 October 2006, communication with \wire\ failed after almost eight years in space. Because of this sad news we will give a brief history of \wire\ at the end of this paper. | The failure of the main mission of the \wire\ satellite shifted focus to the star tracker, which has a small aperture of $52$\,mm and is equipped with a 512$\times$512 pixel SITe CCD. During observations the main target is positioned near the middle of the CCD and the four brightest stars in the 8$\times$8 degree field of view are also monitored. These four secondary stars are chosen automatically by the on-board software. Each observation comprises a time stamp and an 8$\times$8 pixel window centred on the target. Two images per second are collected for each star, resulting in typically one million CCD windows in three weeks. Due to pointing restrictions two fields are observed during each \wire\ orbit. The duty cycle for one star per orbit is optimally 40\%, but can be as low as 20\%. The orbital period has decreased from 96 to 93 minutes over the cause of the mission. The filter reponse of the star tracker is not well defined but Johnson $V+R$ has been suggested \citep{buzasi00}. For more details about \wire\ see Sect.~\ref{sec:epitaph}. \begin{figure*}[h] \centering \includegraphics[width=14cm]{wire_vancouver_hr_markbin.ps} \caption{\label{fig:hr} Hertzsprung-Russell diagram of around 100 stars observed with \wire. The locations of the seven dEBs and the Sun are marked.} \end{figure*} | We have presented on-going work on the observations and modelling of several eclipsing binary systems observed with the \wire\ satellite. The week-long temporal coverage of the targets has made it possible to study dEB systems that are notoriously difficult to observe from the ground, due to their periods being close to an integer number of days. The high precision of the photometry from \wire\ is a huge improvement compared to even the best photometric observations from the ground. The two important advantages of space based photometry are: \begin{itemize} \item No atmospheric scintillation noise \item High stability of the \wire\ star tracker over periods of several weeks. \end{itemize} These data have made it possible to push the limits on the constraints we can put on the theoretical models of these stars, covering the HR diagram from early B to late F type main sequence stars. However, at this level of precision the oscillations of the component stars, although amplitudes are small, need to be taken into account as part of the light curve analysis. In a broader context, the results presented here give an idea of the potential of secondary science on dEBs that can be done with data from the future satellite photometry missions. Many new systems will be discovered with COROT (see \citep{michel05}) and Kepler as discussed by \cite{koch07}. Although the new systems will be much fainter, follow-up spectroscopy and multi-band photometry can still be done from the ground. | 7 | 10 | 0710.3416 |
0710 | 0710.0553_arXiv.txt | The DAMA Collaboration has recently analyzed its data of the extensive WIMP direct search (DAMA/NaI) which detected an annual modulation, by taking into account the channelling effect which occurs when an ion traverses a detector with a crystalline structure. Among possible implications, this Collaboration has considered the case of a coherent WIMP-nucleus interaction and then derived the form of the annual modulation region in the plane of the WIMP-nucleon cross section versus the WIMP mass, using a specific modelling for the channelling effect. In the present paper we first show that light neutralinos fit the annual modulation region also when channelling is taken into account. To discuss the connection with indirect signals consisting in galactic antimatter, in our analysis we pick up a specific galactic model, the cored isothermal-sphere. In this scheme we determine the sets of supersymmetric models selected by the annual modulation regions and then prove that these sets are compatible with the available data on galactic antiprotons. We comment on implications when other galactic distribution functions are employed. Finally, we show that future measurements on galactic antiprotons and antideuterons will be able to shed further light on the populations of light neutralinos singled out by the annual modulation data. | \label{sec:intro} \begin{figure}[t] \centering \vspace{-20pt} \includegraphics[width=1.0\columnwidth]{A1_global.ps} \vspace{-30pt} \caption{WIMP--nucleon scattering cross-section as a function of the WIMP mass. The solid (dashed) line denotes the annual modulation region derived by the DAMA Collaboration with (without) the inclusion of the channeling effect. The two regions contain points where the likelihood- function values differ more than 4$\sigma$ from the null hypothesis (absence of modulation). These regions are obtained by varying the WIMP galactic distribution function (DF) over the set considered in Ref. \cite{bcfs} and by taking into account other uncertainties of different origins \cite{damalast}. The scatter plot represents supersymmetric configurations calculated with the supersymmetric model summarized in the Appendix. The (red) crosses denote configurations with a neutralino relic abundance which matches the WMAP cold dark matter amount ($0.092 \leq \Omega_{\chi} h^2 \leq 0.124$), while the (blue) dots refer to configurations where the neutralino is subdominant ($\Omega_{\chi} h^2 < 0.092$).} \label{fig:00} \end{figure} In a recent paper \cite{damalast} the DAMA Collaboration has analyzed the data of its extensive WIMP direct search (DAMA/NaI) \cite{dama} which measured an annual modulation effect at 6.3 $\sigma$ C. L., by taking into account the channelling effect. This effect occurs when an ion traverses a detector with a crystalline structure \cite{drob}. In Ref. \cite{damalast} implications of channelling have been discussed in terms of a specific modelling of this effect for the case of the DAMA NaI(Tl) detector; it is shown that the occurrence of channelling makes the response of this detector to WIMP-nucleus interactions more sensitive than in the case in which channelling is not included. Therefore, when applied to a WIMP with a coherent interaction with nuclei, the inclusion of the channelling effect implies that the annual modulation region, in the plane of the WIMP-nucleon cross section versus the WIMP mass, is considerably modified as compared to the one derived without including channelling. The extent of the modification depends on the specific model--dependent procedure employed in the evaluation of the channeling effect \cite{damalast}. These properties are shown in Fig. \ref{fig:00}, where the quantity $\sigma^{\rm nucleon}_{\rm scalar}$ denotes the WIMP-nucleon scalar cross-section, $\xi = \rho_{\rm WIMP}/\rho_0$ is the WIMP local fractional matter density and $m_{\chi}$ is the WIMP mass. The dashed line denotes the annual modulation region derived by the DAMA Collaboration without including the channeling effect \cite{dama}. The solid line shows the annual modulation region derived by the same Collaboration when the channeling effect is included as explained in Ref. \cite{damalast}. The regions displayed in Fig. \ref{fig:00} are derived by varying the WIMP galactic distribution function (DF) over the set considered in Ref.\cite{bcfs} and by taking into account other uncertainties of different origins \cite{damalast,nota1}. Fig. \ref{fig:00} shows that the effect of taking channelling into account is that the annual modulation region modifies its contour with an extension towards lighter WIMP masses. Most remarkably, for WIMP masses $\lsim $ 30 GeV, the WIMP-nucleon cross section involved in the annual modulation effect decreases sizeably, up to more than an order of magnitude. As mentioned before, the specific shape of the annual modulation region depends on the way in which channelling is modelled \cite{damalast}. These features are of great importance for a specific dark matter candidate, the light neutralino, which was extensively investigated in Refs. \cite{lowneu,lowdir,ind}. Actually, in these papers we analyzed light neutralinos, {\it i.e.} neutralinos with a mass $m_{\chi} \lsim 50$ GeV, which arise naturally in supersymmetric models where gaugino mass parameters are not related by a GUT--scale unification condition. In Refs. \cite{lowneu,lowdir} it is proved that, when R-parity conservation is assumed, these neutralinos are of great relevance for the DAMA/NaI annual modulation effect. In these papers it is also shown that in MSSM without gaugino mass unification the lower limit of the neutralino mass is $m_{\chi} \gsim$ 7 GeV \cite{hp}. In Fig. \ref{fig:00}, superimposed to the annual modulation regions is the scatter plot of the supersymmetric configurations of our model, whose features are summarized in the Appendix. One sees that, also when the channeling effect is taken into account, the light neutralinos of our supersymmetric model fit quite well the annual modulation region. In the present paper we consider the phenomenological consequences for light neutralinos when the annual modulation region is the one indicated by the solid line in Fig.\ref{fig:00}. More specifically we examine the properties of our supersymmetric population of light relic neutralinos in terms of the possible antimatter components generated by their pair annihilation in the galactic halo. To do this, we have to resort to a specific form for the WIMP DF. We take as our representative DF a standard cored isothermal-sphere model, though we do not mean to associate to this model prominent physical motivations as compared to other forms of DFs. Analyses similar to the one we present here for the cored isothermal-sphere can be developed for other galactic models. We will comment about some of them, selected among those considered in Ref. \cite{bcfs} (we will follow the denominations of this Reference to classify our DFs). The scheme of the present paper is the following. In Sect. II, we show how the model presented in Refs. \cite{lowneu,lowdir,ind} fits the DAMA/NaI annual modulation regions of Ref. \cite{damalast} for the case of the cored isothermal-sphere model. In Sect. III we combine these results with constraints derivable from available data on cosmic antiprotons; we also discuss the sensitivity of upcoming measurements on cosmic antiprotons for investigating the neutralino populations selected by the annual modulation regions. Complementary investigations by measurements of galactic antideuterons are presented in Sect. IV. Conclusions are drawn in Sect.V. The main features of the supersymmetric scheme adopted here are summarized in the Appendix. | In the present paper we have considered the annual modulation regions which the DAMA Collaboration has recently determined, by including also the channelling effect which occurs when an ion traverses a detector with a crystalline structure, such the detector of the DAMA/NaI experiment. The inclusion of the channelling effect implies that the annual modulation region is considerably modified as compared to the one derived without including channelling. The extent of the modification depends on the specific model--dependent procedure employed in the evaluation of the channeling effect. In the present paper we have considered the phenomenological consequences for light neutralinos when the annual modulation region includes the channelling effect as modelled in Ref.\cite{damalast}. We have proved that these annual modulation data are fitted by light neutralinos which arise naturally in supersymmetric models where gaugino mass parameters are not related by the a GUT--scale unification condition. The precise range of the neutralino mass which fits the annual modulation data depends on how the WIMP galactic distribution function is modelled and on a number of other assumptions, such as those mentioned in Sect. II. As an example, we have worked out in detail the case of a cored isothermal sphere DF. For this instance, the neutralino mass stays in the range $m_{\chi} \simeq (7 - 30)$ GeV, for values of the local rotational velocity, $v_0$, and of the local dark matter density, $\rho_0$, in the low-medium side of their own physical ranges, {\it i.e.} $v_0 \simeq$ (170 - 220) km sec$^{-1}$ and $\rho_0 \simeq (0.2 - 0.4)$ GeV cm$^{-3}$. Similar ranges are found also in the case of a Navarro--Frenk--White profile or for an isothermal model with a non-isotropic velocity dispersion. We have then shown that the populations of light neutralinos selected by the annual modulation regions are consistent with present data on galactic antiprotons. We have also derived the intervals of the diffusion parameters which provide this agreement in correlation with the specific galactic halo model. For instance, for neutralinos with a mass of 20 GeV and a cored isothermal model with $v_0 = 170$ km s$^{-1}$ we have $0.55 \lsim \delta \lsim 0.85$ and $L \lsim 3$ kpc when $ \rho_0 = \rho_0^{\rm max} = 0.42$ GeV cm$^{-3}$; instead when $ \rho_0 = \rho_0^{\rm min} = 0.20$ GeV cm$^{-3}$, $L$ may go up to 15 kpc with a range of $\delta$ which progressively shrinks to $\delta \sim 0.70 - 0.75$, when $L$ increases. We have also shown that future measurements of galactic antiprotons and antideuterons will offer, together with the upcoming data from DAMA/LIBRA, very interesting perspectives for further investigating the light neutralino populations selected by the annual modulation data. In case of models with a corotating halo or with triaxial spatial distributions, not investigated in the present paper, also heavier neutralinos can be involved. Finally, a word of caution should be said concerning the fact that the distribution of WIMPs in the Galaxy could deviate from the models mentioned above, mainly because of the presence of streams and/or clumpiness. In such instances, the analysis should be appropriately adapted, along the lines discussed in the present paper. | 7 | 10 | 0710.0553 |
0710 | 0710.0765_arXiv.txt | CCD $VRI$ photometry is presented for SN 2002hh from 14 days after the outburst till day 347. SN 2002hh appears to be normal type IIP supernova regarding both luminosity and the shape of the light curve, which is similar to SN 1999gi. | 7 | 10 | 0710.0765 |
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0710 | 0710.0003_arXiv.txt | We present a study of elemental abundances in a sample of thirteen Blue Compact Dwarf (BCD) galaxies, using the $\sim$10--37$\mu$m high resolution spectra obtained with Spitzer/IRS. We derive the abundances of neon and sulfur for our sample using the infrared fine-structure lines probing regions which may be obscured by dust in the optical and compare our results with similar infrared studies of starburst galaxies from ISO. We find a good correlation between the neon and sulfur abundances, though sulfur is under-abundant relative to neon with respect to the solar value. A comparison of the elemental abundances (neon, sulfur) measured from the infrared data with those derived from the optical (neon, sulfur, oxygen) studies reveals a good overall agreement for sulfur, while the infrared derived neon abundances are slightly higher than the optical values. This indicates that either the metallicities of dust enshrouded regions in BCDs are similar to the optically accessible regions, or that if they are different they do not contribute substantially to the total infrared emission of the host galaxy. | Blue Compact Dwarf Galaxies (BCDs) are dwarf galaxies with blue optical colors resulting from one or more intense bursts of star-formation, low luminosities (M$_B>-$18) and small sizes. The first BCD discovered was I\,Zw\,18 by \citet{Zwicky66}, which had the lowest oxygen abundance observed in a galaxy \citep{Searle72}, until the recent study of the western component of SBS0335-052 \citep{Izotov05}. Although BCDs are defined mostly by their morphological parameters, they are globally found to have low heavy element abundances as measured from their HII regions (1/30\,Z$_\odot$$\sim$1/2\,Z$_\odot$). The low metallicity of BCDs is suggestive of a young age since their interstellar medium is chemically unevolved. However, some BCDs do display an older stellar population and have formed a large fraction of their stars more than 1Gyr ago \citep[see][]{Loose85,Aloisi07}. The plausible scenario that BCDs are young is intriguing within the context of Cold Dark Matter models which predict that low-mass dwarf galaxies, originating from density perturbations much less massive than those producing the larger structures, can still be forming at the current epoch. However, despite the great success in detecting galaxies at high redshift over the past few years, bona fide young galaxies still remain extremely difficult to find in the local universe \citep{Kunth86,Kunth00,Madden06}. This is likely due to the observational bias of sampling mostly luminous more evolved galaxies at high redshifts. If some BCDs are truly young galaxies, they would provide an ideal local laboratory to understand the galaxy formation processes in the early universe. Over the past two decades, BCDs have been studied extensively in many wavelengths using ground-based and space-born instruments \citep [for a review see][]{Kunth00}. In the FUV, the Far Ultraviolet Spectroscopic Explorer (FUSE) has been used to study the chemical abundances in the neutral gas in several BCDs \citep{Thuan02, Aloisi03, Lebouteiller04}. Optical spectra have been obtained for a large number of BCDs and display strong narrow emission lines resulting from the intensive star-formation processes that take place in these systems \citep{Izotov97, Izotov99b, Pustilnik05, Salzer05}. The Infrared Space Observatory (ISO) revealed unexpectedly that despite their low metallicities, BCDs, such as SBS\,0335-052E, could still have copious emission from dust grains \citep{Thuan99, Madden00, Madden06, Plante02}. More recently, the Spitzer Space Telescope \citep{Werner04} has been used to observe these metal-poor dwarf systems in order to study their dust continuum properties and the polycyclic aromatic hydrocarbon (PAH) features \citep{Houck04b, Hogg05, Engelbracht05, Rosenberg06, Wu06, OHalloran06, Hunt06, Wu07}. Finally, radio observations have also been performed for several BCDs to study their HI kinematics and distribution \citep{Thuan04} and thermal/non-thermal continuum emission properties \citep{Hunt05}. \begin{deluxetable*}{lrllllcc} \tabletypesize{\scriptsize} \setlength{\tabcolsep}{0.02in} \tablecaption{Observing Parameters of the Sample\label{tab1}} \tablewidth{0pc} \tablehead{ \colhead{Object Name} & \colhead{RA} & \colhead{Dec} & \colhead{AORKEY} & \colhead{Observation} & \colhead{Redshift} & \multicolumn{2}{c}{On-source Time (Sen)}\\ \colhead{} & \colhead{(J2000)} & \colhead{(J2000)} & \colhead{} & \colhead{Date} & \colhead{} & \colhead{SH} & \colhead{LR} \\ } \startdata Haro11 & 00h36m52.5s & -33d33m19s & 9007104 & 2004-07-17 & 0.0206 & 480 & 240 \\ NGC1140 & 02h54m33.6s & -10d01m40s & 4830976 & 2004-01-07 & 0.0050 & 480 & 240 \\ SBS0335-052E & 03h37m44.0s & -05d02m40s & 11769856 & 2004-09-01 & 0.0135 & 1440 & 960 \\ NGC1569 & 04h30m47.0s & +64d50m59s & 9001984 & 2004-03-01 & $\sim$ 0& 480 & 240 \\ IIZw40 & 05h55m42.6s & +03d23m32s & 9007616 & 2004-03-01 & 0.0026 & 480 & 240 \\ UGC4274 & 08h13m13.0s & +45d59m39s & 12076032 & 2004-10-23 & 0.0015 & 120 & 56 \\ & & & 12626688 & 2004-11-11 & & 120 & 56 \\ IZw18 & 09h34m02.0s & +55d14m28s & 9008640 & 2004-03-27 & 0.0025 & 480 & 240 \\ & & & 16205568 & 2005-12-16 & & 2880 & 1440 \\ VIIZw403 & 11h27m59.9s & +78d59m39s & 9005824 & 2004-12-09 & $\sim$ 0& 480 & 240 \\ Mrk1450 & 11h38m35.6s & +57d52m27s & 9011712 & 2004-12-12 & 0.0032 & 480 & 240 \\ UM461 & 11h51m33.3s & -02d22m22s & 9006336 & 2005-01-03 & 0.0035 & 480 & 240 \\ & & & 16204032 & 2006-01-14 & & 1440 & \nodata \\ SBS1210+537A & 12h12m55.9s & +53d27m38s & 8989952 & 2004-06-06 & \nodata & 480 & 240 \\ Tol1214-277 & 12h17m17.1s & -28d02m33s & 9008128 & 2004-06-28 & 0.0260 & 480 & 240 \\ Tol65 & 12h25m46.9s & -36d14m01s & 4829696 & 2004-01-07 & 0.0090 & 480 & 240 \\ UGCA292 & 12h38m40.0s & +32d46m01s & 4831232 & 2004-01-07 & 0.0010 & 480 & 240 \\ Tol1304-353 & 13h07m37.5s & -35d38m19s & 9006848 & 2004-06-25 & 0.0140 & 480 & 240 \\ Pox186 & 13h25m48.6s & -11d37m38s & 9007360 & 2004-07-14 & 0.0039 & 480 & 240 \\ CG0563 & 14h52m05.7s & +38d10m59s & 8992512 & 2005-05-30 & 0.0324 & 240 & 120 \\ CG0598 & 14h59m20.6s & +42d16m10s & 8992256 & 2005-03-19 & 0.0575 & 480 & 240 \\ CG0752 & 15h31m21.3s & +47d01m24s & 8991744 & 2005-03-19 & 0.0211 & 480 & 240 \\ Mrk1499 & 16h35m21.1s & +52d12m53s & 9011456 & 2004-06-05 & 0.0090 & 480 & 240 \\ {\rm [RC2]}A2228-00 & 22h30m33.9s & -00d07m35s & 9006080 & 2004-06-24 & 0.0052 & 480 & 240 \\ \enddata \tablecomments{The coordinates and redshifts of the objects are cited from the NASA/IPAC Extragalactic Database (NED). In this paper, we only include the analysis of thirteen out of twenty-two sources which have SNRs sufficient for our abundance study. CG0563, CG0598 and CG0752 are included in the original sample as BCD candidates, however, they appear to be more starburst like \citep[see][]{Hao07}. Thus even though they have high SNR, we do not include these three sources in this study.} \end{deluxetable*} Metallicity is a key parameter that influences the formation and evolution of both stars and galaxies. Detailed studies of the elemental abundances of BCDs have already been carried out by several groups \citep{Izotov99b, Kniazev03, Shi05} and the well known metallicity-luminosity relation has also been studied in detail in the environment of dwarf galaxies \citep{Skillman89, Hunter99, Melbourne02}. However, because these studies were performed in the optical, they were limited by the fact that the properties of some of the deeply obscured regions in the star-forming galaxies may remain inaccessible due to dust extinction at these wavelengths. In fact, \citet{Thuan99}, using ISO, have shown that the eastern component of SBS\,0335-052 does have an embedded super star cluster (SSC) that is invisible in the optical while contributing $\sim$75\% to the bolometric luminosity \citep[see also][]{Plante02,Houck04b}, even though it has very low metallicity (12+log(O/H)=7.33), which would in principle imply a low dust content. In addition to probing the dust enshrouded regions, emission in the infrared also has the advantage that the lines accessible at these wavelengths are less sensitive to the electron temperature fluctuations than the corresponding optical lines of the same ion. In the infrared, more ionization stages of an element become available as well. The improved sensitivity of the Infrared Spectrograph (IRS\footnote{The IRS was a collaborative venture between Cornell University and Ball Aerospace Corporation funded by NASA through the Jet Propulsion Laboratory and the Ames Research Center.}) \citep{Houck04a} on Spitzer has enabled us to obtain for the first time infrared spectra for a much larger sample of BCDs than was previously possible \citep{Thuan99, Madden00, Verma03, Martin06}, thus motivating this study to probe the heavy element abundances in BCDs. In this paper, we analyze Spitze/IRS spectra of thirteen BCDs and present elemental abundances of neon and sulfur, which are both primary elements produced by the same massive stars in the nuclear synthesis processes. In section 2, we describe the sample selection, observations and data reduction. We present our results on the chemical abundances in section 3, followed by a comparison of the optical and infrared derived abundances in section 4. We show the interplay between the abundances and PAH emission in section 5. Finally, we summarize our conclusions in section 6. | We have studied the neon and sulfur abundances of thirteen BCDs using Spitzer/IRS high-resolution spectroscopy. Our analysis was based on the fine-structure lines and the hydrogen recombination line detected in the SH spectra of the {\em IRS}, combined with the radio continuum, H$\alpha$ images and integrated optical spectral data in some cases. We find a positive correlation between the neon and sulfur abundances, though sulfur appears to be more under-abundant than neon (with respect to solar). The ratio of Ne/S for our sources is on average 11.4$\pm$2.9, which is consistent with what has been found in other HII regions using infrared data. However, this average ratio appears to be higher than the corresponding optical value of 6.5$\pm$1.8 (in BCDs), which could be due to the adopted ICFs in the optical studies. When comparing the newly derived neon and sulfur abundances with the oxygen abundances measured from the optical lines, we find a good overall agreement. This indicates that there are few completely dust enshrouded HII regions in our BCDs, or if such HII regions are present, they have similar metallicities to the ones probed in the optical. Finally, the infrared derived neon and sulfur abundances also correlate, with some scatter, with the corresponding elemental abundances derived from the optical data. | 7 | 10 | 0710.0003 |
0710 | 0710.0185.txt | Observational and theoretical evidence suggests that coronal heating is impulsive and occurs on very small cross-field spatial scales. A single coronal loop could contain a hundred or more individual strands that are heated quasi-independently by nanoflares. It is therefore an enormous undertaking to model an entire active region or the global corona. %therefore requires simulating the %time-dependent plasma evolution of enormous numbers of %elemental strands. Three-dimensional MHD codes have inadequate spatial resolution, and 1D hydro codes are too slow to simulate the many thousands of elemental strands that must be treated in a reasonable representation. Fortunately, thermal conduction and flows tend to smooth out plasma gradients along the magnetic field, so ``0D models" are an acceptable alternative. We have developed a highly efficient model called Enthalpy-Based Thermal Evolution of Loops (EBTEL) that accurately describes the evolution of the average temperature, pressure, and density along a coronal strand. It improves significantly upon earlier models of this type---in accuracy, flexibility, and capability. It treats both slowly varying and highly impulsive coronal heating; it provides the time-dependent differential emission measure distribution, $D\!E\!M(T)$, at the transition region footpoints; and there are options for heat flux saturation and nonthermal electron beam heating. EBTEL gives excellent agreement with far more sophisticated 1D hydro simulations despite using four orders of magnitude less computing time. It promises to be a powerful new tool for solar and stellar studies. | An abundance of observational and theoretical evidence indicates that much of the corona is highly dynamic and evolves in response to heating that is strongly time-dependent. The evidence further suggests that the cross-field spatial scale of the heating is very small, so that unresolved structure is ubiquitous. In particular, many if not all coronal loops are bundles of thin strands that are heated impulsively and quasi-randomly by nanoflares. It is estimated that a single loop contains several tens to several hundreds of such strands. See \citet{k06} for a detailed justification of these ideas and references to relevant work. Three-dimensional (3D) magnetohydrodynamic simulations are extremely useful for studying the source of coronal heating (instabilities of electric current sheets, reconnection, turbulence, etc.), but they cannot adequately address the complexity that is present in a single coronal loop, much less an entire active region. A more feasible approach is to treat the magnetic field as static and to solve the one-dimensional (1D) hydrodynamic equations along many representative flux strands using an assumed heating rate. The individual strands {\it must} be treated separately. It is not valid to approximate a loop as a monolithic structure with uniform heating corresponding to the average for the component strands. This gives a completely different and incorrect result. There is reason to believe that the diffuse corona that lies between distinct bright loops is also comprised of elemental strands (e.g., \citealt{anb07}). If roughly 100 strands are present in a single loop, then the numbers present in active regions and the global Sun are truly staggering. While it is possible to construct a limited number of model active regions with time-dependent 1D simulations \citep{ww07}, it is not possible to investigate a wide range of values for the coronal heating parameters that must be assumed, such as the dependence on magnetic field strength, loop length, etc. \citep{mdk00}. This is a major limitation, since we are still struggling to identify the properties and physical origin of the heating mechanism. Progress in the foreseeable future must therefore rely on simplified solutions to the hydro equations that treat field-aligned averages and are much less computationally intensive. These are sometimes called ``0D models" because there is only one value of temperature, pressure, and density at any given time in the simulation. 0D models have been developed previously by \citet{fh90} and \citet{kp93}, but the best known is that of \citet{c94}. It has been used to study a variety of topics, including coronal loops \citep{ck97, ck06, kc01,petal06}, flares \citep{rw02,pak02}, post-eruption arcades \citep{rf05}, and active stellar coronae \citep{ck06}. We have learned a great deal with the Cargill model, and our understanding has now advanced to the point where a more accurate and flexible model is required. This article presents an improved 0D model called Enthalpy-Based Thermal Evolution of Loops (EBTEL). As the name suggests, a key aspect of the model is an explicit recognition of the important role that enthalpy plays in the energy budget. EBTEL improves upon the Cargill model in several important ways. First, whereas the Cargill model is limited to an instantaneous heat pulse, EBTEL accommodates any time-dependent heating profile and can include a low-level background heating if desired. Second, EBTEL accounts for thermal conduction cooling and radiation cooling at all times during the evolution. The Cargill model assumes that only one or the other operates at any given time. Third, EBTEL has options for heat flux saturation and nonthermal electron beam heating. Finally, EBTEL is unique among 0D models in that it provides the time-dependent differential emission measure distribution of the transition region footpoints. Emission from the transition region plays a critical role in spatially unresolved observations, such as stellar observations and observations of the solar spectral irradiance, which is important for space weather \citep{lean97}. Note that footpoint emission is not limited to the cooler ($<1$ MK) plasma traditionally associated with the transition region. It can also include hot emissions that originate from the base of very hot loops. The so-called moss seen in the ``coronal" channels of the {\it Transition Region and Coronal Explorer (TRACE)} is an example \citep{betal99,mkb00}. We describe the coronal and transition region parts of EBTEL in the next two sections. We then present example simulations and compare them with corresponding simulations from a 1D model and, in one case, the Cargill model. We conclude with a discussion of EBTEL and the possible significance of the example simulations. | As evidenced by these examples, our simple 0D model is an excellent proxy for more sophisticated and far more computationally intensive 1D hydro simulations. It improves substantially on the 0D models of \citet{c94}, \citet{fh90}, and \citet{kp93}. The Cargill model assumes that heating is instantaneous and that cooling occurs either by thermal conduction or by radiation, but not by both at the same time. The Fisher-Hawley model: (1) predicts abrupt evolutionary changes as the strand evolves between three distinct regimes; (2) does not account for the evaporation that continues well beyond the end of an impulsive heating event; and (3) cannot return to the pre-event state due to unphysical catastrophic cooling. The Kopp-Poletto model shares some similarities with EBTEL, but it treats the flows in a fundamentally different way. Like EBTEL, it equates the enthalpy carried by evaporative upflows with an excess heat flux, but the excess is determined relative to the pre-event state, rather than to the time-varying radiative losses from the transition region. Condensation downflows in the model are given by a density-dependent fraction of the free-fall velocity. In actuality, gravity plays no direct role in condensation, since the downflows are driven by pressure gradient deficits relative to hydrostatic equilibrium, in the same way that evaporative upflows are driven by pressure gradient excesses. Gravity sets the value of the hydrostatic gradient, but it is only the deficit or excess relative to this value that is important for the flows. Inclined strands experience essentially the same condensation and evaporation as do upright strands of the same length. Finally, EBTEL has advantages over all three of the other models in that it provides the $D\!E\!M(T)$ of the transition region and treats nonthermal electron beams and heat flux saturation. One obvious application of EBTEL is to investigate the idea that the basic structural elements of the corona are very thin, spatially unresolved magnetic strands that are heated impulsively. Loops may be bundles of such strands, as reviewed in \citet{k06}, and the diffuse corona may be similarly structured. Differential emission measure distributions are one important test of this idea. Observed $D\!E\!M(T)$ curves from active regions and the quiet Sun tend to be peaked near $10^{6.5}$ and $10^{6.1}$ K, respectively, and to have a slope (temperature power law index) $\ge 0.5$ coolward of the peak \citep{rd81,dm93,betal96}. This is consistent with the coronal $D\!E\!M(T)$ curves of examples 1 and 4 (Figures \ref{fig:dem1tr} and \ref{fig:dem4}). The full loop curves are discrepant, on the other hand, due to the strong contribution from the transition region. The cited observations were made on the disk and should in principle include the transition region component. However, it is possible that absorption from chromospheric material such as spicules significantly attenuates the intensities of transition region lines used to construct the $D\!E\!M(T)$ curves (e.g., \citealt{ddg95}; \citealt{detal99}; \citealt{df82}; \citealt{so79}). We are currently investigating the magnitude of this effect. One of the great mysteries of coronal physics that has come to light in the last few years is the discovery that warm ($\sim 1$ MK) coronal loops are much denser than expected for quasi-static equilibrium and live for much longer than a cooling time. The loops are therefore neither steadily heated nor cooling as monolithic structures. It has been shown that the observed densities and timescales can be explained by bundles of nanoflare heated strands, as long as nanoflares do not all occur at the same time (see \citealt{k06} and references cited therein). Neighboring strands will therefore have different temperatures, and loops are predicted to have multi-thermal cross sections. In particular, emission should be produced at temperatures hotter than 3 MK. Hot loops are sometimes observed at the locations of warm loops, but not always. Example 5 suggests that nonthermal electron beams are a possible explanation for the lack of hot emission. As we have discussed, beams can produce excess densities through evaporation without the need for high temperatures. We have just begun to explore this possibility. For now, we note that the coronal $D\!E\!M(T)$ curve of example 5 (Figure \ref{fig:dem5}) bears a close resemblance to the observed curves reported by \citet{setal01} for a loop seen above the limb. %Other problems now being pursued with EBTEL include modeling the %emission characteristics of coronal arcades \citep{pk07b}, modeling the %light curves of solar flares \citep{rgm07}, investigating the %resolvability of elemental loop strands \citep{pk07a}, and modeling coronal %loops as self-organized critical systems \citep{lk05,kld06}. In this %last study, loop strands are assumed to become tangled by turbulent %photospheric %convection and to release magnetic energy when the misalignment between %adjacent strands reaches a critical angle, as expected for the secondary %instability \citep{dka05}. EBTEL is used to follow the %plasma evolution of the strands. Preliminary results suggest that %this model can reproduce the light curves of loops observed by the %Soft X-ray Imager (SXI) on the GOES-12 satellite \citep{lkm07}. However, %many more simulations with different combinations of model parameters are %necessary before any definitive conclusions can be drawn. This would %be very difficult with a computationally intensive 1D hydro code, but is %no problem with EBTEL. %Another major program we have begun under the partial sponsorship of %NASAs Living With a Star Program is to build realistic %models of active regions and the global Sun. Our approach is to %construct a ``magnetic skeleton" %by extrapolating photospheric magnetograms and to populate many representative %field lines with plasma using EBTEL. Models of this kind have been %published before, but they use static equilibrium strand solutions and cannot %adequately reproduce imaging observations over %a wide range of temperatures \citep{letal04, setal04, metal05, ww06, bw07}. %When the models are adjusted to %resemble soft X-ray images, they %they predict too little warm ($\sim 1$ MK) emission at coronal altitudes %(EUV loops) and too much warm emission at the transition %region footpoints of hot coronal structures (EUV moss). There is %reason to believe that impulsive heating can improve the situation %\citep{kld06, pk07b}. Indeed, \citet{ww07} have built an impressive %model of an active region using the NRLFTM 1D hydro code and %find that impulsive heating gives a significant improvement over %steady heating. The agreement with observations is still not adequate, %however. %Since 1D simulations are computationally very intensive, %Warren and Winebarger were only able to consider one parameterized form for the %coronal heating function. Using EBTEL, we will examine a wide range of %heating parameters, including %the magnitude, duration, and frequency of the nanoflares, and their dependence %on magnetic field strength and field line length %(see \citealt{mdk00}). We will also consider different proportions of %direct heating and nonthermal electron beams. %We estimate that 100 million hydro simulations %but is a manageable task with EBTEL. Identifying the coronal heating %parameters will provide valuable constraints for testing competing %theories of the heating mechanism. It should ultimately lead to a %physics-based operational model for nowcasting and %forecasting the X-ray and UV spectral irradiance. In conclusion, EBTEL is a powerful new tool that can be applied to a variety of problems in which large numbers of evolving strands must be computed. For example, it is now feasible to construct multiple models of nanoflare-heated active regions or entire stars and therefore to examine a wide array of nanoflare parameters (magnitude, lifetime, occurrence rate, dependence on field strength and strand length, etc.). By determining which parameters best reproduce the observations, we can place important constraints on the heating and thereby gain insight into the physical mechanism (e.g., \citealt{mdk00}; \citealt{setal04}; \citealt{ww06}). EBTEL is currently being used to study the emission characteristics of coronal arcades \citep{pk07b}, to explain the light curves of solar flares \citep{rgm07}, and to model coronal loops as self-organized critical systems \citep{lk05,kld06}. Interested users are invited to contact us for a copy of our IDL code. %We end by noting possible future improvements: %the inclusion of kinetic energy, cross-sectional area variation, and %gravitational stratification. %% Included in this acknowledgments section are examples of the %% AASTeX hypertext markup commands. Use \url without the optional [HREF] %% argument when you want to print the url directly in the text. Otherwise, %% use either \url or \anchor, with the HREF as the first argument and the %% text to be printed in the second. | 7 | 10 | 0710.0185 |
0710 | 0710.0373_arXiv.txt | The physics behind the acceleration of the cosmic expansion can be elucidated through comparison of the predictions of dark energy equations of state to observational data. In seeking to optimize this, we investigate the advantages and disadvantages of using principal component analysis, uncorrelated bandpowers, and the equation of state within redshift bins. We demonstrate that no one technique is a panacea, with tension between clear physical interpretation from localization and from decorrelated errors, as well as model dependence and form dependence. Specific lessons include the critical role of proper treatment of the high redshift expansion history and the lack of a unique, well defined signal-to-noise or figure of merit. | } The acceleration of the universe poses a fundamental mystery to cosmology, gravitation, and quantum physics. Understanding the nature of the dark energy responsible for the acceleration relies on careful, robust measurements of the dark energy properties, in particular its equation of state (EOS), or pressure to energy density, ratio that directly enters the Friedmann equation for cosmic acceleration. As scientists design the next generation of dark energy experiments they seek to optimize the measurements for the clearest insight into this unknown physics. Two critical pieces of information will be the value of the EOS at some epoch, such as the present, and a measure of its time variation, in much the way that early universe inflation theories are classified by the value of the spectral index and its running. The best parametrized EOS are physics based and model independent, i.e.\ able to describe dark energy dynamics globally, or at least over a wide range of behaviors. Such EOS are very successful at fitting to data and projecting the results of future experiments, and can be robust to bias against inexact parametrization. Other approaches seek to remove one drawback of parametrized EOS by not assuming a functional form for the time variation, lest the true dark energy model lie outside the apparently wide range of validity of the form, i.e.\ they aim for form independence. Two major avenues for achieving this are decomposition into basis functions or principal components (e.g.\ \cite{HutStark03}, also see \cite{CritPog05,ShapTurn06,SimpBrid06,DickKnoxChu06,Ste06,HutPeir07}) and individual values of the EOS $w(z)$ over finite redshift bins, which become more general as the number of elements increases. However uncertainties in estimation of the EOS properties also grow as the number of principal components or bins increases. This article begins by examining general properties of the cosmological data and its dependence on the EOS in \S\ref{sec:cosdep}. Many of the later, detailed results will already be foreshadowed by this straightforward and general analysis. In \S\ref{sec:pca} we examine principal component analysis of the EOS and in \S\ref{sec:band} uncorrelated bandpowers. Bins of EOS in redshift is investigated in \S\ref{sec:bin}, including figures of merit for quantifying the uncertainties. Further concentration on the crucial role of the high redshift EOS, and the risk of biasing parameter estimation, occurs in \S\ref{sec:whi}. We consider physical constraints on EOS properties in \S\ref{sec:wlimit} and summarize our results and conclude in \S\ref{sec:concl}. | } The dark energy equation of state properties contain clues crucial to understanding the nature of the acceleration of the cosmic expansion. Deciphering those properties from observational data involves a combination of robust analysis and clear interpretation. We considered three approaches -- principal components, uncorrelated bandpowers, and binning; none of the approaches provides a panacea. In particular, we identify issues of dependence on basis functions, binning variables, and baseline models. The three approaches are not truly nonparametric and physical interpretation (not merely the values) of the results in the two decorrelated basis techniques depends on model, priors, and data, indeed even on an implicitly assumed functional form. Nevertheless, principal components can give a useful guide to the qualitative sensitivity, the best constrained aspects, of the data. The uncorrelated bin approach unfortunately does not truly deliver uncorrelated bandpowers for the equation of state. This approach using the square root of the Fisher matrix does not tightly localize the information (without a strong prior), making the interpretation nontrivial. This property of nonlocality is inherent in the cosmological characteristics. One might prefer to stay with the original binned equations of state used as the initial step for this technique, which are readily interpreted. Conversely, if the modes can be localized, the interpretation is easy, but in that case the original Fisher matrix is close to diagonal and thus the original bins almost uncorrelated. Hence, again, one might as well stay with the bin parameters which have a clear meaning. Indeed the goal is understanding the physics, not obtaining particular statistical properties. Decorrelated parameters that are not readily interpretable physically are of limited use; for example one still prefers to analyze the cosmic microwave background in terms of physical quantities such as physical matter density and spectral tilt rather than the principal axes of the eigenvectors. Note that the uncertainty on the EOS behavior $\sigma(w(z))$ is the same whether calculated by PCA (if all modes are kept), uncorrelated bands, or binned EOS, since the same information is in the data. We also emphasize that the modes most clearly determine the effect on the equation of state, not the weights, which are often the only quantity displayed. Moderately localized, even all positive, weights do not guarantee a localized physical effect. A further caution is that locality and positivity of weights can owe more to prior restrictions, especially the treatment of the high redshift equation of state, than to the data itself. Assuming a fixed value for the high redshift equation of state has major, widespread impacts on the results, ranging from strongly misestimated uncertainties to spurious localization to bias in the derived cosmology. We emphasize that it is essential to fit for the high redshift behavior in order not to be misled. Adding CMB data and marginalizing over a new, high redshift bin removes the ill effects of bias but ``cancels out'', providing no new constraints; multiple data points for $z>2$ are required, such as from high redshift distances or weak lensing measurements of the mass growth behavior. Assuming that dark energy is negligible at $z>2$ is also effectively assuming a functional form -- precisely what the use of eigenmodes was supposed to avoid. Indeed, functional forms do not have many of the basis, model, binning, etc.\ dependences of eigenmodes, while principal components are in turn not fully form independent. If one assumes a functional form to obtain informative constraints on the equation of state, one must indeed choose the form to represent robustly the physical behavior (as has been shown to be widely the case for $w(a)=w_0+w_a(1-a)$ by \cite{Linder03,Linder0708}), and carefully check the range of validity of the conclusions by examining other forms. A good complementary analysis tool would be the binned equation of state approach examined here. Regardless of the form of analysis, only a finite amount of information can be extracted from even next generation data. As has been concluded for functional equations of state and principal component analysis \cite{LinHut05}, the analysis here in terms of binned equation of state indicates that only two physically informative parameters can be fit with realistic accuracy. However, we identify several issues in the PCA and uncorrelated bin approaches that cause accuracy or signal to noise criteria to be ill defined. Similar difficulties arise in condensing the physical information on dark energy to a single figure of merit; the number is quite sensitive to cosmologically irrelevant aspects like the binning used (as well as very dependent on the treatment of the high redshift dark energy behavior). In conclusion, physically motivated fitting of the equation of state such as the $w_0$-$w_a$ parametrization in complement with a binned equation of state approach (perhaps with physical constraints such as outlined in \S\ref{sec:wlimit}) have the best defined, clearest to interpret, and robust insights of the approaches we considered. With any method, one must use caution regarding the influence of priors and fit the dark energy physics over the entire expansion history. | 7 | 10 | 0710.0373 |
0710 | 0710.0780_arXiv.txt | $Chandra$ and $XMM-Newton$ observations of the Cartwheel galaxy show $\sim{}17$ bright X-ray sources ($\gtrsim{}5\times{}10^{38}$ erg s$^{-1}$), all within the gas-rich outer ring. We explore the hypothesis that these X-ray sources are powered by intermediate-mass black holes (IMBHs) accreting gas or undergoing mass transfer from a stellar companion. To this purpose, we run $N$-body/SPH simulations of the galaxy interaction which might have led to the formation of Cartwheel, tracking the dynamical evolution of two different IMBH populations: halo and disc IMBHs. Halo IMBHs cannot account for the observed X-ray sources, as only a few of them cross the outer ring. Instead, more than half of the disc IMBHs are pulled in the outer ring as a consequence of the galaxy collision. However, also in the case of disc IMBHs, accretion from surrounding gas clouds cannot account for the high luminosities of the observed sources. Finally, more than 500 disc IMBHs are required to produce $\lesssim{}15$ X-ray sources via mass transfer from very young stellar companions. Such number of IMBHs is very large and implies extreme assumptions. Thus, the hypothesis that all the observed X-ray sources in Cartwheel are associated with IMBHs is hardly consistent with our simulations, even if it is still possible that IMBHs account for the few ($\lesssim{}1-5$) brightest ultraluminous X-ray sources (ULXs). | Ring galaxies are among the most fascinating objects in the Universe. The Cartwheel galaxy is surely the most famous of them, and also the biggest (with its optical diameter of $\sim{}$50-60 kpc) and the most studied. Cartwheel has been thoroughly observed in almost every band: H$\alpha{}$ (Higdon 1995) and optical (Theys \& Spiegel 1976, Fosbury \& Hawarden 1977) images, red continuum (Higdon 1995), radio line (Higdon 1996) and continuum (Higdon 1996; Mayya et al. 2005), near- (Marcum, Appleton \& Higdon 1992) and far-infrared (Appleton \& Struck-Marcell 1987a), line spectroscopy (Fosbury \& Hawarden 1977) and X-ray (Wolter, Trinchieri \&{} Iovino 1999; Gao et al. 2003; Wolter \& Trinchieri 2004; Wolter, Trinchieri \& Colpi 2006). Cartwheel exhibits a double ringed shape, with some transversal 'spokes' (Theys \& Spiegel 1976, Fosbury \& Hawarden 1977; Higdon 1995), which have been detected only in few of the $\lesssim{}300$ known ring galaxies (Arp \& Madore 1987; Higdon 1996). The outer ring is rich of HII regions, especially in its southern quadrant, and dominates the H$\alpha{}$ emission (Higdon 1995). This implies that Cartwheel is currently undergoing an intense epoch of star formation (SF, with SF rate $\sim{}20-30\,{}M_\odot{}$ yr$^{-1}$; Marston \& Appleton 1995; Mayya et al. 2005), almost entirely confined to the outer ring. Inner ring, nucleus and spokes are nearly devoid of gas and dominated by red continuum emission, indicating a relatively old, low- and intermediate-mass stellar population (Higdon 1995, 1996; Mayya et al. 2005). Cartwheel is located in a small group of 4 galaxies. All of its 3 companions (known as G1, G2 and G3) are less massive than Cartwheel, and host less than 20 per cent of the total gas mass in the group. The analysis of X-ray data shows another intriguing peculiarity: most of the point sources detected by $Chandra$ are in the outer ring, and particularly concentrated in the southern quadrant (Gao et al. 2003; Wolter \& Trinchieri 2004). According to Wolter \& Trinchieri (2004) 13 out of 17 sources associated with Cartwheel are in the outer ring, the remaining 4 being close to the inner rim of the ring or to the optical spokes\footnote{The total number of point X-ray sources in the Cartwheel group is 24 (Wolter \& Trinchieri 2004), but 6 of them are associated with the companion galaxies G1 and G2, whereas 1 of them is probably background/foreground contamination.}. Gao et al. (2003) noted that all the five strongest H$\alpha{}$ knots in the ring are associated with an X-ray source, indicating a possible correlation between X-ray sources and young star-forming regions. Furthermore, most of the observed sources have isotropic X-ray luminosity $L_X\gtrsim{}10^{39}$ erg s$^{-1}$, belonging to the category of ultraluminous X-ray sources (ULXs). Given the distance of Cartwheel ($\sim{}124$ Mpc for an Hubble constant $H_0=73$ km s$^{-1}$ Mpc$^{-1}$), the data are incomplete for $L_X\lesssim{}5\times{}10^{38}$ erg s$^{-1}$, and the faintest sources in the sample have $L_X\sim{}10^{38}$ erg s$^{-1}$. Then, almost all the X-ray sources detected in Cartwheel are close to the ULX range. Many theoretical studies, both analytical (Lynds \& Toomre 1976; Theys \& Spiegel 1976; Appleton \& Struck-Marcell 1996) and numerical (Theys \& Spiegel 1976; Appleton \& Struck-Marcell 1987a, 1987b; Hernquist \& Weil 1993; Mihos \& Hernquist 1994; Struck 1997; Horellou \& Combes 2001; Griv 2005) were aimed to explain the origin of Cartwheel and its observational features. In the light of these papers, the origin of propagating rings in Cartwheel and similar galaxies can be explained by galaxy collisions with small impact parameter (Theys \& Spiegel 1976; Appleton \& Struck-Marcell 1987a, 1987b; Hernquist \& Weil 1993; Mihos \& Hernquist 1994; Struck 1997; Horellou \& Combes 2001). Among the companions of Cartwheel, both G1 and G3 are good candidates for this interaction, having a small projected impact parameter, being not too far from Cartwheel, and showing a disturbed distribution of neutral hydrogen (Higdon 1996). An alternative, less investigated model explains the rings with disc instabilities (Griv 2005). Models of galaxy collisions explain quite well most of Cartwheel properties. However, none of the previous theoretical studies has investigated the nature of the X-ray sources, and especially of the ULXs, observed in Cartwheel. The nature of the ULXs is still not understood. It has been suggested that they are high-mass X-ray binaries (HMXBs) powered by stellar mass black holes (BHs) with anisotropic X-ray emission (King et al. 2001; Grimm, Gilfanov \& Sunyaev 2003; King 2006) or with super-Eddington accretion rate (e.g. King \&{} Pounds 2003; Socrates \& Davis 2006; Poutanen et al. 2007). However, some ULXs, especially the brightest ones ($L_X\gtrsim{}10^{40}$ erg s$^{-1}$), show characteristics which are difficult to reconcile with the hypothesis of beamed emission, such as the presence of an isotropically ionized nebula (e.g. Pakull \& Mirioni 2003; Kaaret, Ward \& Zezas 2004) or of quasi periodic oscillations (e.g. Strohmayer \& Mushotzky 2003). Then, it has been proposed that some ULXs (or at least the brightest among them; Miller, Fabian \& Miller 2004) might be associated with intermediate-mass black holes (IMBHs), i.e. BHs with mass $20\lesssim{}m_{\rm BH}/M_\odot{}\lesssim{}10^{5}$, accreting either gas or companion stars (Miller et al. 2004; see Mushotzky 2004 and Colbert \& Miller 2005 for a review). On the other hand, most of ULXs can be explained with the properties of stellar mass BHs (see Roberts 2007 for a review and references therein). In this paper we present a new, refined $N-$body/SPH model of the Cartwheel galaxy, which reproduces the main features of Cartwheel. The aim of this paper is to use the $N-$body/SPH model in order to check whether the IMBH hypothesis is viable to explain all or a part of the X-ray sources observed in Cartwheel. In fact, in the last few years the hypothesis that most of the ULXs are powered by IMBHs has became increasingly difficult to support, as the observational features of the majority of ULXs are consistent with accreting stellar mass BHs (Roberts 2007). It would be interesting to see whether also the dynamical simulations agree with this conclusion\footnote{We will not consider other possible scenarios (e.g. the production of ULXs by beamed emission in HMXBs), due to intrinsic limits of $N-$body methods.}. In Section 2 we describe our simulations. In Section 3 we discuss the evolution of our models and the dynamics of either halo or disc IMBHs. In Section 4 we investigate the possibility that the X-ray sources in Cartwheel are powered by IMBHs accreting gas or stars. In the last case we consider both the accretion from old stars (which produce only transient sources) and from young stars, formed after the galaxy collision (which can lead also to persistent sources). We also investigate the hypothesis that new disc IMBHs are formed in very young clusters after the galaxy collision. Our findings are summarized in Section 5. | In this paper we investigated the possible connection among IMBHs and the $\sim{}17$ (Gao et al. 2003; Wolter \& Trinchieri 2004) bright X-ray sources detected in the outer ring of Cartwheel. Recent observations show that models based on beamed emission or super-Eddington accretion in HMXBs including stellar mass BHs can explain most of ULXs, apart from the brightest ones (Roberts 2007 and references therein). However, the observations cannot definitely exclude that all the ULXs are powered by IMBHs. Thus, in our paper we checked whether the IMBHs can account for all or only for a part of the ULXs observed in Cartwheel. We simulated the formation of a Cartwheel-like ring galaxy via dynamical interaction with an intruder galaxy. In this simulation we also integrated the evolution of 100 IMBHs particles. We considered two different models of IMBH formation, i.e. IMBHs born as relics of population III stars (and distributed as a concentrated halo population) and IMBHs formed via runaway collapse of stars (and distributed as an exponential disc). For these models, we investigated both gas accretion from surrounding molecular clouds and mass transfer from a stellar companion. The main results of this study are the following. \subsection{Halo IMBHs} IMBHs born as the relics of population III stars, if they are distributed as a halo population, cannot contribute to the X-ray sources, neither via gas accretion nor via mass transfer in binaries. In particular, the luminosity produced by halo IMBHs accreting gas is always many orders of magnitude smaller than that of the observed sources, even if we assume that halo IMBHs have a large mass ($m_{\rm BH}=10^3\,{}M_\odot{}$) and a high radiative efficiency ($\eta{}=0.1$). This is due to the fact that only a small fraction of halo IMBHs passes through the disc (only for a short lapse of time), and even these IMBHs have a high ($v\gtrsim{}100$ km s$^{-1}$) relative velocity with respect to gas particles. Similarly, halo IMBHs cannot accrete mass from stars, as the probability that they acquire a companion is very low. \subsection{Disc IMBHs} IMBHs born from the runaway collapse of stars should be a disc population. Under overoptimistic assumptions ($n_g=10^2\,{}{\rm cm}^{-3}$, $m_{\rm BH}=10^3\,{}M_\odot{}$ and $\eta{}=0.1$) these IMBHs can produce, via gas accretion, X-ray sources with $L_X\lesssim{}10^{39}$ erg s$^{-1}$, a factor of $\sim{}10$ fainter than the brightest ULXs observed in Cartwheel. Thus, also disc IMBHs accreting gas cannot explain the observed X-ray sources in Cartwheel. Our model of gas accreting IMBHs contains many rough assumptions; but most of them are upper limits, strengthening our result. However, a more realistic treatment of the local properties of the gas would be helpful, to understand the physical mechanisms of gas accretion onto IMBHs. On the other hand, runaway collapse born IMBHs are hosted in dense young clusters. In such environment, it is easy for the IMBH to capture a stellar companion. Blecha et al. (2006) estimated that a $\sim{}100\,{}M_\odot{}$ IMBH undergoes mass transfer from a companion star for about the $3$ per cent of the cluster lifetime. Previous papers (Portegies Zwart et al. 2004; Patruno et al. 2005) have shown that IMBHs accreting from low-mass ($\lesssim{}10\,{}M_\odot{}$) and high-mass ($\gtrsim{}10\,{}M_\odot{}$) companions generate transient (with a bright phase of only a few days) and persistent bright X-ray sources, respectively. As $10\,{}M_\odot{}$ stars have a lifetime of $\sim{}30-40\,{}$ Myr, only IMBHs hosted in sufficiently young star clusters can generate persistent X-ray sources. Then, IMBHs hosted in clusters born before the dynamical encounter with the intruder (i.e. more than 100 Myr ago) can produce only transient sources. We estimated that, out of 100 IMBHs which were present before the dynamical encounter with the intruder galaxy, only $\lesssim{}2-3$ are expected to undergo mass transfer from low-mass companions at present, producing a comparable number of transient X-ray sources. As observations show that at least 4 X-ray sources in the Cartwheel ring are persistent over a time-scale of 6 months (Wolter et al. 2006), we conclude that pre-encounter formed IMBH binaries are not sufficient to explain the data. We considered the possibility that pre-encounter IMBHs capture massive stars produced after the encounter with the intruder. In this case, under overoptimistic assumptions, 100 pre-encounter IMBHs can produce $\lesssim{}1$ X-ray source, either persistent or not. Finally, we hypothesized that very young ($<40$ Myr) star clusters, formed after the encounter, generate IMBHs at their centre. Under this hypothesis, $500-1000$ IMBHs are required to produce $\sim{}15-30$ bright ($10^{36}-10^{41}$ erg s$^{-1}$) X-ray sources, some of them persistent and some transient% . This scenario might account for the $\sim{}17$ observed X-ray sources in the Cartwheel ring. It is also in agreement with the fact that many ULXs observed in Cartwheel are associated with bright $H_\alpha{}$ spots, i.e. active star-forming regions (Gao et al. 2003). The birth of $\sim{}1000$ IMBHs (each one of 100 $M_\odot{}$) in $\sim{}$40 Myr implies an IMBH formation rate of $2.5\times{}10^{-3}\,{}M_\odot{}\,{}{\rm yr}^{-1}$, that is a factor of $\sim{}10^4$ lower than the SF rate. This rate is acceptable for runaway collapse scenarios, as Portegies Zwart \& McMillan (2002) show that $\sim{}$0.1 per cent of the mass of the parent young cluster merges to form the IMBH. However, we stress that only the runaway collapse scenario, under extreme assumptions, can explain the formation of such a huge number of IMBHs in the disc. Thus, our simulations suggest that IMBHs can hardly account for all the ULXs observed in Cartwheel. On the other hand, it is possible that only the few brightest sources in Cartwheel are powered by IMBHs, while the other ones are either beamed HMXBs, super-Eddington accreting stellar mass BHs or a blending of multiple fainter sources. For example, $\sim{}30$ IMBHs are expected to form in the ring, in order to produce only 1 very bright ULX, such as the source N.10 in Cartwheel. These results agree with the semi-analytical model by King (2004), who showed that IMBHs cannot explain all the ULXs in Cartwheel. However, King (2004) concludes that $>3\times{}10^4$ IMBHs are required to produce the observed number of ULXs, $\sim{}30-60$ times more than in our analysis. This apparent discrepancy is due to the fact that King (2004) assumes that the IMBHs power only transient sources (Kalogera et al. 2004), and thus he has to introduce a $\sim{}10^{-2}$ duty-cycle. However, Patruno et al. (2005) showed that IMBHs accreting from young massive stars ($\gtrsim10\,{}M_\odot{}$) produce non-transient sources, increasing the expected duty-cycle. In conclusion, new $Chandra$ and $XMM-Newton$ observations of Cartwheel could partially solve the mystery of Cartwheel X-ray sources, investigating which sources are transient, which variable, and which persistent. Deeper observations are also needed to resolve possible blended sources. In the future, it would be interesting to search whether other ring galaxies host as many bright X-ray sources as Cartwheel and whether these sources are similarly concentrated in the outer ring. | 7 | 10 | 0710.0780 |
0710 | 0710.0749_arXiv.txt | {}{We investigate the structure of a field around the position of V838~Mon as seen in the lowest CO rotational transitions. We also measure and analyse emission in the same lines at the position of V838~Mon.}{Observations have primarily been done in the $^{12}$CO $J = 2$$\to$$1$ and $J=3$$\to$$2$ lines using the KOSMA telescope. A field of 3.4 squared degrees has been mapped in the on-the-fly mode in these transitions. Longer integration spectra in the on-off mode have been obtained to study the emission at the position of V838~Mon. Selected positions in the field have also been observed in the $^{12}$CO $J = 1$$\to$$0$ transition using the Delingha telescope.}{In the observed field we have identified many molecular clouds. They can be divided into two groups from the point of view of their observed radial velocities. One, having $V_{\rm LSR}$ in the range $18-32\,{\rm km\,s}^{-1}$, can be identified with the Perseus Galactic arm. The other one, having $V_{\rm LSR}$ between $44-57\,{\rm km\,s}^{-1}$, probably belongs to the Norma-Cygnus arm. The radial velocity of V838~Mon is within the second range but the object does not seem to be related to any of the observed clouds. We did not find any molecular buble of a $1\degr$ dimension around the position of V838~Mon claimed in van~Loon et~al. An emission has been detected at the position of the object in the $^{12}{\rm CO}\ J=2$$\to$$1$ and $J=3$$\to$$2$ transitions. The emission is very narrow (FWHM~$\simeq$~1.2~km\,s$^{-1}$) and at $V_{\rm LSR} = 53.3$~km\,s$^{-1}$. Our analysis of the data suggests that the emission is probably extended.}{} | } The outburst of V838 Mon was discovered in the beginning of January~2002. Initially thought to be a nova, the object appeared unusual and enigmatic in its nature. The eruption, as observed in the optical, lasted about three months (Munari et~al. \cite{muna02}, Kimeswenger et~al. \cite{kimes02}, Crause et~al. \cite{crause03}). After developing an A--F supergiant spectrum at the maximum at the beginning of February~2002, the object evolved to lower effective temperatures and in April~2002 it practically disappeared from the optical, remaining very bright in the infrared. A detailed analysis of the evolution of the object in the outburst and decline can be found in Tylenda (\cite{tyl05}). Different outburst mechanisms have been proposed to explain the eruption of V838~Mon. They include an unusual nova (Iben \& Tutukov \cite{it92}), a late He-shell flash (Lawlor \cite{law05}) and a stellar merger (Soker \& Tylenda \cite{soktyl03}). These models have critically been discussed in Tylenda \& Soker (\cite{tylsok06}) and the authors conclude that the only mechanism that can satisfactorily account for the observational data is a collision and merger of a low-mass pre-main-sequence star with an $\sim$$8\,M_\odot$ main-sequence star. V838 Mon is surrounded by diffuse matter which gave rise to a spectacular light-echo phenomenon (e.g. Bond et~al. \cite{bond03}). Bond et~al. claim that the matter comes from previous eruptions of the object. However Tylenda (\cite{tyl04}), Crause et~al. (\cite{crause05}), as well as Tylenda, Soker \& Szczerba (\cite{tss05}) argue that the echoing matter is of interstellar origin. The latter is consistent with recent findings, namely that V838~Mon is a member of a young open cluster (Afsar \& Bond \cite{afsar}) and that the total mass of the echoing matter is probably of the order of $100\ M_\odot$ (Banerjee et~al. \cite{baner06}). van Loon et al. (\cite{loon04}) have analyzed archive infrared and radio data on the sky around V838~Mon and claimed discovery of multiple shells ejected by the object in the past. In particular, from a compilation of CO galactic surveys done in Dame et~al. (\cite{dame01}) van~Loon et~al. have suggested that V838~Mon is situated within a bubble of CO emission with a diameter of $\sim$$ 1\degr$. These results have been critically discussed in Tylenda et~al. (\cite{tss05}), who have concluded that the shells of van~Loon et~al. are not realistic. Deguchi et~al. (\cite{degu05}) have discovered an SiO maser emission from V838~Mon. The main component is narrow and centered at $V_{\rm LSR} \simeq 54\ {\rm km}\,{\rm s}^{-1}$, which is thought to be a radial velocity of the object itself. Further observations of Claussen et~al. (\cite{clauss06}) have shown that the SiO maser is variable and that most of the emission comes from a region smaller than a milliarcsecond. In the present paper we report on results of our observations of V838~Mon and its nearby vicinity in the $^{12}$CO $J = 1$$\to$0, 2 $\to$1 and 3$\to$2 transitions. We describe the observational material and discuss results on the CO emission from a field of 3.4 squared degrees around the position of V838~Mon. One of the goal of this survey is to verify the existence of the CO bubble around V838~Mon claimed in van~Loon et~al. (\cite{loon04}). Measurements of the CO emission obtained at the position of V838~Mon are also presented, analysed and discussed. A preliminary analysis of the data has been done in Kami\'{n}ski et~al. (\cite{kmst}). | } Our on-the-fly maps in the CO rotational transitions, especially those done in the (2$\to$1) line with the KOSMA telescope, have allowed us to identify 29 molecular clouds around the position of V838~Mon (see Appendix~\ref{append}). As far as we know, the region mapped with the KOSMA telescope has not been observed before with a sensitivity and a spatial resolution comparable to our survey. The data presented in Appendix~\ref{append} can be used in studies of molecular matter in the outer Galaxy. We do not confirm the existence of a molecular bubble of a 1 degree dimension claimed in van~Loon et~al. (\cite{loon04}), which was probably an artefact resulted from merging two surveys of different resolution and sensitivity in the data used by van~Loon et~al. Our maps did not reveal any molecular cloud in the near vicinity of V838~Mon. The nearest CO emission has been detected at $\sim$$40$~arcmin from the position of the object (see Fig.~\ref{map_2}). Thus the nearest dense molecular cloud is located at least at a distance of $\sim$$80$~pc from V838~Mon (assuming a 7~kpc distance to V838~Mon -- see below). From the noise level in the (2$\to$1) map (see Table~\ref{tech_tab}) we can put an upper limit to $I_{\rm CO}=\int T_{\rm mb}\,dV$ of $1.26~{\rm K~km~s}^{-1}$ ($3 \sigma_{\rm rms}$) at the position of V838~Mon. Assuming the $N$(H$_2$) to $I_{\rm CO}$ conversion factor of $X_{\rm CO} = 5.1$ (in units $10^{20}\,{\rm molecules}\ {\rm cm}^{-2}\,{\rm K}^{-1}\,{\rm km}^{-1}\,{\rm s}$, typical value for the NC clouds discussed in Appendix~\ref{append}) we obtain an upper limit to the column density of $N({\rm H}_2) = 6.4 \times 10^{20}~{\rm cm}^{-2}$. This upper limit is an order of magnitude lower than typical column densities observed for molecular clouds in the vicinity of the Sun (e.g. Blitz \cite{blitz}). Thus we can conclude that V838~Mon is not located in a typical molecular cloud. However, as presented in Sect.~\ref{kosma}, long integrations in the on-off mode have allowed us to detect an emission in the CO (2$\to$1) and (3$\to$2) transitions at the position of V838~Mon. The question that arises is: what is the origin of this emission? Certainly it cannot come from the matter ejected in the 2002 outburst. That matter was expanding with large velocities, a few hundred km~s$^{-1}$, while our profiles have a FWHM of $\sim$$1\,{\rm km\,s}^{-1}$. The $V_{\rm LSR}$ position and width of the CO lines similar to those of the SiO maser (Deguchi et~al. \cite{degu05}, Claussen et~al. \cite{clauss06}) suggest that both emissions (CO and SiO) originate from the same place, e.g from the remnant of the 2002 outburst. However, in this case the KOSMA telescope would see a point source, which does not seem to be the case. The April 2006 observations show detectable emissions at the 1~arcmin offsets from the V838~Mon position (see Table~\ref{res2_tab}). The HPBW of the KOSMA beam at 230~GHz is 130~arcsec. If the CO emission source were a point source at the (0,0) position, than the intensity measured at the (0,1) and ($0,-1$) offsets would be about twice fainter than that at the central position. Table~\ref{res2_tab} shows that this ratio is larger, i.e. $0.6-0.8$. Unfortunantely the accuracy of the measurements at the 1~arcmin offsets was not high, so we cannot say that the results in Table~\ref{res2_tab} are definitively inconsistent with a point source emission. However, there are other arguments in favour of the idea that the CO emission is extended and/or situated outside the V838~Mon position. Our Delingha observations in the (1$\to$0) transition (see Sect.~\ref{delingha}) have not detected any emission at the V838~Mon position and the upper limit was 0.39~K. Assuming that the intensity in the (1$\to$0) line is comparable to that in the (2$\to$1) line (which is usually the case in molecular clouds) and that we observe a point source at the V838~Mon position than from the measured $T_{\rm mb}$ in the (2$\to$1) transition in December~2005 (see Table~\ref{res_tab}), taking into account different beamsizes of the two telescopes (55~arcsec in Delingha versus 130~arcsec in KOSMA), the expected value of $T_{\rm mb}$ in the (1$\to$0) Delingha observations would be $\sim$$1.8$\,K. This is 4.5 times higher than the observed upper limit. Thus either the (2$\to$1)/(1$\to$0) ratio is exceptionally large ($> 4.5$) or the CO source is situated outside the Delingha beam but still inside the KOSMA beam. The later interpretation is supported by our possible detection of an emission with Delingha at three positions $\sim$$1$\,arcmin from the object (see Table~\ref{resd_tab}). It is also supported by a finding of Deguchi et~al. (\cite{degu06}) who, using the Nobeyama telescope, tried to measure the CO (1$\to$0) emission from a few points in a field around V838~Mon. A signal was detected from a position 30~arcsec north of V838~Mon at a velocity very close to that of the (2$\to$1) line in our KOSMA observations. No emission was however recorded at the position of the object. The above discussion allows us to conclude that the CO emission, clearly seen in our KOSMA observations, most probably originates not from V838~Mon itself but from a region (regions) situated, typically, $\sim$$1$\,arcmin from the V838~Mon position. There are two possible ways of explaining the emission in this case. One is that the observed emission is a faint part of a larger CO structure belonging, for instance, to a molecular cloud. However, as discussed above, our CO maps have not revealed any dense CO cloud of similar $V_{\rm LSR}$ closer than $\sim$$40$\,arcmin from the position of V838~Mon. All the detected and possibly detected CO emission at and near the position of V838~Mon lie well within the optical light echo of V838~Mon. Hence the second possibility, namely that the CO emission comes from the same matter that gave rise to the light echo and that it was the light flash from the 2002 eruption which induced the emission. Then the observed narrowness of the line profile would be a strong argument that the matter is of interstellar origin rather than being ejected by V838~Mon in previous eruptions. The observed $V_{\rm LSR} = 53.3\,{\rm km\,s}^{-1}$ of the CO emission would imply a kinematical distance of $\sim$$7.0$\,kpc (using the Galactic rotational curve of Brandt \& Blitz \cite{bb93}). This can be compared to $6.1 - 6.2$~kpc found by Bond et~al. (\cite{bond06}) and $\sim$$8$~kpc advocated in Tylenda (\cite{tyl05}). Rushton et al. (\cite{rush03}) searched for CO emission from V838~Mon about a year after the 2002 eruption. No measurable signal was detected in the three lowest rotational transitions and the upper limit was $T_A^* \la 25-40$~mK. It is not straightforward to compare this result with our findings as the observations of Rushton et~al. were done about 3~years before ours and the object probably evolved significantly during this time span. Nevertheless, given the beamwidth of the telescopes used by Rushton et~al. (HPBW~$\le 45\,\arcsec$), their negative result at the position of the object does not preclude a possibility that a significant emission was present in a near vicinity of the object, but outside the telescope beam. More sensitive observations with a higher spatial resolution are required to distinguish between the above discussed possibilities and to futher investigate the problem of the CO emission from V838~Mon. | 7 | 10 | 0710.0749 |
0710 | 0710.5393_arXiv.txt | The Zeeman pattern of Mn~{\sc i} lines is sensitive to hyperfine structure (HFS) and, because of this reason, they respond to hG magnetic field strengths differently from the lines commonly used in solar magnetometry. This peculiarity has been employed to measure magnetic field strengths in quiet Sun regions. However, the methods applied so far assume the magnetic field to be constant in the resolution element. The assumption is clearly insufficient to describe the complex quiet Sun magnetic fields, biasing the results of the measurements. The diagnostic potential of Mn~{\sc i} lines can be fully exploited only after understanding the sense and the magnitude of such bias. We present the first syntheses of Mn~{\sc i} lines in realistic quiet Sun model atmospheres. Plasmas varying in magnetic field strength, magnetic field direction, and velocity, contribute to the synthetic polarization signals. The syntheses show how the Mn~{\sc i} lines weaken with increasing field strength. In particular, kG magnetic concentrations produce \mni{5538} circular polarization signals (Stokes~$V$) which can be up to two orders of magnitude smaller than what the weak magnetic field approximation predicts. As corollaries of this result, (1) the polarization emerging from an atmosphere having weak and strong fields is biased towards the weak fields, and (2) HFS features characteristic of weak fields show up even when the magnetic flux and energy are dominated by kG fields. For the HFS feature of \mni{5538} to disappear the filling factor of kG fields has to be larger than the filling factor of sub-kG fields. Since the Mn~{\sc i} lines are usually weak, Stokes~$V$ depends on magnetic field inclination according to the simple consine law, a scaling that holds independently of the magnetic field strength. Atmospheres with unresolved velocities produce very asymmetric line profiles, which cannot be reproduced by simple one-component model atmospheres. Inversion techniques incorporating complex magnetic atmospheres must be implemented for a proper diagnosis. Using the HFS constants available in the literature we reproduce the observed line profiles of nine lines with varied HFS patterns. Consequently, the uncertainty of the HFS constants do not seem to limit the use of Mn~{\sc i} lines for magnetometry. | \label{intro} When the polarimetric sensitivity and the angular resolution exceed the required threshold, magnetic signals show up almost everywhere on the solar photosphere. The signals in supergranulation cell interiors are particularly weak, however, most of the solar surface is in the form of cell interiors and, therefore, these weak signals probably set the total unsigned magnetic flux and magnetic energy existing in the photosphere at any given time \citep[e.g.,][]{unn59,ste82,yi93,san98c,san04,sch03b}. The importance of these ubiquitous quiet Sun magnetic fields depends, to a large extent, on the magnetic field strengths characterizing them. For example, the magnetic flux and the magnetic energy scale as powers of the field strength, and the connectivity between photosphere and corona is probably provided by the magnetic concentrations with the largest field strengths \citep[e.g.,][]{dom06}. Unfortunately, measuring quiet Sun magnetic field strengths is not a trivial task. All measurements are based on the residual polarization left when a magnetic field of complex topology is observed with finite angular resolution \citep[e.g.,][]{emo01,san03,tru04}. Measuring implies assuming many properties of the unresolved complex magnetic field and, in doing so, the measurements become model dependent and prone to bias. Depending on the technique used for measuring, the real quiet Sun exhibits weak daG fields \citep[e.g.,][]{ste82,fau93,bom05}, intermediate hG fields \citep[e.g.,][]{lin99,kho02,lop06}, or strong kG fields \citep[e.g.,][]{san00,lit02,dom03b}. Such discrepancies can be naturally understood if the true quiet Sun contains a continuous distribution of field strengths going all the way from zero to two kG. Different techniques are biased differently and, therefore, they tend to pick out a particular part of the distribution. This scenario is very much consistent with realistic numerical simulations of magneto-convection \citep[][]{cat99a,stei06,vog05,vog07}. In order to provide a comprehensive observational description of the quiet Sun magnetic field strengths, one has to combine different methods carefully chosen to have complementary biases. Then the full distribution can be assambled taking the biases into account \cite[][]{dom06}. In an effort to complement the existing magnetic field strength diagnostic techniques, \citet{lop02} proposed the use of spectral lines whose Zeeman patterns are sensitive to hyperfine structure (HFS). The formalism to deal with the HFS of spectral lines in magnetic atmospheres was developed more than thirty years ago by \citet{lan75}. According to such formalism, the polarization of the HFS lines vary with magnetic field strength very differently from the lines commonly used in solar magnetometry. This unusual behavior was invoked by \citet{lop02} when proposing the use of HFS Mn~{\sc i} lines as a diagnostic tool for magnetic field strengths. L\'opez Ariste and coworkers have applied the idea to measure magnetic field strengths in quiet Sun regions \citep{lop06,lop07,ase07}. Since the number of observables is limited, they minimize the statistical error of the measuremet by minimizing the number of free parameters to be tuned. Milne-Eddington atmospheres (ME) are used to synthesize the polarization of the lines. The magnetic field is assumed to be constant and, therefore, the measurements provide some kind of weighted average of the true field strengths existing in the resolution elements. As we pointed out above, the topology of the quiet sun magnetic field is complex, with a distribution of field strengths and polarities coexisting in a typical resolution element. Then the ill-defined average provided by the Mn~{\sc i} lines is expected to be biased toward a particular range of field strengths, as it happens with the rest of the quiet Sun measurements (e.g., the NIR Fe~{\sc i} lines overlook kG magnetic concentrations, \citealt{san00}, \citealt{soc03}; the traditional visible Fe~{\sc i} lines exaggerate the contribution of kG fields, \citealt{bel03}, \citealt{mar06}; the Hanle scattering polarization signals are not sensitive to hG and kG, \citealt{fau01}, \citealt{san05}). The bias presented by the HFS Mn~{\sc i} lines is so far unknown, and the existing and forthcoming measurements based on those lines will be fully appreciated only when the sources of systematic effects are properly acknowledged and quantified. In order to explore the magnitude and the sense of the expected effects, we undertake the synthesis of Mn~{\sc i} lines in a number of realistic quiet Sun atmospheres with complex magnetic field distributions. The main trends are presented here. The paper is organized as follows. Section~\ref{code} describes the software developed to carry out the syntheses. First, a ME code is needed to compare our syntheses with the results in the literature. Then, a plane parallel one-dimensional (1D) code allows us to explore the influence of realistic thermodynamic conditions on the polarization of lines with HFS. Finally, a MIcro-Structured Magnetic Atmosphere (MISMA) code provides additional realism to the modeling since it includes coupling between magnetic field strengths and thermodynamic conditions, magnetic fields varying along and across the line-of-sight (LOS), mixed polarities in the resolution element, etc. All these codes are based on the original FORTRAN routine written by \citet{lan78}. The analysis is focused on the line most often used in observations, namely, \mni{5538}. We discuss its intensity and circular polarization, since the linear polarization signals are very weak and they remain undetected so far. Single component and multi-component syntheses of this line are described in \S~\ref{scs} and \S~\ref{mcs}, respectively. An exploratory attempt to consider unresolved mixed polarities and velocity fields is carried out in \S~\ref{comp_obs}. The polarization of other Mn~{\sc i} lines is also considered in \S~\ref{other}. The implications of these syntheses are discussed in \S~\ref{conclusions}. | The Zeeman pattern of the Mn~{\sc i} lines depends on hyperfine structure (HFS), which confers them a sensitivity to hG magnetic field strengths different from the lines traditionally used in solar magnetometry. This peculiarity has been used to measure magnetic field strengths in quiet Sun regions (see \S~\ref{intro}). The methods applied so far assume the magnetic field strength to be constant in the resolution element, an approximation driven by feasibility rather than based on physical or observational arguments. Actually, it is not a good approximation since the magnetic fields of the quiet Sun are expected to vary on very small spatial scales, with field strengths spanning from zero to 2kG. Under these extreme conditions, all diagnostic techniques employed in quiet Sun magnetometry are strongly biased towards a particular range of field strengths, and the Mn~{\sc i} signals are not expected to be the exception. Consequently, the diagnostic content of the Mn~{\sc i} lines cannot be fully exploited unless their biases are properly understood. Such task is undertaken in the paper by exploring the response of Mn~{\sc i} lines in a number of realistic quiet Sun scenarios. Three complementary LTE synthesis codes have been written and tested (ME, 1D and MISMA; \S~\ref{code}). They provide a number of relevant results, the first one being the ability to reproduce all observed Mn~{\sc i} HFS patterns. We reproduce the observed unpolarized line profile of nine assorted lines corresponding to all HFS sensitivities (\S~\ref{other}). The study is focused on the line most often used in magnetometry, \mni{5538}, however its behavior should be representative of the other lines. According to the weak magnetic field approximation, the Stokes $V$ signals scale with the longitudinal component of the magnetic field, i.e., the magnetic field strength times the cosine of the magnetic field inclination with respect to the LOS. We verify that the scaling on cosine holds very tightly. The scaling on the magnetic field strength, however, breaks down soon (when $B \ge 400$~G for \mni{5538}). Even for ME atmospheres, the weak field approximation predicts a Stokes~$V$ signal twice as large as the synthetic one for $B\simeq$1.5~kG. When the expected coupling between the thermodynamic conditions and the magnetic field strengths is taken into account, the dimming of the kG Stokes~$V$ signals can be as large as two orders of magnitude (see Figs.~\ref{maxv} and \ref{phi}). The dimming of the polarization signals formed in kG magnetic concentrations affects all Mn~{\sc i} lines (\S~\ref{scs}), and it has significant observational implications. If the resolution elements contains both weak and strong fields, then the kG fields tend to be under-represented in the average profile. We have modelled the bias assuming a multi component atmosphere, where the synthetic signals are weighted means of the Stokes profiles corresponding to each single field strength. The weight is given by the fraction of atmosphere filled by each field strength, i.e., by the magnetic field strength probability density function PDF (equations~[\ref{pdfIV}] and [\ref{pdfIVb}]). According to our modeling, even when the (unsigned) magnetic flux and the magnetic energy are dominated by kG magnetic fields, the Stokes~$V$ profile of \mni{5538} can show HFS reversal at line core characteristic of hG fields (Fig.~\ref{vsme}). A pure morphological inspection of the Stokes profiles does not suffice to infer which is the dominant magnetic field strength in the resolution element. For the HFS hump of \mni{5538} to disappear the kG filling factor has to be larger than the sub-kG filling factor and, consequently, when the HFS hump disappears the magnetic flux and magnetic energy of the atmosphere are completely dominated by kG fields (\S~\ref{proxy}). Detecting Stokes~$V$ profiles with HFS features indicates the presence of hG fields in the resolution element. However, this sole observation does not tell whether the hG field strengths dominate. There seem to be two extreme alternatives to exploit the diagnostic potential of these lines. First, improving the spatial resolution of the observations to a point where the quiet Sun magnetic structures can be regarded as spatially resolved. Unfortunately, this possibility does not seem to feasible at present. Realistic simulations of magneto-convection indicate that quiet Sun magnetic fields are uniform only at the diffusive length scales \citep[e.g.][]{cat99a,vog07}, which are of the order of a few km in the photosphere \citep[e.g.][]{sch86}. These scales are very far from the angular resolution of the present measurements, and even much smaller than the length-scale for the radiative transfer average along the LOS \citep{san96}. We prefer the alternative possibility, namely, developing inversion techniques where complex magnetic atmospheres are included into the diagnostics. Using appropriate constraints, the number of free parameters required to describe such atmospheres can be maintained within reasonable limits \citep[e.g., the model MISMAs in ][]{san97b}. Dealing with unresolved velocities also favors detailled inversion codes (\S~\ref{comp_obs}). We are presently working on these improvements needed to develop the diagnostic technique pioneered by \citeauthor{lop02}. | 7 | 10 | 0710.5393 |
0710 | 0710.3326_arXiv.txt | {I provide a short review of the properties of Narrow-line Seyfert 1 (NLS1) galaxies across the electromagnetic spectrum and of the models to explain them. Their continuum and emission-line properties manifest one extreme form of Seyfert activity. As such, NLS1 galaxies may hold important clues to the key parameters that drive nuclear activity. Their high accretion rates close to the Eddington rate provide new insight into accretion physics, their low black hole masses and perhaps young ages allow us to address issues of black hole growth, their strong optical FeII emission places strong constraints on FeII and perhaps metal formation models and physical conditions in these emission-line clouds, and their enhanced radio quiteness permits a fresh look at causes of radio loudness and the radio-loud radio-quiet bimodality in AGN. } \resumen{Favor de proporcionar un resumen en espa\~nol. \\ } \addkeyword{Galaxies: Active} \addkeyword{Galaxies: Seyfert} \begin{document} | \label{sec:intro} Narrow-line Seyfert 1 galaxies are a subclass of active galactic nuclei (AGN). Their spectra exhibit exceptional emission-line and continuum properties. The most common NLS1 defining criterion is the width of the broad component of their optical Balmer emission lines in combination with the relative weakness of the [OIII]$\lambda$5007 emission (FWHM$_{\rm H\beta} < 2000$ km/s and [OIII]/H$\beta_{\rm totl}$ $<$ 3; Osterbrock \& Pogge 1985, Goodrich 1989){\footnote{ While it is clear that a strict cutoff in line width (FWHM$_{\rm H\beta} < 2000$ km/s) is a gross simplification of any classification scheme, this historical value is still most commonly adopted for practical purposes. Suggestions have been made that more advanced NLS1 classification schemes would, for instance, incorporate the source luminosity (e.g., Laor 2000, Veron-Cetty et al. 2001). According to Sulentic et al. (2008 and references therein), AGN properties appear to change more significantly at a broad line width of FWHM$_{\rm H\beta} \approx 4000$ km/s.}}. NLS1 galaxies typically show strong FeII emission which anticorrelates in strength with the [OIII] emission, and with the width of the broad Balmer lines. Often the presence of FeII emission is added as further NLS1 classification criterion and Veron et al. (2001) suggest the use of an intensity ratio FeII/H$\beta_{\rm totl} > 0.5$. NLS1 galaxies as AGN with the smallest Balmer lines from the Broad Line Region (BLR) and the strongest FeII emission, cluster at one extreme end of AGN correlation space. It is expected that such correlations provide some of the strongest constraints on, and new insights in, the physical conditions in the centers of AGN and the prime drivers of activity, and the study of NLS1 galaxies is therefore of particular interest. For instance, observations and interpretations hint at smaller black hole masses in NLS1 galaxies, and as such their black holes represent an important link with the elusive intermediate mass black holes, which have been little studied so far. Accreting likely at very close to the maximum allowed values, NLS1 galaxies are important test-beds of accretion models. This paper provides a short overview of the multi-wavelength properties of NLS1 galaxies and major models to explain them. | -- {\sl{It seems hard to resist the feeling that nature is telling us something important here, but we do not yet know what it is.}} ~~~~~~~~~~~~~~Lawrence et al. (1997) \vskip0.3cm While our knowledge has increased significantly in the last decade, important questions are still open. For instance, what are sufficient, what are necessary conditions for the onset of NLS1 activity ? For instance, the question is raised whether there are two types of Seyfert galaxies with low black hole masses: there is the unavoidable low-black-hole mass extension of BLS1 galaxies. Such systems would have their FWHM$_{\rm H\beta}$ fall below the formal cutoff value of 2000 km/s. Does low black hole mass already imply the emergence of some or all of the typical observed NLS1 characteristics ? Or is there a separate class of NLS1 galaxies ? While many individual NLS1 galaxies have been studied in great detail, we still need larger samples free of selection biases and well-suited BLS1 comparison samples in order to identify robust trends. Correlation space will ultimately have to be expanded to include the radio and infrared properties of NLS1 galaxies, as well as the properties of their host galaxies. On the theoretical/modeling side, interesting questions that persist or have emerged are related to mechanisms of super-Eddington accretion, the simultaneous presence of (TeV) blazar activity and high accretion rate in extreme radio loud NLS1 galaxies, and mechanisms of fueling and feedback in NLS1 galaxies. The study of NLS1 galaxies will continue to provide important contributions to our understanding of AGN and their cosmic evolution. | 7 | 10 | 0710.3326 |
0710 | 0710.1609_arXiv.txt | We examine the early phases of two near-limb filament destabilizations involved in coronal mass ejections on 16 June and 27 July 2005, using high-resolution, high-cadence observations made with the Transition Region and Coronal Explorer (TRACE), complemented by coronagraphic observations by Mauna Loa and the SOlar and Heliospheric Observatory (SOHO). The filaments' heights above the solar limb in their rapid-acceleration phases are best characterized by a height dependence $h(t)\propto t^m$ with $m$ near, or slightly above, 3 for both events. Such profiles are incompatible with published results for breakout, MHD-instability, and catastrophe models. We show numerical simulations of the torus instability that approximate this height evolution in case a substantial initial velocity perturbation is applied to the developing instability. We argue that the sensitivity of magnetic instabilities to initial and boundary conditions requires higher fidelity modeling of all proposed mechanisms if observations of rise profiles are to be used to differentiate between them. The observations show no significant delays between the motions of the filament and of overlying loops: the filaments seem to move as part of the overall coronal field until several minutes after the onset of the rapid-acceleration phase. | Observations of the early rise phase of filaments and their overlying fields can in principle help constrain the mechanisms involved in the destabilization of the magnetic configuration through comparison with numerical simulations (e.g., \citep{fan2005}; \citep{torok+kliem2005}; \citep{williams+etal2005}; and references therein), because the detailed evolution depends sensitively on the model details. For example, a power-law rise with an exponent $m=2.5$ was obtained for a slender flux tube in the two-dimensional version of the catastrophe model (\citep{Priest&Forbes02}). An MHD instability triggered by an infinitesimal perturbation implies an exponential rise, as was verified, for example, for a three-dimensional flux rope subject to a helical kink instability (\citep{Torok&al04}, \citep{torok+kliem2005}). The same holds for the torus (expansion) instability (TI), which starts as a $\sinh(t)$ function (\citep{Kliem&Torok06}) that is very similar to a pure exponential early on. The CME rise in a breakout model simulation was well described by a parabolic profile (\citep{Lynch&al04}). The early rise phase of erupting filaments is best observed near the solar limb using high-resolution data, both in space and in time. Such data can be obtained by, for example, Big Bear Solar Observatory H$\alpha$ observations (e.g., \citep{kahler+etal1988}), the Mauna Loa K-coronameter (e.g., \citep{gilbert+etal2000}), the Nobeyama Radioheliograph (e.g., \citep{gopalswamy+etal2003}; \citep{kundu+etal2004}), and the Transition Region and Coronal Explorer, TRACE (e.g., \citep{vrsnak2001}; \citep{gallagher+etal2003}; \citep{goff+etal2005}; \citep{sterling+moore2004}; \citep{sterling+moore2005}; \citep{williams+etal2005}). In those few cases where observers had the field of view for an appropriate diagnostic to attempt to establish whether the high loops or the filaments were accelerated first, the temporal resolution often was not adequate (see, e.g., \citep{sterling+moore2004}, who use the standard 12-min.\ cadence of SOHO/EIT). These studies show that filaments that are about to erupt often --~but not always~-- exhibit a slow initial rise during which both the filament and the overlying field expand with velocities in the range of $1-15$\,km/s. Then follows a rapid-acceleration phase during which velocities increase to a range of $100$\,km/s up to over $1000$\,km/s. The rapid-acceleration phase finally transitions into a phase with a nearly constant velocity or even a deceleration into the heliosphere. The height evolution immediately following the onset of the rapid acceleration phase is often approximated by either an exponential curve (e.g., \citep{gallagher+etal2003}; \citep{goff+etal2005}; \citep{williams+etal2005} --~who also show systematic deviations from that fit up to 2$\sigma$ in position~-- ) or by a constant-acceleration curve (e.g., \citep{kundu+etal2004}; and \citep{gilbert+etal2000} --~ who show one case in which a third-order curve improves the fit to the earliest evolution, and leave others for future analysis); \citet{kahler+etal1988} fit curves for the acceleration $a=ct^b$ to the first $10-50$\,Mm for four erupting filaments, but do not list the best-fit values. \cite{alexander&al02} find a best fit for the height of the early phase of a CME observed in X-rays by YOHKOH's SXT of the form $h_0+v_0t+ct^{3.7\pm0.3}$. For 184 prominence events observed by the Nobeyama Radioheliograph, \citet{gopalswamy+etal2003} show that higher in the corona velocity profiles include decelerating, constant velocity, and accelerating ones for heights from $\sim 50$\,Mm to 700\,Mm above the solar surface. In many cases, the detailed study of the evolution of the early phase is hampered by insufficient temporal coverage or by gaps between the fields of view of two complementing instruments that can be as large as a few hundred Mm. This results in substantial uncertainties in the height evolution. \citet{vrsnak2001}, for example, concludes that ``[t]he main acceleration phase \ldots\ is most often characterized by an exponential-like increase of the velocity'', but notes that polynomial or power-law functions fit at comparable confidence levels. In this study, we examine two events displaying the early destabilization and acceleration of ring filaments leading to coronal mass ejections. The high cadence down to 20\,s, and the high spatial resolution of 1\,arcsec, for the early evolution result in relatively small uncertainties in the height profiles. This enables a sensitive test of the height evolution against exponential, parabolic, and power-law fits. We find that a power-law with exponent near $3$, or slightly higher, is statistically preferred in both cases. As no published model matches that profile, we experiment with a numerical model for the torus instability, and find that this model can indeed approximate the observations provided that a sufficiently large initial velocity perturbation is applied (without which an exponential-like profile would be found). This finding reminds us of the sensitivity of developing instabilities to both initial and boundary conditions, and shows that the models, particularly their parametric dependencies, need to be worked out in greater detail in order to use observations of the height-time observations to differentiate successfully between competing models. | \label{sec:conclusions} We study two well-observed filament eruptions, and find that their rapid acceleration phases are well fit by a cubic height-time curve that implies a nearly constant jerk for $10-15$\,minutes, followed by a transition to a terminal velocity of $\sim 750$\,km/s and $\sim 1250$\,km/s, respectively. Simulations of a torus instability (TI) can reproduce such a behavior, provided that a substantial initial velocity perturbation is introduced. Without that perturbation, an exponential rise profile would be found. We note that the initial slow rise and the onset of the subsequent rapid acceleration phase are shared between the filament and overlying loop structures: neither \referee{leads the other to within the temporal resolution. For characteristic Alfv{\'e}n speeds over active regions of $\sim 1,000$\,km/s, the propagation of a perturbation over the separation of $\sim 75,000$\,km would require only $\sim 1.2$\,min., which corresponds to only one or two exposures. Thus the observations allow for Alfv{\'e}nic propagation of a signal between filament and overlying loops, but suggest no longer-term differential evolution.} We observe no significant changes in the separation of erupting filament and overlying loops within that interval (Figs.~\ref{fig:1c} and~\ref{fig:2c}). After that, the distance increases in the 2005/06/16 eruption, suggesting the overlying field moves to the side for some time faster than the filament rises. For the 2007/07/27 eruption, the distance stays the same for one loop and decreases for another for up to 10~min after the start of the rapid acceleration phase, which reflects the significant sideways motion component of the rising filament. The observed configuration of the filament and high loops may be part of a larger overall destabilizing field configuration. Our numerical modeling has assumed that, in the rapid acceleration phase, the overlying field starts to move rapidly only as a consequence of the flux rope's destabilization. This is consistent with the data. However, we cannot exclude that the filament and the overlying field were destabilized simultaneously by a process different from the one considered here. More study is needed to establish whether the common evolution of the filament and high loops has a significant diagnostic value as to the cause of the instability. Comparison with other model studies in the literature leads us to conclude that the catastrophe model and the TI model are both marginally consistent with the observations of the two erupting filaments. The catastrophe model predicts a power-law exponent near the lower edge of the range of acceptable fits, but we have to allow for the possibility that changing that model's details may change the acceleration profile. \referee{In order to yield the observed nearly cubic power-law rise (with $m$ slightly exceeding 3), our TI model requires an initial perturbation velocity that is in agreement with the observed rise velocity at the onset of the rapid-acceleration phase. If a nearly exact cubic rise were to be matched, however, initial velocities moderately exceeding the observed ones, by a factor $\approx1.5$, were required. In any case, our modeling is consistent with the observed velocities after the first few minutes of the eruption.} Having established that the model for the TI instability is very sensitive to the initial conditions, we should of course also acknowledge that it depends sensitively on the model details itself. These include the details of the external field and of the rates and locations of the reconnection that occurs behind the erupting filament. That such reconnection occurs in reality is suggested for both events by the occurrence of brightenings mainly at the bottom side of the filaments at the onset of the rapid-acceleration phase. These brightenings develop later into the streaks used for position determination in Sect.~\ref{sec:Observations}. The onset of reconnection even before the rapid-acceleration phase of the filament eruption on 27 July 2005 is strongly suggested by precursor soft and hard X-ray emission during about 04:00--04:30~UT, whose analysis revealed heating to 15~MK and the acceleration of non-thermal electrons to energies $>10$~keV (\citep{Chifor&al06}). The observed rise velocity early in the filament eruption may be an underestimate of the true expansion velocity of the hoop formed by the flux rope: the filament channel in the pre-eruption phase of AR\,10775 is strongly curved, and one of the two possible channels in AR\,10792 is too (ambiguity exists here because the eruptions occurred very near the limb, so that the configurations of the filament channels can only be observed some days before and after the events, respectively). If the initial expansion of the flux rope would have a strong component in the general direction of the inclined plane of the curved filament channel rather than be purely normal to the solar surface, projection effects could cause us to underestimate the expansion velocity in particular early in the evolution. In addition to that, we must realize that the TI model assumes a flux rope that stands normal to the solar surface and that erupts radially. Future more detailed modeling will have to show how deviations from that affect the evolution of the eruption. The fact that the torus-instability model yields qualitatively different rise profiles (exponential vs.\ power law) in different parts of parameter space, cautions against expectations that precise measurements of the rise profile of filament eruptions by themselves permit a determination of the driving process: the non-linearities in the eruption models clearly require high-fidelity modeling if such observations are to be used to differentiate successfully between competing models. Our initial modeling discussed here suggests that the torus instability is a viable candidate mechanism for at least some filament eruptions in coronal mass ejections. \referee{Given the dependence of nonlinear models on the details of boundary and initial conditions, it will be necessary to investigate how other models for erupting filaments compare to the data, as well as how the fidelity of our modeling of the torus instability can be improved before we can reach definitive conclusions about the mechanism(s) responsible for filament eruptions in general.} | 7 | 10 | 0710.1609 |
0710 | 0710.4506_arXiv.txt | The hydrogen-deficiency in extremely hot post-AGB stars of spectral class PG1159 is probably caused by a (very) late helium-shell flash or a AGB final thermal pulse that consumes the hydrogen envelope, exposing the usually-hidden intershell region. Thus, the photospheric elemental abundances of these stars allow to draw conclusions about details of nuclear burning and mixing processes in the precursor AGB stars. We compare predicted elemental abundances to those determined by quantitative spectral analyses performed with advanced non-LTE model atmospheres. A good qualitative and quantitative agreement is found for many species (He, C, N, O, Ne, F, Si, Ar) but discrepancies for others (P, S, Fe) point at shortcomings in stellar evolution models for AGB stars. PG1159 stars appear to be the direct progeny of [WC] stars. | The PG1159 stars are a group of 40 extremely hot hydrogen-deficient post-AGB stars. Their effective temperatures ($T_{\rm eff}$) range between 75\,000 and 200\,000~K. Many of them are still heating up along the constant-luminosity part of their post-AGB evolutionary path in the HR diagram ($L \approx 10^4$L$_\odot$) but most of them are already fading along the hot end of the white dwarf cooling sequence (with $L$ {\raisebox{-0.4ex}{${\stackrel{>}{\scriptstyle \sim}} $}} 10\,L$_\odot$). Luminosities and masses are inferred from spectroscopically determined $T_{\rm eff}$ and surface gravity ($\log g$) by comparison with theoretical evolutionary tracks. The position of analysed PG1159 stars in the ``observational HR diagram'', i.e., the $T_{\rm eff}$--$g$ diagram, are displayed in Fig.\,\ref{fighrd}. The high-luminosity stars have low $\log g$ ($\approx$\,5.5) while the low-luminosity stars have a high surface gravity ($\approx$\,7.5) that is typical for white dwarf (WD) stars. The derived mean mass is 0.57\,M$_\odot$, a value that is practically identical to the mean mass of WDs \citep{be:07}. The PG1159 stars co-exist with hot central stars of planetary nebulae and the hottest hydrogen-rich (DA) white dwarfs in the same region of the HR diagram. About every other PG1159 star is surrounded by an old, extended planetary nebula. For a recent review with a detailed bibliography see \citet{werner:06}. What is the characteristic feature that discerns PG1159 stars from ``usual'' hot central stars and hot WDs? Spectroscopically, it is the lack of hydrogen Balmer lines, pointing at a H-deficient surface chemistry. The proof of H-deficiency, however, is not easy: The stars are very hot, H is strongly ionized and the lack of Balmer lines could simply be an ionisation effect. In addition, every Balmer line is blended by a Pickering line of ionized helium. Hence, only detailed modeling of the spectra can give reliable results on the photospheric composition. The high effective temperatures require non-LTE modeling of the atmospheres. Such models for H-deficient compositions have only become available in the early 1990s after new numerical techniques have been developed and computers became capable enough. \begin{figure*}[bth] \vspace{-3.1cm} \begin{center} \epsfxsize=1.0\textwidth \epsffile{kwerner02.eps} \end{center} \vspace{-6.9cm} \caption{\label{fig:hrd} Complete stellar evolution track with an initial mass of 2\,M$_\odot$ from the main sequence through the RGB phase, the HB to the AGB phase, and finally through the post-AGB phase that includes the central stars of planetary nebulae to the final WD stage. The solid line represents the evolution of a H-normal post-AGB star. The dashed line shows a born-again evolution of the same mass, triggered by a very late thermal pulse, however, shifted by approximately $\Delta \log T_{\rm eff} = -0.2$ and $\Delta \log~L/$L$_\odot = - 0.5$ for clarity. The double-loop structure of the path is the consequence of a hydrogen-ingestion flash. The ``$\star$'' symbol shows the position of PG1159$-$035 \citep[from][]{werner:06}. } \end{figure*} The first quantitative spectral analyses of optical spectra from PG1159 stars indeed confirmed their H-deficient nature \citep{werner:91}. It could be shown that the main atmospheric constituents are C, He, and O. The typical abundance pattern is C=0.50, He=0.35, O=0.15 (mass fractions). It was speculated that these stars exhibit intershell matter on their surface, however, the C and O abundances were much higher than predicted from stellar evolution models. It was further speculated that the H-deficiency is caused by a late He-shell flash, suffered by the star during post-AGB evolution, laying bare the intershell layers. The re-ignition of He-shell burning brings the star back onto the AGB, giving rise to the designation ``born-again'' AGB star \citep{iben:83a}. If this scenario is true, then the intershell abundances in the models have to be brought into agreement with observations. By introducing a more effective overshoot prescription for the He-shell flash convection during thermal pulses on the AGB, dredge-up of carbon and oxygen into the intershell can achieve this agreement \citep{herwig:99c}. Another strong support for the born-again scenario was the detection of neon lines in optical spectra of some PG1159 stars \citep{werner:94}. The abundance analysis revealed Ne=0.02, which is in good agreement with the Ne intershell abundance in the improved stellar models. If we do accept the hypothesis that PG1159 stars display former intershell matter on their surface, then we can in turn use these stars as a tool to investigate intershell abundances of other elements. Therefore these stars offer the unique possibility to directly see the outcome of nuclear reactions and mixing processes in the intershell of AGB stars. Usually the intershell is kept hidden below a thick H-rich stellar mantle and the only chance to obtain information about intershell processes is the occurrence of the third dredge-up. This indirect view onto intershell abundances makes the interpretation of the nuclear and mixing processes very difficult, because the abundances of the dredged-up elements may have been changed by additional burning and mixing processes in the H-envelope (e.g., hot-bottom burning). In addition, stars with an initial mass below 1.5~M$_\odot$ do not experience a third dredge-up at all. The central stars of planetary nebulae of spectral type [WC] are believed to be immediate progenitors of PG1159 stars, representing the evolutionary phase between the early post-AGB and PG1159 stages. This is based on spectral analyses of [WC] stars which yield very similar abundance results (see papers by Crowther, Todt, and Gr\"afener in these proceedings). \begin{figure*} \begin{center} \epsfxsize=1.0\textwidth \epsffile{kwerner03.eps} \end{center} \vspace{-0.5cm} \caption{\label{fig:psifn} Detail from FUSE spectra of two relatively cool PG1159 stars (see labels). Note the following features. The \ion{F}{vi}~1139.5~\AA\ line which is the first detection of fluorine at all in a hot post-AGB star; the \ion{P}{v} resonance doublet at 1118.0 and 1128.0~\AA, the first discovery of phosphorus in PG1159 stars; the \ion{N}{iv} multiplet at 1132~\AA. Also detected are lines from \ion{Si}{iv} and \ion{S}{vi}. The broader features stem from \ion{C}{iv} and \ion{O}{vi} \citep{reiff:07}.} \end{figure*} | It has been realized that PG1159 stars exhibit intershell matter on their surface, which has probably been laid bare by a late final thermal pulse. This provides the unique opportunity to study directly the result of nucleosynthesis and mixing processes in AGB stars. Abundance determinations in PG1159 stars are in agreement with intershell abundances predicted by AGB star models for many elements (He, C, N, O, Ne, F, Si, Ar). For other elements, however, disagreement is found (Fe, P, S) that points at possible weaknesses in the evolutionary models. Generally, the abundance patterns clearly support the idea that [WC] stars are direct progenitors of PG1159 stars. | 7 | 10 | 0710.4506 |
0710 | 0710.1786_arXiv.txt | {A large fraction of the stellar mass is found to be located in groups of the size of the Local Group. Evolutionary status of poor groups is not yet clear and many groups could still be at an early dynamical stage or even still forming, especially the groups containing spiral and irregular galaxies only.} {We carry out a photometric study of a poor group of late-type galaxies around IC 65, with the aim: (a) to search for new dwarf members and to measure their photometric characteristics; (b) to search for possible effects of mutual interactions on the morphology and star-formation characteristics of luminous and faint group members; (c) to evaluate the evolutionary status of this particular group.} {We make use of our $BRI$ CCD observations, DPOSS blue and red frames, and the 2MASS $JHK$ frames. In addition, we use the \ion{H}{i} imaging data, the far-infrared and radio data from the literature. Search for dwarf galaxies is made using the SExtractor software. Detailed surface photometry is performed with the MIDAS package. } {Four LSB galaxies were classified as probable dwarf members of the group and the $BRI$ physical and model parameters were derived for the first time for all true and probable group members. Newly found dIrr galaxies around the IC 65 contain a number of \ion{H}{ii} regions, which show a range of ages and propagating star-formation. Mildly disturbed gaseous and/or stellar morphology is found in several group members. } {Various structural, dynamical, and star-forming characteristics let us conclude that the IC 65 group is a typical poor assembly of late-type galaxies at an early stage of its dynamical evolution with some evidence of intragroup (tidal) interactions.} | The main contributions of this paper are as follows: \begin{enumerate} \item We have selected four LSB dwarf companion candidates of the IC 65 group of galaxies on deep DPOSS frames according to their surface brightnesses, colours and morphology. \item The $B, R$ and $I$ band surface photometry is presented for the first time for all certain and probable members of the group. The bright group members were studied in the NIR $J, H$ and $K$ bands, too. An image gallery and the deduced $SB$ and colour profiles are shown, permitting the detailed morphological analysis of the galaxies studied. Their relevant physical and model characteristics are determined. \item Dynamical masses and star-forming characteristics of the bright group members are estimated using the new optical photometry and the available NIR, FIR, \ion{H}{i} and radio data. The probable evolutionary status of the group is discussed. \end{enumerate} An analysis of the available photometric and kinematic data of individual galaxies with emphasis to study the possible mutual interactions between the group members leads to the following results: \noindent $\bullet$ The available \ion{H}{i} imaging data show that all bright members and at least one dwarf companion have a nearly normal gaseous fraction with \ion{H}{i} mass to blue luminosity ratios in the range of 0.3 -- 1.0, consistent with their morphological type. The outer \ion{H}{i} isophotes of the IC 65, and especially of the UGC 622 appear disturbed, in agreement with perturbation estimates.\\ $\bullet$ The optical morphology of the bright galaxies generally appears to be regular, with barely significant disturbances in isophotes of the outer stellar disk of the IC 65 and UGC 622.\\ $\bullet$ All bright group members (except PGC~138291, which we could not study in such detail) consist of many blue star-forming knots and plumes, especially UGC~608. A comparison of the surface photometry with stellar population models of Bruzual \& Charlot (\cite{bruzual03}) indicates that these blue knots must have formed recently. % The available data do not allow to establish whether they formed simultaneously, e.g. in star-bursts possibly triggered by interactions.\\ $\bullet$ Two dIrr galaxies around the IC 65 both contain a number of \ion{H}{ii} regions, which show a range of stellar ages and provide an evidence of propagating star-formation. One of these galaxies - A~0101+4744 is a confirmed member of the group; the second one - A~0100+4734 appears to be located in front of the group.\\ $\bullet$ The brightest galaxies in both subgroups can fuel their current star-forming rates of $\sim$ 1 - 2 ${\cal M}_\odot$ yr$^{-1}$ for about the next 3 - 7 Gyr. The IC~65 group of galaxies obviously belongs to the class of less evolved groups. It is composed of late type spiral and irregular galaxies arranged in two subgroups. No massive early type galaxies are present. No hot gas has been detected in it by the ROSAT survey. % Some morphological irregularities and signs of enhanced SF in its members could be indicative of recent/ongoing mutual interactions. Yet, the individual group members have retained much of their initial gas component. A few available velocities point to a short crossing time of only $\sim$ 0.1~$H_0^{-1}$. However, this hardly means that the group has already reached a stable (virialized) configuration. The evidence, discussed above, lets us conclude that the IC~65 group of galaxies is a dynamically young system at a still relatively early stage of its collapse. | 7 | 10 | 0710.1786 |
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0710 | 0710.3783_arXiv.txt | We report the first results of the GammeV experiment, a search for milli-eV mass particles with axion-like couplings to two photons. The search is performed using a ``light shining through a wall'' technique where incident photons oscillate into new weakly interacting particles that are able to pass through the wall and subsequently regenerate back into detectable photons. The oscillation baseline of the apparatus is variable, thus allowing probes of different values of particle mass. We find no excess of events above background and are able to constrain the two-photon couplings of possible new scalar (pseudoscalar) particles to be less than $3.1\times 10^{-7} \mbox{ GeV}^{-1}$ ($3.5\times 10^{-7} \mbox{ GeV}^{-1}$) in the limit of massless particles. | The key to this experiment is the short 5 ns, 160 mJ pulses of 532 nm light emitted with a repetition rate of 20 Hz by our light source, a frequency-doubled Continuum Surelite I-20 Nd:YAG laser. As described below, the small duty cycle enables a large reduction in the detector noise via coincidence counting. The laser light is vertically polarized and when needed, a halfwave plate is used to obtain horizontal polarization. The laser pulses are sent through a vacuum system (diagrammed in Fig.~\ref{F:apparatus}) designed around an insulating warm bore inserted into a 6 m Tevatron superconducting dipole magnet. The magnet produces a 5 T vertical field uniform across the aperture of the $48 \mbox{ mm}$ inner diameter warm bore. A ``wall'' consisting of a high-power laser mirror on the end of a long (7 m) hollow stainless steel ``plunger'' is inserted into the warm bore. The mirror may be placed at various positions within the magnet by sliding the plunger. The plunger mirror projects the reflected wave into a photon state and the transmitted wave into a (pseudo-)scalar state, provided that the scalar is sufficiently weakly-interacting to pass through the material of the mirror. The mirror also functions to reflect the incident laser power out of the magnet to prevent heating of the magnet coils. The mirror is mounted on a welded stainless steel cap on the end of the plunger in order to prevent stray photons from passing through. Thus, the beam passing through the end of the plunger is a pure scalar beam. These scalars can then oscillate back into photons through the remaining magnetic field region inside the $35 \mbox{ mm}$ inner diameter hollow plunger. Upon exiting the magnetic field region, the interaction ceases and the photon-scalar wavefunction is frozen. This combined wavefunction then propagates $\sim 6\mbox{ m}$ into a dark box where a Hamamatsu H7422P-40 photomultiplier tube (PMT) module is used to detect single photons in coincidence with the laser pulses. As described below, a high signal to noise ratio is achieved due to the very short pulses emitted by the laser. The photon-scalar transition probability may be written in convenient units as: \begin{eqnarray} \label{E:conversionprob1} P_{\gamma \rightarrow \phi} &=& \frac{4 B^2\omega^2}{M^2 (\Delta m^2)^2} \sin^2 \left( \frac{\Delta m^2 L}{4\omega} \right) \\ \label{E:conversionprob2} &\approx& \frac{4 B^2\omega^2}{M^2 m_\phi^4} \sin^2 \left( \frac{m_\phi^2 L}{4\omega} \right) \\ \label{E:conversionprob3} &=& 1.5\times 10^{-11} \frac{(B/\mbox{Tesla})^2(\omega/\mbox{eV})^2}{(M/10^5 \mbox{ GeV})^2 (m_\phi/10^{-3} \mbox{ eV})^4} \nonumber \\ & & \times \sin^2 \left( 1.267 \frac{(m_\phi/10^{-3} \mbox{ eV})^2 (L/\mbox{m})}{(\omega/\mbox{eV})} \right) \end{eqnarray} where $B$ is the strength of the external magnetic field, $\omega$ is the inital photon energy, $L$ is the magnetic oscillation baseline. The mass-squared difference between the scalar mass and the effective photon mass, $\Delta m^2=m_\phi^2-m_\gamma^2$, characterizes the mismatch of the phase velocities of the photon wave and the massive scalar wave and determines the characteristic oscillation length. While the photon does not really gain a mass in a normal dielectric medium, the phase advance may be modelled with an effective imaginary mass $m_\gamma^2 = -2 \omega^2 (n-1)$, where $n$ is the index of refraction \cite{vanBibber:1988ge}. Both the warm bore and the interior of the plunger are pumped to moderate vacuum pressures of less than $10^{-4} \mbox{ Torr}$, and a conservative estimate gives $\sqrt{-m_\gamma^2} < 10^{-4} \mbox{ eV}$. Therefore, starting with Eqn.~\ref{E:conversionprob2} we assume that the contribution from the effective photon mass is negligible for larger values of $m_\phi$ near the PVLAS region. As can be seen from Eqn.~\ref{E:conversionprob3}, the meter scale baseline provided by typical accelerator magnets is well-suited for probing the milli-eV range of possible particle masses. This fact can be a curse as well as a boon because for a monochromatic laser beam, a fixed magnet length may accidentally coincide with a minimum in the oscillation rather than a maximum. Indeed, this is a possible reason why the BFRT experiment \cite{Ruoso:1992nx} did not see the PVLAS signal even though they had similar sensitivity. GammeV's plunger design allows us to change the oscillation baseline and thus scan through all possible values of the scalar mass in the milli-eV range without any regions of diminished sensitivity. The total conversion and regeneration probability contains two factors of Eqn.~\ref{E:conversionprob2}, corresponding to the pre-mirror and post-mirror magnetic field regions of lengths $L_1$ and $L_2$. The total probability varies as $\sin^2 \left(\frac {m_\phi^2 L_1}{4\omega} \right) \sin^2 \left(\frac{m_\phi^2 L_2}{4\omega} \right)$ where $L_1+L_2=6\mbox{ m}$. \begin{figure}[t] \includegraphics[width=0.48\textwidth]{gammevcartoon.eps} \caption{\label{F:apparatus} Diagram (not to scale) of the experimental apparatus. The initial vacuum chamber consists of a 10 m insulating warm bore which is offset by 1.6 m from the end of the 6 m magnetic field region, and is sealed to the sliding plunger via a double o-ring assembly. The sliding plunger has a range of motion of 1.9 m, and contains an independent vacuum chamber. The vacuum window at the far end slides within a stationary, long dark box.} \end{figure} To detect regenerated photons we use a $51\mbox{ mm}$ diameter lens to focus the beam onto the $5 \mbox{ mm}$ diameter GaAsP photocathode of the PMT. The alignment is performed using a low power green helium-neon alignment laser and a mock target. The alignment is verified both before and after each data-taking period by replacing the sealed plunger with an open-ended plunger, re-establishing the vacuum, and firing the Nd:YAG laser onto a flash paper target. An optical transport efficiency of $92\%$ is measured using the ratio of laser power transmitted through the open-ended plunger and through the various optics and vacuum windows, to the initial laser power, using the same power meter in both cases to remove systematic effects. The quantum efficiency of the photocathode is factory-measured to be $38.7\%$ while the collection efficiency of the metal package PMT is estimated to be $70\%$. The PMT pulses are amplified by 46 dB and then sent into a NIM discriminator. Using a highly attenuated LED flasher as a single photon source, the discriminator threshold is optimized to give 99.4$\%$ efficiency for triggering on single photo-electron pulses while also efficiently rejecting the lower amplitude noise. By studying the trigger time distribution, the deadtime fraction due to possible multiple rapid PMT pulses is found to be negligible (0.001$\%$). Thus, we estimate the total photon transport and counting efficiency to be $(25\pm 3)\%$. Using this threshold, and the built-in cooler to cool the photocathode to $0^\circ$C, we measure a typical dark count rate of 130 Hz. \begin{figure}[t] \includegraphics[width=0.48\textwidth]{ttime_leaky_100ns.eps} \caption{\label{F:timing} PMT trigger times for the four run configurations, shown relative to the expected time distribution of photons as calibrated from the leaky mirror data.} \end{figure} To perform the coincidence counting we use two Quarknet boards \cite{quarknet} \cite{quarknetgps} with 1.25 ns timing precision, referenced to a GPS clock. The Quarknet boards determine the absolute time of the leading edge of time-over-threshold triggers from the PMT and from a monitoring photodiode that is located inside the laser box. The clocks on the laser board and on the PMT board are synchronized using an external trigger from a signal generator. The absolute timing between the laser pulses and the PMT traces is established by removing the plunger with the mirror, and allowing the laser to shine on the PMT through several attenuation stages consisting of two partially reflective (``leaky'') mirrors, a pinhole, and multiple absorptive filters mounted directly on the aperture of the PMT module. The $10^{19}$ photons per second emitted by the laser are thus attenuated to a corresponding PMT trigger rate of less than 0.1 Hz for this timing calibration and to provide an {\it in situ} test of the data acquisition system. The regenerated photons should arrive at the same time as the straight-through photons since milli-eV particles are also highly relativistic. For coincidence counting between the laser pulses and the PMT, a 10 ns wide window is chosen and includes $99\%$ of the measured photon time distribution shown in Fig.~\ref{F:timing}. The coincident dark count rate can be estimated to be $R_{\rm{noise}}=20\mbox{ Hz}\times 130\mbox{ Hz}\times 10\mbox{ ns} = 2.6\times 10^{-5} \mbox{ Hz}$. This noise rate is negligible to the expected signal rate of $\sim 2\times 10^{-3} \mbox{ Hz}$ estimated from the central values of the PVLAS parameters. \begin{table} \begin{tabular}{|l|c|c|c|c|} \hline Configuration&$\#$ photons&Est.Bkgd&Candidates&g[GeV$^{-1}$]\\ \hline Horiz.,center&$6.3\times 10^{23}$&$1.6$&1&$3.4\times 10^{-7}$\\ Horiz.,edge&$6.4\times 10^{23}$&$1.7$&0&$4.0\times 10^{-7}$\\ Vert.,center&$6.6\times 10^{23}$&$1.6$&1&$3.3\times 10^{-7}$\\ Vert.,edge&$7.1\times 10^{23}$&$1.5$&2&$4.8\times 10^{-7}$\\ \hline \end{tabular} \caption{\label{F:table} Summary of data in each of the 4 configurations.} \end{table} | 7 | 10 | 0710.3783 |
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0710 | 0710.4450.txt | In several studies of Low Luminosity Active Galactic Nuclei (LLAGNs), we have characterized the properties of the stellar populations in LINERs and LINER/HII Transition Objects (TOs). We have found a numerous class of galactic nuclei which stand out because of their conspicuous 0.1--1 Gyr populations. These nuclei were called ''Young-TOs'' since they all have TO-like emission line ratios. To advance our knowledge of the nature of the central source in LLAGNs and its relation with stellar clusters, we are carrying out several imaging projects with the Hubble Space Telescope (HST) at near-UV, optical and near-IR wavelengths. In this paper, we present the first results obtained with observations of the central regions of 57 LLAGNs imaged with the WFPC2 through any of the V (F555W, F547M, F614W) and I (F791W, F814W) filters that are available in the HST archive. The sample contains 34$\%$ of the LINERs and 36$\%$ of the TOs in the Palomar sample. The mean spatial resolution of these images is 10 pc. With these data we have built an atlas that includes structural maps for all the galaxies, useful to identify compact nuclear sources and, additionally, to characterize the circumnuclear environment of LLAGNs, determining the frequency of dust and its morphology. The main results obtained are: 1) We have not found any correlation between the presence of nuclear compact sources and emission-line type. Thus, nucleated LINERs are as frequent as nucleated TOs. 2) The nuclei of "Young-TOs" are brighter than the nuclei of "Old-TOs" and LINERs. These results confirm our previous results that Young-TOs are separated from other LLAGNs classes in terms of their central stellar population properties and brightness. 3) Circumnuclear dust is detected in 88$\%$ of the LLAGNs, being almost ubiquitous in TOs. 4) The dust morphology is complex and varied, from nuclear spiral lanes to chaotic filaments and nuclear disk-like structures. Chaotic filaments are as frequent as dust spirals; but nuclear disks are mainly seen in LINERs. These results suggest an evolutionary sequence of the dust in LLAGNs, LINERs being the more evolved systems and Young-TOs the youngest. | \label{sec:Introduction} %{\bf\ojo\ojon ROSA: YOUR FILE CAME WITH FUNNY LINE BREAKS, SO CHECK %THAT I HAVE NOT UNDONE SOME OF THE PARAGRAPH BREAKS INADVERTEDLY...} Low-luminosity active galactic nuclei (LLAGNs) comprise 30\%\ of all bright galaxies (B$\leq$12.5) and are the most common type of AGN (Ho, Filippenko \& Sargent 1997a, hereafter HFS97). These include LINERs, and transition-type objects (TOs, also called weak-[OI] LINERs). These two types of LLAGNs have similar emission line ratios in [OIII]/H$\beta$, [NII]/H$\alpha$, and [SII]/H$\alpha$, but [OI]/H$\alpha$ is lower in TOs than in LINERs. LLAGNs constitute a rather mixed class and different mechanisms have been proposed to explain the origin of the nuclear activity, including shocks, and photoionization by a non-stellar source, by hot stars or by intermediate age stars (e.g. Ferland \& Netzer 1983; Filippenko \& Terlevich 1992; Binnette et al. 1994; Taniguchi, Shioya \& Murayama 2000). Because we do not know yet what powers them and how they are related to the Seyfert phenomenon, LLAGNs have been at the forefront of AGN research since they were first systematically studied by Heckman (1980). Are they all truly ``dwarf'' Seyfert nuclei powered by accretion onto nearly dormant supermassive black holes (BH), or can some of them be explained at least partly in terms of stellar processes? If LLAGNs were powered by a BH, they would represent the low end of the AGN luminosity function in the local universe and would also establish a lower limit to the fraction of galaxies containing massive BHs in their centers. If, on the contrary, LLAGNs were powered by nuclear stellar clusters, their presence would play an important role in the evolution of galaxy nuclei. Therefore, it is fundamental to unveil the nature of the central source in LLAGNs. There is clear evidence that at least some LINERs harbor a bona fide AGN, and they may be considered the faint end of the luminosity function of Seyfert galaxies. It has been found that about 20\%\ of the nearby LINERs have a weak broad H$\alpha$ emission component similar to those found in type 1 Seyferts (Ho, Filippenko \& Sargent 1997b). In a few of these LINERs the H$\alpha$ line shows a double peak component (e.g. Storchi-Bergmann et al. 1997; Shields et al. 2000; Ho et al. 2000) suggesting that they are powered by an accreting black hole (BH). It has also been found that X-ray emission in LINERs has a nonthermal origin associated with an AGN (e.g. Terashima, Ho \& Ptak 2000). Recent works based on higher spatial resolution data taken with Chandra show that only in half of the LLAGNs observed the X-ray emission is associated with compact nuclear cores (Satyapal et al. 2004; Dudik et al. 2006; Gonz\'alez-Mart\'\i n et al. 2006). This is consistent with VLA radio observations, which detect unresolved radio cores in half of the LINERs (Nagar et al. 2000, 2002). Finally, a monitoring study at near-UV wavelengths by Maoz et al. (2005) finds UV variability in a significant fraction of the 17 LLAGNs observed. The variation of the UV fluxes may be interpreted as the manifestation of low rate or low radiative efficiency accretion onto a supermassive BH. On the other hand, detection of stellar wind absorption lines in the ultraviolet spectra of some TOs (Maoz et al. 1998; Colina et al. 2002) has proven unequivocally the presence of young stellar clusters in the nuclear region. Additional evidence comes from optical studies, in which we have focused on the study of the stellar population in the nuclei and circumnuclear region of LLAGNs to establish the role of stellar processes in their phenomenology. Ground-based (Cid Fernandes et al.\ 2004; Cid Fernandes et al.\ 2005) and HST+STIS spectra (Gonz\'alez Delgado et al.\ 2004) have shown that the contribution of an intermediate age stellar population is significant in a sizable fraction of the TO population. These studies identified a class of objects, called ``Young-TOs'', which are clearly separated from LINERs in terms of the properties and spatial distribution of the stellar populations. They have stronger stellar population gradients, a luminous intermediate age stellar population concentrated toward the nucleus ($\sim$100~pc) and much larger amounts of extinction than LINERs. These objects, which underwent a powerful star formation event $\sim$ 1 Gyr ago, could correspond to post-Starburst nuclei or to evolved counterparts of the Seyfert 2 with a composite nucleus, characterized too by harboring nuclear starbursts (Gonz\'alez Delgado et al. 2001; Cid Fernandes et al. 2001, 2004). HST imaging of the nuclei of these Seyfert 2 galaxies shows that the UV emission is resolved into stellar clusters (Gonz\'alez Delgado et al. 1998) that are similar to those detected in starburst galaxies (Meurer et al.\ 1995). Nuclear stellar clusters are a common phenomenon in spirals, having been detected in 50-70\% of these sources (Carollo et al. 1998, 2002; Boeker et al. 2002, 2004). Therefore stellar clusters are a natural consequence of the star formation processes in the central region of spirals. On the other hand, evidence has been accumulating during the past few years about the ubiquity of BH in the nuclei of galaxies. Furthermore, the tight correlation of the BH mass and stellar velocity dispersion (Ferrarese \& Merrit 2000; Gebhardt et al.\ 2000) implies that the creation and evolution of a BH is intimately connected to that of the galaxy bulge. Recently, in a HST survey in the Virgo Cluster, C\^ot\'e et al. (2006) have detected compact sources in a comparable fraction of elliptical galaxies. These compact stellar clusters, referred to as nuclei by the authors (see also Ferrarese et al 2006a), have masses that scale directly to the galaxy mass, in the same way as do the BH masses in high luminosity galaxies (Ferrarese et al. 2006b). Therefore, a natural consequence of the physical processes that formed present-day galaxies should be the creation of a compact massive object in the nucleus, either a BH and/or a massive stellar cluster. To determine the nature of the nuclear source of active galaxies and their evolution we are carrying out several projects with HST+ACS imaging at the near-UV wavelengths a sample of Seyferts (ID. 9379, PI. Schmitt, Mu\~noz-Mar\'\i n et al. 2007) and LLAGNs (ID. 10548, PI. Gonz\'alez Delgado). These observations are complemented with WFPC2 optical data retrieved from the HST archive. The high angular resolution provided by HST is crucial to determine the physical properties of the nuclei, the central structure of these galaxies, as well as to study the circumnuclear environment of AGNs. The main goals of these studies are to determine the frequency of nuclear and circumnuclear stellar clusters in AGNs, and whether they are more common in Seyferts, TOs or LINERs; to characterize the intrinsic properties of these clusters and to study whether there is evolution from Seyferts to TOs and LINERs. In addition, the frequency of dust and its morphology can also provide relevant information about the origin of nuclear activity. Dust is a valuable probe of the presence of cold interstellar gas in galaxies, and it is very sensitive to the perturbations that drive the gas toward the center and feed the AGN. Here, we present the initial results obtained for LLAGNs based on archival visible and red images obtained with the WFPC2. The paper is organized as follows: section 2 presents the sample selection, and 3 the characteristics of the observations. Sections 4, 5 and 6 describe the imaging atlas, the dust morphology and the central properties of the galaxies. Finally, the summary and conclusions are presented in section 7. %%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%%SEC%%% | LLAGNs, that include LINERs and TOs, are the most common type of AGN. What powers them is still at the forefront of AGN research. To unveil the nature of the central source we are constructing a panchromatic atlas of the inner regions of these galaxies, which will be used to determine their nuclear stellar population. To this end we have already carried out a near-UV snapshot survey of nearby LLAGNs with the ACS on board HST, that is complemented with optical and near-IR images available in the HST archive. In this paper, the first of a series, we present observations of 57 LLAGNs imaged with the WFPC2 through any of the V (F555W, F547M, F606W) and/or I (F791W, F814W) bands. These objects comprise 36$\%$ of the original HFS97 LLAGNs sample, and correspond to those for which there are WFPC2 images available in the HST archive and whose circumnuclear stellar population we have already studied spectroscopically (Papers I--III). The subset of objects studied here follows the same distance and morphological type distribution of the complete HFS97 LLAGN sample. Classifying the objects in strong-[OI] and weak-[OI] ([OI]/H$\alpha$)$\leq$ 0.25), this subset includes 34$\%$ and 36$\%$ of the strong- and weak-[OI], respectively, of the whole HFS97 LLAGN sample. Following our results obtained from the analysis of the circumnuclear stellar population, this sub-sample contains 17 Young-TOs, 20 Old-TOs, 18 Old-LINERs and 2 Young-LINERs. Young-TOs or Young-LINERs are LLAGNs in which intermediate age stars contribute significantly to the nuclear blue-optical continuum. The dearth of Young-LINERs in the sample can be understood from the result obtained in our spectroscopic studies, which show that the overwhelming majority of LINERs harbor an old stellar population. With these data we have built an atlas that includes the structural map for all the images, and color maps for those galaxies for which images in two filters are available. We have identified those galaxies that have nuclear compact sources, and we have studied the circumnuclear environment of LLAGNs. We have found circumnuclear dust in 88$\%$ of the LLAGNs, but this fraction is somewhat larger (95$\%$) in weak-[OI] LLAGNs. The dust morphology is quite complex, from nuclear spiral lanes, to chaotic filaments and to nuclear disk-like structures. Chaotic filaments are as frequent as the well organized dust spirals; but disks are mainly seen in strong-[OI] LLAGNs. The dust concentration (simply graded by its location relative to a radius of 100-200 pc, is similar in weak- and in strong-[OI] because the fraction of LLAGNs with dust located only in the inner part is larger in Old-LLAGNs than in Young-LLAGNs. These results suggest an evolutionary dust sequence from Young-TOs to Old-LLAGNs. We have found that LINERs and TOs have both similar central magnitude and surface brightness, but LLAGNs with young and intermediate age populations are brighter than Old-TOs and LINERs. We have not found any correlation between the presence of nuclear compact sources and the emission line spectral type, ie., LINERs are as frequently nucleated as TOs. However, the centers of Young-TOs are brighter than the centers of Old-TOs and LINERs. The difference in magnitude and surface brightness can be even larger if we account for internal extinction, since Young-TOs are dustier. This result indicates that Young-TOs are separated from other type of LLAGNs also in terms of their central brightness, in addition of the properties and spatial distribution of the stellar population. These data have been very useful to study the circumnuclear environment of LLAGNs, and to identify which of these galaxies have a nuclear compact source. The fact that compact sources are as frequent in LINERs as in TOs, confirms again that LLAGNs are a mixed bag of objects. These results also suggest that the central morphology alone is not sufficient to elucidate the origin of their central source, and it cannot be used to ascertain whether LLAGNs are powered by AGNs or stellar clusters. These data will be complemented with near-UV (ACS) and near-IR (NICMOS) images to provide a panchromatic atlas of the inner regions of LLAGNs and to further investigate the origin of the nuclear sources and their relation with stellar clusters. {\bf Acknowledgements} We thanks the referee for her/his suggestions that helped to improve the paper. RGD and EP acknowledge support from the Spanish Ministerio de Educaci\'on y Ciencia through the grant AYA2004-02703. The data used in this work come from observations made with NASA/ESA Hubble Space Telescope, obtained from the STScI data archive. Basic research in radio astronomy at the NRL is supported by 6.1 base funding. We also thank support from a joint CNPq-CSIC bilateral collaboration grant. %%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%%REF%%% | 7 | 10 | 0710.4450 |
0710 | 0710.4992_arXiv.txt | We establish a nonminimal Einstein-Yang-Mills-Higgs model, which contains six coupling parameters. First three parameters relate to the nonminimal coupling of non-Abelian gauge field and gravity field, two parameters describe the so-called derivative nonminimal coupling of scalar multiplet with gravity field, and the sixth parameter introduces the standard coupling of scalar field with Ricci scalar. The formulated six-parameter nonminimal Einstein-Yang-Mills-Higgs model is applied to cosmology. We show that there exists a unique exact cosmological solution of the de Sitter type for a special choice of the coupling parameters. The nonminimally extended Yang-Mills and Higgs equations are satisfied for arbitrary gauge and scalar fields, when the coupling parameters are specifically related to the curvature constant of the isotropic spacetime. Basing on this special exact solution we discuss the problem of a hidden anisotropy of the Yang-Mills field, and give an explicit example, when the nonminimal coupling effectively screens the anisotropy induced by the Yang-Mills field and thus restores the isotropy of the model. | The discussion of a nonminimal coupling (NMC) of gravity with fields and media has a long history. The most intensely this topic has been studied in connection with the problem of nonminimal coupling of gravity and scalar field, which has numerous cosmological applications. The details of the investigations of this problem are discussed, e.g., in the review of Faraoni {\it et al\/}.\cite{FaraR} The development of the theory of NMC of gravity and scalar field $\phi$ has started by the introduction of the term $\xi \phi^2 R$ to the Lagrangian ($R$ is the Ricci scalar). In Ref.~\refcite{Chernikov} the special choice $\xi = 1/6$ has been motivated by the conformal invariance; in Ref.~\refcite{Callan} this quantity was considered as an arbitrary parameter of the model. Such a model has been widely used for the cosmological applications, in which $\xi$ played a role of extra parameter of inflation (see, e.g., Refs.~\refcite{Abbott}--\refcite{Fara4}). In Refs.~\refcite{HDehnen1}--\refcite{HDehnen4} the gauge-invariant term $\alpha {\bf \Phi}^{+}{\bf \Phi}R $ has been introduced instead of $\xi \phi^2 R$ in the context of non-Abelian gauge theory (${\bf \Phi}$~is a multiplet of scalar complex Higgs fields interacting with gravity and spinor matter.) Subsequent generalizations have been related to the replacement of $\xi \phi^2$ by the function $f({\Phi}^2)$ (see, e.g., Refs.~\refcite{Bergmann}--\refcite{Steinh}), as well as, to the inserting of the terms of the type $F({\Phi}^2, {\cal R})$ both linear and nonlinear in the Ricci scalar, Ricci and Riemann tensors (see, e.g., Refs.~\refcite{Linde}--\refcite{Inagaki}). The idea of nonminimal derivative coupling introduced in Ref.~\refcite{Amen3} and developed further in Refs.~\refcite{Capo1,Capo2} has enriched the NMC modeling by the terms $\phi_{,ij..}$. Nonminimal cosmological models based on the formalism of derivative coupling are the multi-parameter ones and have supplementary abilities for a fitting of observational data. Let us note that the NMC of gravity and scalar field leads to the modifications of both the Klein-Gordon and the Einstein equations, and such modifications are of interest for various inflation scenarios. Thus, the modeling of nonminimal interactions of scalar and gravitational fields is one of the well established and physically motivated branch of modern cosmology. Natural extension of the nonminimal theory from the models with scalar fields coupled to curvature to the models describing scalar fields interacting with gauge fields has the same sound motivation and can disclose new aspects of cosmological dynamics. The study of the nonminimal coupling of gravity with electromagnetic field has another motivation and another history. In 1971 Prasanna\cite{Prasa1} introduced the invariant $R^{ikmn}F_{ik}F_{mn}$ ($R^{ikmn}$ is the Riemann tensor, $F_{ik}$ is the Maxwell tensor) as a possible element of a Lagrangian, and then in Ref.~\refcite{Prasa2} obtained the corresponding nonminimal one-parameter modification of the Einstein-Maxwell equations. In 1979 Novello and Salim\cite{Novello1} proposed to insert the gauge non-invariant terms $R A^k A_k$ and $R^{ik}A_i A_k$ in the Lagrangian ($A_k$ is an electromagnetic potential four-vector). A qualitatively new step has been made by Drummond and Hathrell in Ref.~\refcite{Drum}, where the one-loop corrections to the quantum electrodynamics (QED) are obtained, which take into account the nonminimal coupling of gravity and electromagnetism. The Lagrangian of such a theory happens to contain three fundamental $U(1)$-gauge-invariant scalars $R^{ikmn}F_{ik}F_{mn}$, $R^{ik}g^{mn}F_{im}F_{kn}$ and $RF_{mn}F^{mn}$ with coefficients reciprocal to the square of the electron mass. This Lagrangian had no arbitrary parameters, but curvature induced modifications of the electrodynamic equations gave the impetus to wide discussions about the formal structure of the nonminimal Lagrangian, basic evolutionary equations, breaking the conformal invariance and the properties of the photons, coupled to curvature in different gravitational backgrounds (see, e.g., Refs.~\refcite{Acci1}--\refcite{Lafrance}). The last paper revived, as well, the interest to the paradigm: curvature coupling and equivalence principle, various aspects of which are now discussed (see, e.g., Refs.~\refcite{Prasa3,Solanki}). The QED-motivation of the use of the generalized Maxwell equations can also be found in the papers of Kosteleck\'y and colleagues.\cite{Kost1,Kost2} The effect of birefringence induced by curvature, first discussed in Ref.~\refcite{Drum}, and some of its consequences for the electrodynamic systems have been investigated in Refs.~\refcite{Balakin1}--\refcite{Balakin5} for the case of pp-wave background. The generalization of the idea of nonminimal interactions to the case of torsion coupled to the electromagnetic field has been made in Refs.~\refcite{Hehl1,Hehl2} (see, also, Ref.~\refcite{Hehl3} for a review on the problem). To summarize we stress that the study of electrodynamic systems nonminimally coupled to the gravity field poses a natural question about curvature induced variations of photon velocity in the cosmological background. Since the interpretation of observational data in cosmology depends essentially on the velocity of photon propagation during different cosmological epochs, the modeling of nonminimal electrodynamic phenomena seems to be well motivated and interesting from physical point of view. Concerning the nonminimal Einstein-Yang-Mills (EYM) theory, we can distinguish between two different ways to establish it. The first way is the direct nonminimal generalization of the Einstein-Yang-Mills (EYM) theory. In the framework of this approach Horndeski\cite{Horn} and M\"uller-Hoissen\cite{MH} obtained the nonminimal one-parameter EYM model from a dimensional reduction of the Gauss-Bonnet action. Now the Gauss-Bonnet models are of great interest in connection with the problem of dark energy (see, e.g., the Gauss-Bonnet model with nonminimal scalar field\cite{OdinDE}). Thus, the non-Abelian multi-parameter extensions of nonminimal models are also well motivated, since they give a chance to explain the accelerated expansion of the Universe without addressing to exotic substance. We follow the alternative way, which is connected with a non-Abelian generalization of the nonminimal Einstein-Maxwell theory along the lines proposed by Drummond and Hathrell\cite{Drum} for the linear electrodynamics. Based on the results of Ref.~\refcite{BL05} a three-parameter gauge-invariant nonminimal EYM model linear in curvature is considered.\cite{1BZ06}\cdash\cite{BDZ07} Our goal is to formulate a nonminimal Einstein-Yang-Mills-Higgs (EYMH) theory, and this process, of course, also admits different approaches. In fact, the nonminimal EYMH theory should accumulate the ideas and methods both from the nonminimally extended EYM theory and from the nonminimally extended scalar field theory. Initial attempt to develop nonminimal EYMH theory can be found, for instance, in Ref.~\refcite{Bij}, where the scalar Higgs field is nonminimally coupled with gravity via $\xi {\Phi}^2 R$ term, and the Higgs field ${\bf \Phi}$ is included into the Lagrangian of the Yang-Mills field in a composition with a square of the Yang-Mills potential: ${\Phi}^2 A_k^{(a)} A^k_{(a)}$. Such a theory is not gauge-invariant. In this paper we establish a new six-parameter nonminimal Einstein-Yang-Mills-Higgs model. First three coupling parameters, $q_1$, $q_2$ and $q_3$, describe a nonminimal interaction of Yang-Mills field and gravitational field. The fourth and fifth parameters, $q_4$ and $q_5$, describe the so-called gauge-invariant nonminimal derivative coupling of the Higgs field with gravity. Since the gauge-invariant derivative, $\D_m {\Phi}^{(a)}$, contains the potential of the Yang-Mills field, the corresponding nonminimal term is associated with ``triple'' interaction, namely, gravitational and scalar fields, gauge and scalar fields, and gauge and gravitational fields. The sixth parameter, $\xi$, is the well-known coupling parameter nonminimally connecting gravitational and scalar fields via the term $\xi R {\Phi}^2$. Of course, this model is only one of a wide class of the nonminimal EYMH models. As for its motivation and possible physical applications, one can see that on the one hand, the interest to a six-parameter nonminimal EYMH model is based on the sound results obtained earlier in the framework of partial nonminimal models (Einstein-Maxwell, Einstein-Yang-Mills and scalar field theories), on the other hand, the six-parameter model under discussion shows new specific solutions of cosmological type, which can not appear in more simple models. The paper is organized as follows. In Sec. \ref{Formalism} we formulate the nonminimal EYMH model, which contains six phenomenological coupling parameters, and establish the nonminimally extended Yang-Mills, Higgs and Einstein equations. In Sec. \ref{IsModel} we apply the introduced master equations to the spacetime with constant curvature and obtain the specific relationships between coupling constants, which turn the extended equations for the gauge field and scalar field into identities. In Subsec. \ref{dSsptime} we discuss the exact solutions to the nonminimal EYMH equations attributed to the isotropic cosmological model with Yang-Mills field, characterized by hidden anisotropy. | \label{Discussion} \noindent \par 1. The main mathematical result of the presented paper is the establishing of a new self-consistent nonminimal system of master equations for the coupled Yang-Mills, Higgs and gravity fields from the gauge-invariant nonminimal Lagrangian (\ref{1act}). The obtained mathematical model contains six arbitrary parameters, and, thus, admits a wide choice of special sub-models interesting for the applications to the nonminimal cosmology (isotropic and anisotropic) and nonminimal colored spherical symmetric objects. The applications require the phenomenological coupling constants $q_1$, $q_2,\ \dots,\ q_5$ and $\xi$ to be interpreted adequately. Following the idea, discussed in Ref.~\refcite{HDehnen4}, we intend not to introduce ``new constants of Nature'', but to relate the phenomenological parameters with the constants well-known in the High Energy Particle Physics, on the one hand, and with the constants of cosmological origin, on the other hand. Indeed, in the specific cosmological model, established above, the sixth phenomenological parameter $\xi$ is expressed in terms of the square of the effective mass of the Higgs bosons $\mu$ and constant curvature $K$, $\xi = \frac{\mu}{12K}$. Other parameters are expressed in terms of $K$ (see (\ref{q})). Since in the de Sitter model the Hubble constant is $H=\sqrt{K}$, one can say that $q_1$, $q_2,\ \dots,\ q_5$ are connected with $H$. Analogously, one can consider the equality $H^2=K=\frac{\Lambda}{3}$ and thus, one can say that they are connected with the cosmological constant $\Lambda$. In any case the parameters of nonminimal coupling $q_1$, $q_2,\ \dots,\ q_5$ can be expressed in terms of cosmological parameters $K$, $H$ or $\Lambda$, and define a specific radius of curvature coupling, $r_q \equiv \frac{1}{\sqrt{K}}$ and the corresponding time parameter $t_q \equiv r_q/c$. 2. The curvature coupling modifies the master equations for the Yang-Mills and Higgs fields. According to (\ref{Heqs}) a new tensor ${H}^{ik}_{(a)}$ appears (see (\ref{HikR})), which is an analog of the induction tensor in the Maxwell theory\cite{Maugin}. This means that the curvature coupling of the non-Abelian gauge field with gravity acts as a sort of quasi-medium with a nonminimal susceptibility tensor ${\cal R}^{ikmn}$ (see (\ref{HikR})). As well, the curvature coupling modifies the master equations for the Higgs field, and the tensor $\Re^{mn}$, according to (\ref{Heq}), can be indicated as a simplest nonminimal susceptibility tensor for the Higgs field, and the vector ${\Psi}^{m}_{(a)}$ (see (\ref{21Heq})) can be defined as scalar induction. For the specific set of coupling constants (see (\ref{AQU}), (\ref{q})) the non-Abelian induction $H^{ik}_{(a)}$ and the scalar induction ${\Psi}^{m}_{(a)}$ can turn into zero, despite the fact that the Yang-Mills field strength ${F}^{ik}_{(a)}$ and the Higgs field ${\Phi}^{(a)}$ are non-vanishing. This means that, when (\ref{AQU}) holds, the possibility exists to satisfy the nonminimally extended Yang-Mills and Higgs equations for arbitrary ${F}^{ik}_{(a)}$ and ${\Phi}^{(a)}$. This possibility gives, in principle, a new option for modeling physical processes in Early Universe and shows very interesting analogy between this nonminimal model and resonance phenomena in plasma physics. Indeed, when we deal with plasma waves (for instance, with the longitudinal waves) one can see that electric induction $\vec{D}$ is connected with the longitudinal electric field $\vec{E}_{||}$ with the frequency $\omega$ by the relation $\vec{D}= \varepsilon_{||}\vec{E}_{||}$. Here $\varepsilon_{||}$ is the longitudinal dielectric permittivity, the simplest expression for this quantity can be obtained in the limit of long waves and gives $\varepsilon_{||}= 1-\frac{\Omega^2_{p}}{\omega^2}$, where $\Omega_{p}$ is the well-known plasma frequency. When $\omega=\Omega_{p}$, one obtains $\vec{D}=0$ and electrodynamic equations are satisfied for arbitrary $\vec{E}_{||}$. Analogous feature can be found in the nonminimal model described above (see Eq. (\ref{1simpa})). Indeed, the quantity $K$ with the dimensionality of squared frequency ($c=1$) can be regarded as an analog of $\Omega^2_{p}$, the quantity $2(6q_1+3q_2+q_3)$ can be indicated as $1/\omega^2$, then the term $1-2K(6q_1+3q_2+q_3)$ plays a role of effective permittivity scalar $\varepsilon_q$. When this effective permittivity scalar vanishes, i.e., when the constants of nonminimal coupling are connected with the constant curvature $K$ according to (\ref{AQU}), we obtain the resonance case, for which the Yang-Mills and Higgs equations are satisfied identically for arbitrary strength field tensor $F^{ik}_{(a)}$ and Higgs multiplet $\Phi^{(a)}$, the color induction $H^{ik}_{(a)}$ being equal to zero. 3. The vector potential of the Yang-Mills field $ A^{(a)}_i$ enters the master equations via the gauge covariant derivative $\hat{D}_k$, thus, the gauge field generates an anisotropy in the spacetime. Such an anisotropy, in general case, breaks down the symmetry of the model and produces the isotropy violation. Nevertheless, as it was shown above, the nonminimal coupling can effectively screen the anisotropy and guarantee the symmetry conservation. In the framework of this model one can speak about hidden anisotropy of the Yang-Mills field, keeping in mind that non-Abelian gauge field enters the master equations for the gravity field in the isotropic combinations only. | 7 | 10 | 0710.4992 |
0710 | 0710.3867_arXiv.txt | High speed collisions, although current in clusters of galaxies, have long been neglected, as they are believed to cause little damages to galaxies, except when they are repeated, a process called ``harassment". In fact, they are able to produce faint but extended gaseous tails. Such low-mass, starless, tidal debris may become detached and appear as free floating clouds in the very deep HI surveys that are currently being carried out. We show in this paper that these debris possess the same apparent properties as the so-called ``Dark Galaxies", objects originally detected in HI, with no optical counterpart, and presumably dark matter dominated. We present a numerical model of the prototype of such Dark Galaxies -- VirgoHI21 --, that is able to reproduce its main characteristics: the one-sided tail linking it to the spiral galaxy NGC 4254, the absence of stars, and above all the reversal of the velocity gradient along the tail originally attributed to rotation motions caused by a massive dark matter halo and which we find to be consistent with simple streaming motions plus projection effects. According to our numerical simulations, this tidal debris was expelled 750~Myr ago during a fly-by at 1100~km~s$^{-1}$ of NGC~4254 by a massive companion which should now lie at a projected distance of about 400~kpc. A candidate for the intruder is discussed. The existence of galaxies that have never been able to form stars had already been challenged based on theoretical and observational grounds. Tidal collisions, in particular those occurring at high speed, provide a much more simple explanation for the origin of such putative Dark Galaxies. | With the availability of unprecedented deep HI blind surveys, a population of apparently free floating HI clouds without any detected stellar counterpart has become apparent \citep{meyer04,davies04,things,alfalfa1,alfalfa2}. It has been suggested that a fraction of them could be ``dark galaxies", a putative family of objects that would consist of a baryonic disk rotating in a dark matter halo, but that would differ from normal galaxies by being free of stars, having all their baryons under the form of gas. They would thus be ``dark'' in the optical and most other wavelengths, but visible through their HI emission, contrary to pure ``dark matter'' halos. Such dark galaxies would be extreme cases of Low Surface Brightness Galaxies (LSBs), a class of objects that have a particularly faint stellar content compared to their gaseous and dynamical masses \citep[e.g.,][]{carignanfreeman88}. The formation of low mass dark galaxies is actually predicted by $\Lambda$-CDM models \citep[e.g.,][]{vandenbosch03,tully05}. \citet{taylor05} provided theoretical arguments against the existence of galaxies that would have remained indefinitely stable against star formation, unless they are of very low mass, at least a factor ten below than that of classical dwarf galaxies. If they exist, the dark galaxies are predicted to have a low dynamical mass and HI content. In the Local Group, some possibly rotating, high-velocity clouds were speculated to be dark galaxies \citep{simon04,simon06}. Further away, \citet{davies06} argued that most previous HI blind surveys were not sensitive enough to rule out the existence of dark galaxies. And indeed, while the HIPASS survey failed at detecting HI~clouds without optical counterparts \citep{doyle05}, deeper recent HI observations, in particular with the Arecibo telescope, have revealed a number of dark galaxy candidates \citep{alfalfa2}. Among these free-floating low-mass HI clouds, one object located in the outerskirts of the Virgo Cluster has attracted much attention and discussion: VirgoHI21 \citep[][see Fig.~1]{davies04,minchin05}. Despite a HI mass of only $\sim 10^8$~M$_\sun$, this elongated gaseous structure, mapped with the Westerbork Synthesis Radio Telescope (WSRT) by \citet[][hereafter M07]{minchin07}, exhibits a velocity gradient as large as 220~km~s$^{-1}$ (see Fig.~\ref{fig:pv-vhi21}). {\it Assuming} that the observed HI velocities trace rotation, the inferred dynamical mass would be as large as $\sim 10^{11}$~M$_\sun$. The object shows no optical counterpart, even on deep HST images (M07). With such extreme properties, VirgoHI21 has become the prototype for dark galaxies, although its high dynamical mass is atypical even in models predicting the existence of Dark Galaxies. If real, an object like VirgoHI21 could tidally disturb the galaxies in their neighborhood, as investigated by \cite{karachentsev06}. Actually VirgoHI21 itself lies at about 150~kpc from the massive spiral galaxy NGC~4254 (M~99), to which it is connected by a faint HI~filament. This structure could in principle be a bridge linking the two galaxies and would then appear as a sign of a tidal interaction between them (M07). \begin{figure*} \centering \includegraphics[width=\textwidth,angle=0]{f1.pdf} \caption{The system VirgoHI21/NGC 4254. {\it Left} The distribution of the atomic hydrogen is superimposed in blue on a true color optical SDSS image of the field acquired through the WIKISKY.org project. The HI maps actually combine two data set: the observations obtained with the Arecibo telescope as part of the ALFALFA project \citep[courtesy of B. Kent,][]{haynes07}, which are sensitive enough to show the whole extent of the gaseous tail; the observations obtained at WSRT \citep[courtesy of R. Minchin,][]{minchin07} which are not as deep, but have a much better spatial resolution. The HI maps were smoothed and the WRST data have been deconvolved by the elongated telescope beam. {\it Right} Color coded first moment, velocity, field of the HI tail as mapped by Arecibo. The observed velocity range in km~s$^{-1}$ is indicated to the right.} \label{fig:vhi21} \end{figure*} \begin{figure} \centering \includegraphics[width=\columnwidth,angle=0]{f2.pdf} \caption{Observed Position (x-axis) Velocity (y-axis) diagram along the position angle 22 degrees corresponding to the main direction of the HI bridge. As for Figure~\ref {fig:vhi21}, the WSRT (gray) and Arecibo (light blue) data set have been superimposed to show all available information.}\label{fig:pv-vhi21} \end{figure} However, starless isolated gas clouds, showing a large velocity spread, are not necessarily genuine dark galaxies. Ram pressure can strip gas away from spirals in the vicinity of clusters, a process that does not affect stars. Interaction with an external field, for instance that of another galaxy, can expulse large amounts of material from the disk in the form of gas-rich tidal tails and debris. In that vein, \citet[][hereafter B05]{bekki05} have suggested that interactions between flying-by (i.e. interacting without merging) galaxies orbiting in a potential well similar as the one produced by the Virgo Cluster form tidal tails that after some time can resemble isolated gas clouds containing little stars. Furthermore tidal tails are the place of large streaming motions, as shown for instance by the observations of the merger prototype NGC~7252 by \citet{hibbard94} and the models of \citet{bournaud04} and B05. The remaining HI structures of galaxy collisions can thus not only appear as isolated HI clouds, but also exhibit large velocity gradients that mimic those expected for rotating disks within dark matter haloes. In these conditions, serious doubts have raised whether VirgoHI21 is really a dark galaxy and not simply the result of a tidal interaction or of an harassment process in the Virgo Cluster (B05). They have recently been reinforced by the publication by \citet{haynes07} of a deep HI map of the field, obtained with the Arecibo Telescope as part of the ALFALFA survey \citep{alfalfa1,alfalfa2}. It reveals that VirgoHI21 actually lies within an even larger HI structure that extends further to the North, in the opposite direction of NGC~4254. This further suggests that the thin and long HI feature is in fact a tidal tail emanating from the spiral and that VirgoHI21 is just a denser cloud within it and thus not a dark galaxy. On the basis of HI spectra obtained with the Effelsberg telescope and a numerical model including the cluster ram pressure, \citet{vollmer05} suggested that NGC~4254 has indeed interacted recently with a companion. However this study mostly focussed on the internal properties of the spiral and in particular the formation of VirgoHI21 was not modeled. Nevertheless, several arguments against a tidal origin for the HI cloud have been raised and addressed in detail in M07. The main ones regard the absence of a suitable interacting companion, the nonexistence of a counter tail -- a feature that is generally present in tidal interactions --, the total lack of stars in the HI tail/bridge, and, above all, the remarkable 200~km~s$^{-1}$ velocity gradient associated to VirgoHI21, which seems to be reversed and amplified with respect to the large scale velocity field along the rest of the HI structure. We will show in this paper that most of these criticisms actually apply to low-velocity encounters and not to the high-velocity ones, which are common in the cluster environment. High-velocity collisions have been neglected so far because they seem to cause little disturbances to stellar disks unless they are very numerous and participate to an harassment process \citep[e.g.,][]{moore96}. To further investigate the role of high speed collisions in the formation of tidal debris, especially the gaseous ones, we have carried out a series of numerical simulations. We illustrate their impact showing a numerical model reproducing the morphology and kinematics of VirgoHI21. The numerical simulations are presented in Section~2. In Section~3, we compare the properties of tidal tails formed in high- and low-velocity galaxy encounters. In Section~4, we present the model which best fits VirgoHI21. Its actual nature is discussed in Section~5. Our conclusions and the implications for the detection of dark galaxies are summarized in Section~6. | \subsection{From tidal tails to fake dark galaxies} \label{fakedark} As shown in Section~\ref{sect:highvel}, tidal tails, especially those formed during high-velocity encounters, share many of the properties expected for dark galaxies: starless HI features, strong velocity gradients due in one case to streaming motions and in the other to the presence of a massive dark matter halo. If furthermore some gaseous condensations are present in the tail, they may resemble isolated dark galaxies: indeed the bridge to the parent galaxy can be very faint and hard or impossible to detect on moderately deep HI maps. Tidal tails do not necessarily have uniform profiles; denser regions can even lie in their outermost regions \citep[e.g.,][]{DBM04}. This can be because large pre-existing clouds from the parent disk are moved into the tail where they can form local overdensities \citep{elmegreen93}, or because some parts of an initially uniform tail condense under the effect of gravity \citep{BH92}. That these regions will lie far from the progenitor disk is a natural consequence of the extended flat rotation curves of spirals \citep{DBM04}. Depending on their location and mass, these denser parts of tails can even be self-gravitating, form stars and become finally independent objects with the mass of dwarf galaxies: the so-called Tidal Dwarf Galaxies \citep[TDGs,][for a review]{duc07}. In such cases, the internal motions within the young gravitationally bound object, in particular its rotation, will induce an additional velocity gradient. On the other hand, the less massive condensations that have not reached the critical HI column density threshold to form stars will appear as detached HI clouds without any stellar counterpart. Many of the known free-floating HI clouds, which are considered as Dark Galaxy candidates, have been found in clusters. In this environment, high velocity encounters have a high probability to occur, due to the large velocity dispersion of the cluster galaxies. Thus intra-cluster low-mass HI clouds of tidal origin could then be common but more dedicated studies should check this. Of course, this mechanism applies only if the parent galaxy had before the collision an extended HI disk. This requires in particular that it has not already crossed the cluster core where ram pressure would have contributed to truncate its gaseous disk. Tidal material generally quickly falls back onto the parent spiral, but this can take more than 2~Gyr in the outer parts \citep[e.g.,][]{BD06}. The cluster tidal field can even prevent tidal debris from falling back \citep{Mihos04}. This leaves time for fake dark galaxies of tidal origin to be observed while the interloper galaxy can be far away: at 1000~km~s$^{-1}$, it can be at a projected distance up to 2~Mpc two billion years later. \bigskip \subsection{The nature of VirgoHI21} After having argued that dark galaxies may in general be mistaken with tidal features and thus be fake ones, we discuss more specifically the nature of VirgoHI21, presenting pro and con arguments for the different scenarios proposed so far for its origin. \subsubsection{Tidal debris?} M07 argued that the VirgoHI21 + HI~bridge system cannot be a tidal tail from the spiral NGC~4254, because of the following reasons: \begin{itemize} \item Interacting galaxies generally have pair of tidal tails, and indeed the models in B05 have two-tailed morphologies, while NGC~4254 has no counter-tail. \item The tidal HI clouds in B05 models are star-poor but not star-free, while HST optical observations show VirgoHI21 being completely dark. \item The velocity gradient along the HI~bridge gets reversed around the VirgoHI21 cloud, which seems to rotate in a direction opposite to the rest of the HI bridge. This reversal is said by M07 not to be explained by interaction models as those of B05. \item The velocity spread ($\Delta V=200$~km~s$^{-1}$) of VirgoHI21 could only be explained by an encounter at a comparable velocity (according to M07), implying that the interloper should not be further away than a few arcminutes and would have been identified. Moreover, low relative velocities are rare in the Virgo Cluster. \item High velocity encounters are more common in this environment but velocities as large as $\sim$~1000~km~s$^{-1}$ are ``far too large to generate tidal features such as bridges and tails" (M07) . \end{itemize} If indeed VirgoHI21 is a genuine dark galaxy, as claimed by M07, the HI~bridge would then be a tail expulsed from the dark galaxy by an interaction with NGC~4254 or the cluster field. We note here that our numerical model of VirgoHI21 which suggests that the tidal debris rather emanate from NGC~4254 addresses each of these above-mentioned concerns: \begin{itemize} \item The interaction occurred about 750~Myr ago, so that the counter-tail from the spiral is not necessarily expected to be observed anymore. By that time the interloper might also be far away. \item The HI cloud formed in our model during a high--velocity encounter is free of old stars pre-existing to the galaxy interaction, and is not dense enough to form new stars. \item The velocity spread ($\Delta V=200$~km~s$^{-1}$) of the HI structure does not imply that the galaxy interaction had a similar velocity. It is in fact accounted for by a high-velocity encounter. \item The reversal of the velocity gradient along the tail is reproduced thanks to projection effects. \end{itemize} Whereas the global kinematics of VirgoHI21 and its bridge can be simply explained by streaming motions along a tidal structure, in detail, the model and the observations show some differences. They may be noted when comparing the observed (Fig.~\ref{fig:pv-vhi21}) and simulated (Fig.~\ref{fig:posvel}) Position--Velocity diagrams along the tail; in particular, as put forward by M07, the velocity gradient towards VirgoHI21 is locally larger within the HI cloud, while in our model it is not more enhanced at this location than further away near the tip of the tail. This local difference is actually not a real concern for our scenario. First, part of the local amplification of the gradient can result from the self-gravity of the VirgoHI21 cloud itself, that could be somewhat denser than in our model. In particular, as recently shown by \citet{B07}, tidal debris may contain a significant dark component that has not been included in the simulations. Alternatively, the velocity field can be disturbed by objects in the neighborhood, for instance by the nearby dwarf galaxy, SDSS J121804.26+144510.4, also known as object~C (see Fig.~\ref{fig:vhi21} and M07). The nature of this dwarf is actually unclear (see Appendix). We suppose here that it is a pre-existing object, physically interacting with the system. Object~C has a visible HI mass of $2\times 10^7$~M$_\sun$. We included its possible influence in a simple model. We describe it as a dwarf spheroid with a Plummer profile, total mass of $4\times 10^8$~M$_\sun$ (to include the dark matter mass) and scale-length of 5~kpc. It is located 8~kpc East from VirgoHI21 on the sky plane, and we assume it lies 20~kpc below VirgoHI21 in the radial direction. Simulating the whole trajectory of this dwarf in the simulation from $t=0$ would introduce too many parameters. Since we just want to illustrate its possible local effect, we simply add it as a fixed mass in the late stage, linearly increasing its mass from $0$ at $t=600$~Myr to the final value at $t=700$. The velocity perturbation induced by object~C is shown on Figure~\ref{fig:posvel}. The model qualitatively reproduces the local amplification of the velocity gradient at the position of VirgoHI21, and in particular the ``S-shape" of the velocity profile visible in the Position-Velocity diagram (see Fig.~\ref{fig:pv-vhi21}). Tuning the parameters of the simulation may help to further reproduce the exact kinematical feature. Obviously other objects, such as NGC~4262 -- the gas--rich galaxy to the East in Figure~\ref{fig:vhi21} -- might have also crossed the trajectory of the tidal tail and slightly interacted with it. In any case, whatever the real explication is, the ``S-shape" of the velocity profile of VirgoHI21 is not inconsistent with a tidal origin. \bigskip A second critical issue is the identification of the interloper responsible for the collision. In our high-velocity scenario, the interloper now lies at a projected distance of 400~kpc to the WNW of NGC~4254. A massive spiral is in fact present near this position: NGC~4192 (M~98). This galaxy, which is seen close to edge-on, has a maximum rotation velocity corrected for inclination of 236~km~s$^{-1}$, significantly higher than that of NGC~4254 (193~km~s$^{-1}$, according to the HyperLeda database), making it compatible with the 1.5:1 mass ratio used in our simulation. In our model, the interloper is today approaching us with a radial velocity of 715~km~s$^{-1}$ w.r.t. the target spiral. In the real Universe, NGC~4192 is indeed approaching, but with a relative velocity of $\sim 2000$~km~s$^{-1}$ with respect to NGC~4254, i.e. much larger than in our model. This difference could be explained by the tidal field of the cluster which is not taken into account in our model. Adding it would introduce too many additional free parameters since the radial position and tangential velocity of the studied objects with respect to the cluster are unknown. The cluster gravitational field can modify some details in the properties of tidal tails in the long term \citep{Mihos04}, but not their fundamental characteristics. However the cluster tidal field can have a more dramatic effect on the relative orbit of the distant galaxy pair over the nearly 750~Myr period since the encounter. Indeed, a cluster with the mass typical of Virgo, $M_\mathrm{C}=10^{15}$~M$_\sun$, at a distance $R_\mathrm{C}=1$~Mpc from the pair, and a typical separation between the two galaxies $d\sim 300$~kpc (which is the average 3-D separation from the time of the encounter until today) would create a tidal acceleration \begin{equation} a_\mathrm{t} \simeq \frac{G M_\mathrm{C} d} {{R_\mathrm{C}} ^3} \end{equation} and, over a timescale of $T=750$~Myr, a relative velocity impulsion of \begin{equation} \Delta V_\mathrm{t} \simeq \frac{G M_\mathrm{C} d} {{R_\mathrm{C}} ^3} \times T \end{equation} which would come in addition to the relative velocity found in our model. The values above imply $\Delta V \simeq 1100$~km~s$^{-1}$. This additionnal velocity difference would account for the observed relative velocities of NGC~4192/4254, and goes in the right direction if NGC~4254 lies behind the Virgo Cluster center. NGC~4192 is thus a fully possible interloper, in spite of its large distance and relative velocity. This galaxy does not show a strongly disturbed HI disk on moderately deep HI maps \citep{Chung05}. One reason could be unfavorable internal/orbital parameters, for instance a retrograde orbit. Alternatively, faint tidal debris may be present, but only visible on deep HI observations, similar to those that revealed the existence of VirgoHI21. However, we do not intent to claim that NGC 4192 is the only possible interloper. Other combination of orbits/projection/age can certainly reproduce the properties of VirogHI21, and a galaxy flying away at $\sim 1000$~km~s$^{-1}$ during $\sim 1$~Gyr can be at a projected distance of 1~Mpc today, possibly even at the center of the Virgo Cluster. The number of massive interloper candidates is then large, making hard to identify the real culprit. This is anyway not required for our demonstration that VirgoHI21 can be a tidal debris, since we have shown that possible interlopers do exist. \begin{figure} \centering \includegraphics[width=\columnwidth]{f6.pdf} \caption{Position--Velocity diagram of the gas for Model 7 at $t=$750 Myr when it best reproduces the actual morphology and velocity field of the system VirgoHI21/NGC 4254. The arrow indicates the axis of the ``Position" direction. The PV diagram of a model where ``Object C" has been artificially added during the simulation is also shown. }\label{fig:posvel} \end{figure} \subsubsection{A kinematically decoupled Tidal Dwarf Galaxy?} The presence of a strong velocity gradient in a tidal tail, if not due to streaming motions, may actually pinpoint the presence of a gravitationally bound object that need not be a pre-existing dark-matter dominated galaxy. Massive substructures in tidal tails often become kinematically decoupled, self-gravitating and form new stars, becoming rotating Tidal Dwarf Galaxies. This is the case for VCC~2062, a TDG candidate in Virgo \citep{Duc07b}. VirgoHI21 appears as a gas condensation within a tidal tail; it is currently not a TDG since it is starless, but could be its gaseous progenitor. Whether stars will be formed in this structure later is questionable. If the surface column density has remained unusually very low during the first several hundreds of Myr after the formation of VirgoHI21, it is unlikely that star-formation is ignited later on. On the other hand, the dynamical collapse time of the cloud, $1/\sqrt{G \rho}$ \citep{elmegreen02}, is as large as 400 -- 500 Myr for an initially resting system and an average volume mass density in our modeled cloud of $\rho \sim 10^{-3}$~M$_{\sun}$~pc$^{-3}$ (this is also about the density of the real VirgoHI21 cloud assuming a vertical scale-height of $\sim$ 300~pc). Comparing this time scale to the age of the cloud, about 500~Myr (it appears in the model at $t=200-300$~Myr), one may conclude that VirgoHI21 would barely have had the time to collapse and form stars even in the most favorable conditions and incidentally that the absence of stars today is not dependent on an arbitrary choice of the threshold. In other words, the system could still be contracting under the effect of its internal gravity today, and begin to form stars later-on {\it if} its density comes to exceed the star formation threshold. However, the main argument against VirgoHI21 being {\it yet} a TDG fully responsible for the observed velocity gradient is the large dynamical mass inferred from the rotation curve: in galaxies made out of collisional debris, the dynamical mass should be of the same order as the luminous one, even if the presence of dark baryons may cause some differences between them \citep{B07}. In the case of VirgoHI21, the dynamical mass inferred from the velocity curve is more than a factor of 3000 greater than the luminous one, i.e. that of the HI component. Clearly, streaming motions provide a much more reasonable explanation for the kinematical feature observed near VirgoHI21, if indeed this object is of tidal origin. \subsubsection{Harrasment by the cluster field?} In the group environment, \cite{bekki05b} proposed a scenario in which the group tidal field is able to strip gas from HI-rich galaxies, explaining the presence of isolated intergalactic HI clouds. Following this idea, B05 presented a model in which the combined action of galaxy-galaxy interactions and the cluster tidal field produce debris with properties similar to Dark Galaxies. \citet{haynes07} even suggested that VirgoHI21 and the whole HI structure would result from just the long-term harassment by the large-scale cluster potential. However, the tidal field exerted by a structure of mass $M$ and typical scale $R$ scales as $M/R^3$. The tidal field of the Virgo Cluster ($10^{15}$~M$_\sun$, 1~Mpc) at the present distance of NGC~4254 is then more than ten times smaller than that of the interloper galaxy in our interaction model ($2\times 10^{12}$~M$_\sun$, 50~kpc). It is just unlikely that the cluster field can develop a tail as long as the galaxy interaction can do. The harassment process has a longer timescale than the galaxy pair interaction, but over long timescales the orientation changes, which hardly accounts for the single, thin and long tail around NGC~4254. This structure is more typical of a short and violent interaction like a close galaxy encounter than a weaker and longer process like the harassment by the global cluster field. \subsubsection{Ram pressure stripping?} \label{rampressure} Ram pressure exerted by the cluster hot gas may also expulse gas from the outer regions of spiral disks and create isolated HI clouds without any optical counterpart. Yet, this scenario suffers fundamental concerns, also pointed out by M07, in particular: \begin{itemize} \item structures known to result from ram-pressure stripping are rather short and thick as observed in Virgo \citep{crowl05,lucero05,vollmer06,chung07} and suggested by hydrodynamical simulations \citep[e.g.,][]{Roediger07}, while the HI bridge of VirgoHI21 is much thinner and longer. \item the kinematics is not directly explained too, in particular the reversing velocity gradient in the HI~bridge, which in the context of a tidal interaction results from streaming motions along a curved tail (in 3-D) more or less seen edge-on. This may also be the case for ram pressure but should be demonstrated by a model. \end{itemize} \cite{vollmer05} put forward the role of ram pressure, combined with a tidal interaction, in shaping the internal gas distribution, with its $m=1$ structure, and velocity field of NGC~4254. This partly explains why, in the innermost regions, the detailed kinematics of the spiral presents some differences with that of our model which did not take into account the intracluster medium. More recently, \cite{Kantharia07} presented a low radio frequency continuum map of the galaxy which is best explained invoking a ram pressure scenario. However so far, its possible contribution on the properties of the HI bridge and VirgoHI21 has not yet been investigated. The two above-mentioned papers do not claim that their origin is ram-pressure, and indeed our pure tidal model is able to reproduce these features provided that the HI disk was originally much more extended than the optical radius -- an hypothesis which was not adopted in \cite{vollmer05}. \subsubsection{A dark galaxy?} Showing that virgoHI21 {\it can} be a tidal debris taking the appearance of a dark galaxy does not directly rule out the possibility that it is a real dark galaxy. The dark galaxy hypothesis however suffers several difficulties unexplained so far: \begin{itemize} \item the maximal velocity gradient is not centered on the peak of the HI emission, but actually lies on one side of the VirgoHI21 cloud. This is unexpected for a rotating HI disk within a massive dark halo. One may propose that the on-going interaction with NGC~4254 causes this asymmetry in the velocity field, but that NGC~4254 is massive and close enough to induce such major disturbances remains to be shown. \item within the assumption that VirgoHI21 is a real dark galaxy, the HI bridge is a tidal tail expulsed from it and captured by the more massive galaxy NGC~4254 (M07). The gas in the HI bridge, falling onto NGC~4254 from 150~kpc away, is not expected to have the same velocity as the local gas settled in rotation: it could be on retrograde, polar or direct orbits but with different velocities. However, the base of the HI bridge has a radial velocity coherent with the outer disk of NGC~4254 to which it is morphologically connected: the velocity step between the base of the tidal tail and the outer disk is in fact less than $50$~km~$^{-1}$, much smaller than the circular velocity there \citep[see data in M07 and also][]{phookun93}. This further suggests that the HI material in the bridge comes from NGC~4254. \end{itemize} These facts are naturally explained by the tidal scenario proposed for VirgoHI21. Whether they can also be addressed with the Dark Galaxy hypothesis remains to be demonstrated, in particular with a numerical model. Although challenged, the putative existence of Dark Galaxies as massive as VirgoHI21, has fostered a number of follow-up works. As discussed by \citet{karachentsev06} such invisible ghost objects should tidally perturb galaxies in their neighborhood, explaining why a fraction of apparently isolated spiral stellar disks seem to show signs of an external perturbation. We note however that other mechanisms may account for them, such as accretion of diffuse gas \citep{bournaud05m1}. | 7 | 10 | 0710.3867 |
0710 | 0710.0675_arXiv.txt | We investigate the number and type of pulsars that will be discovered with the low-frequency radio telescope LOFAR. We consider different search strategies for the Galaxy, for globular clusters and for galaxies other than our own. We show an all-sky Galactic survey can be optimally carried out by {\em incoherently} combining the LOFAR stations. In a 60-day all-sky Galactic survey LOFAR can find over a thousand pulsars, probing the local pulsar population to a very deep luminosity limit. For targets of smaller angular size, globular clusters and galaxies, the LOFAR stations can be combined coherently, making use of the full sensitivity. Searches of nearby northern-sky globular clusters can find large numbers of low luminosity millisecond pulsars (eg.\ over 10 new millisecond pulsars in a 10-hour observation of M15). If the pulsar population in nearby galaxies is similar to that of the Milky Way, a 10-hour observation can find the 10 brightest pulsars in M33, or pulsars in other galaxies out to a distance of 1.2Mpc. | Since the discovery of the first four pulsars with the Cambridge radio telescope, an ongoing evolution of telescope systems has doubled the number of known radio pulsars roughly every 4 years. The next step in radio telescope evolution will be the use of large numbers of low-cost receivers that are combined to form an interferometer or a single dish. These telescopes, LOFAR \citep{ls07}, the Allen Telescope Array \citep{bow07} and the SKA \citep{kram04}, create new possibilities for pulsar research. Here we outline and compare strategies for targeting normal and millisecond pulsars, both in the disk and globular clusters of our Galaxy, and in other galaxies. | Because of its large area and wide beam on the sky, LOFAR probes the local population of pulsars to a very deep luminosity limit. A 60-day Galactic survey at 140MHz can find over a thousand new pulsars, disclosing the local low-luminosity population. If we add all antennae coherently the sensitivity increases even further; with this setup, millisecond pulsars in nearby globular clusters can be detected to much lower flux limits than previously possible. Assuming the pulsar population in other galaxies is similar to that in ours, we can detect periodicities or giant pulses from extragalactic pulsars up to several Mpcs away. | 7 | 10 | 0710.0675 |
0710 | 0710.3730_arXiv.txt | Models of terrestrial planet formation in the presence of a migrating giant planet have challenged the notion that hot-Jupiter systems lack terrestrial planets. We briefly review this issue and suggest that hot-Jupiter systems should be prime targets for future observational missions designed to detect Earth-sized and potentially habitable worlds. | Since the discovery of the first extrasolar planets \citep{wolszczan,mayor}, astronomical techniques and observational baselines have advanced to the point where over 200 extrasolar planetary systems have been identified \citep{butler}. Most detected exoplanets are in the giant planet mass range and it is now clear that our solar system is but one variant within a great diversity of planetary system architectures. One of the most surprising discoveries has been of a population of giant planets, the so-called \emph{hot-Jupiters}, found orbiting in a region of extreme insolation very close ($r < 0.1~\mathrm{AU}$) to their central stars and well within the radius of the original nebular snowline ($r \approx 3 - 5~\mathrm{AU}$) where giant planets are thought to form \citep{pollack}. Hot-Jupiters are not uncommon: they amount to about a quarter of exoplanet discoveries, and are thought to provide evidence that protoplanets can migrate over large radial distances via tidal interactions with the protoplanetary disk \citep[e.g.][]{lin1,lin2,ward2,nelson1}. Since the disk gas is observed to disperse within the first few Myr of the system's existence \citep{haisch}, giant planets must form and migrate through the inner system within this period, which is considerably less than the $\sim$~10--100~Myr thought to be required to complete terrestrial planet formation \citep{chambers2,kleine,halliday,obrien}. Test particle studies have shown that terrestrial planets external to a hot-Jupiter would have stable orbits \citep{jones}, and \citet{raymond1} have shown that they should be able to form, in the presence of a giant planet already at $\sim 0.1$~AU, from any available pre-planetary material with a period ratio roughly $>$~3. However, until recently it has been a common assumption that terrestrial planets could not have formed in hot-Jupiter systems due to the disruptive effect of the giant planet's migration which is deemed to have cleared the inner system of planet-forming material \citep[e.g.][]{armitage1}, prompting claims that the observed abundance of hot-Jupiters could be used to constrain the general abundance of habitable planets \citep{ward1}, and even their galactic location \citep{lineweaver1}. This picture has been challenged by the work of two groups who have modeled terrestrial planet formation concurrently with, and following, an episode of giant planet migration \citep{fogg1,fogg2,fogg3,fogg4,raymond2,mandell}. Their findings suggest that inner system solids disks are \emph{not} destroyed by the intrusion of a migrating giant planet and that terrestrial planet formation can resume in the aftermath and run to completion. In this paper, we briefly describe our model of terrestrial planet formation and show some typical results. | Our models predict that terrestrial planets might be routinely expected in hot-Jupiter systems, including within their habitable zones, and may be detectable by forthcoming missions such as Kepler, Darwin, SIM PlanetQuest and TPF. | 7 | 10 | 0710.3730 |
0710 | 0710.1503_arXiv.txt | {Observations of water lines are a sensitive probe of the geometry, dynamics and chemical structure of dense molecular gas. The launch of Herschel with on board HIFI and PACS allow to probe the behaviour of multiple water lines with unprecedented sensitivity and resolution.} {We investigate the diagnostic value of specific water transitions in high-mass star-forming regions. As a test case, we apply our models to the AFGL\,2591 region.} {A multi-zone escape probability method is used in two dimensions to calculate the radiative transfer. Similarities and differences of constant and jump abundance models are displayed, as well as when an outflow is incorporated.} {In general, for models with a constant water abundance, the ground state lines, i.e., $\mathrm{1_{10}}$-$\mathrm{1_{01}}$, $\mathrm{1_{11}}$-$\mathrm{0_{00}}$, and $\mathrm{2_{12}}$-$\mathrm{1_{01}}$, are predicted in absorption, all the others in emission. This behaviour changes for models with a water abundance jump profile in that the line profiles for jumps by a factor of $\sim$\,10-100 are similar to the line shapes in the constant abundance models, whereas larger jumps lead to emission profiles. Asymmetric line profiles are found for models with a cavity outflow and depend on the inclination angle. Models with an outflow cavity are favoured to reproduce the SWAS observations of the $\mathrm{1_{10}}$-$\mathrm{1_{01}}$ ground-state transition. PACS spectra will tell us about the geometry of these regions, both through the continuum and through the lines.} { It is found that the low-lying transitions of water are sensitive to outflow features, and represent the excitation conditions in the outer regions. High-lying transitions are more sensitive to the adopted density and temperature distribution which probe the inner excitation conditions. The Herschel mission will thus be very helpful to constrain the physical and chemical structure of high-mass star-forming regions such as AFGL\,2591.} | We have constructed models to examine the excitation of water in the high-mass star-forming region AFGL\,2591. Depending on the adopted density, temperature and abundance structure, a completely different set of line profiles and strengths is found. Hence, the line profiles are very sensitive to the adopted physical and chemical structure. We have found that ({\it i}) the ground-state transitions 1$_\mathrm{10}$-1$_\mathrm{01}$, 2$_\mathrm{12}$-1$_\mathrm{01}$ and 1$_\mathrm{11}$-0$_\mathrm{00}$, with relatively low upper energy levels ($\lesssim$\,110\,K), become highly optically thick in the outer regions. The line profiles for these transitions, are mainly dominated by the emission from the outer regions, and are therefore not useful to put constraints on the water abundance in the inner regions. However, ({\it ii}) the emission from lines with higher upper energy levels is dominated by the emission originating in the inner regions, and are therefore useful to probe the water abundance in the warm inner regions. ({\it iii}) For models with an outflow cavity, the outflow feature (blue peak less strong than the red peak) is best seen in the ground-state transitions of o- and p-H$_\mathrm{2}$O. ({\it iv}) The influence of a moderate disk (few 100 AU in size) in the centre of the AFGL\,2591 region does not change the water line profiles and strengths within the Herschel beam. The Herschel mission will thus greatly help to understand the structure of high-mass protostellar objects, and consequently the formation process of high-mass stars. | 7 | 10 | 0710.1503 |
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0710 | 0710.1359_arXiv.txt | In a systematic study, we compare the density statistics in high-resolution numerical experiments of supersonic isothermal turbulence, driven by the usually adopted solenoidal (divergence-free) forcing and by compressive (curl-free) forcing. We find that for the same rms Mach number, compressive forcing produces much stronger density enhancements and larger voids compared to solenoidal forcing. Consequently, the Fourier spectra of density fluctuations are significantly steeper. This result is confirmed using the $\Delta$-variance analysis, which yields power-law exponents $\beta\!\sim\!3.4$ for compressive forcing and $\beta\!\sim\!2.8$ for solenoidal forcing. We obtain fractal dimension estimates from the density spectra and $\Delta$-variance scaling, and by using the box counting, mass size and perimeter area methods applied to the volumetric data, projections and slices of our turbulent density fields. Our results suggest that compressive forcing yields fractal dimensions significantly smaller compared to solenoidal forcing. However, the actual values depend sensitively on the adopted method, with the most reliable estimates based on the $\Delta$-variance, or equivalently, on Fourier spectra. Using these methods, we obtain $D\!\sim\!2.3$ for compressive and $D\!\sim\!2.6$ for solenoidal forcing, which is within the range of fractal dimension estimates inferred from observations ($D\!\sim\!2.0\dots2.7$). The velocity dispersion to size relations for both solenoidal and compressive forcings obtained from velocity spectra follow a power law with exponents in the range $0.4\dots0.5$, in good agreement with previous studies. | Observations provide velocity dispersion to size relations for various molecular clouds (MCs), which document the existence of supersonic random motions on scales larger than $\sim\!0.1\,\mathrm{pc}$ \citep[e.g.,][]{Larson1981,Myers1983,PeraultFalgaronePuget1986,SolomonEtAl1987,FalgaronePugetPerault1992,HeyerBrunt2004}. These motions are associated with compressible turbulence \citep[e.g.,][]{ElmegreenScalo2004,ScaloElmegreen2004,MacLowKlessen2004} in the interstellar medium \citep{Ferriere2001} and exhibit a single turbulent cascade or spatially separated coexisting inertial ranges \citep{PassotPouquetWoodward1988} similar to the kinetic energy cascade of incompressible \citet{Kolmogorov1941c} turbulence. However, there are various physical processes (e.g., self-gravity, magnetic fields, nonequilibrium chemistry) and especially the compressibility of the gas, that alter the scaling laws \citep[e.g.,][]{Fleck1996} and statistics \citep[e.g., intermittency corrections measured by][]{HilyBlantFalgaronePety2008} established for incompressible turbulence. The physical origin and characteristics of the turbulent fluctuations are still a matter of debate. To advance on the question of how turbulence isdriven in the interstellar medium, we present results of high-resolution numerical experiments of supersonic isothermal turbulence comparing two distinct and extreme ways of driving the turbulence in a systematic study: 1) solenoidal forcing (divergence-free or rotational forcing), and 2) compressive forcing (curl-free or dilatational forcing). Various numerical and analytical studies have provided important insight into the statistics of supersonic isothermal turbulence \citep[e.g.,][]{PorterPouquetWoodward1992,Vazquez1994,PadoanNordlundJones1997,PassotVazquez1998,StoneOstrikerGammie1998,MacLow1999,Klessen2000,OstrikerStoneGammie2001,BoldyrevNordlundPadoan2002,LiKlessenMacLow2003,PadoanJimenezNordlundBoldyrev2004,JappsenEtAl2005,BallesterosEtAl2006,KritsukEtAl2007,LemasterStone2008}. Most of these studies use purely solenoidal or weakly compressive kinetic energy injection mechanisms (forcing) to excite turbulent motions. In the present study, we aim at comparing the usual case of solenoidal (divergence-free) forcing with the case of fully compressive (curl-free) forcing. The actual way of turbulence production in real MCs is expected to be far more complex compared to what we can model with the present simulations, probably consisting of a convolution of various agents producing turbulence, and mixtures of solenoidal and compressive modes \citep[e.g.,][]{ElmegreenScalo2004,MacLowKlessen2004}. Here, we systematically investigate the extreme cases of purely solenoidal versus purely compressive energy injection. Analyzing the density correlation statistics and fractal structure obtained in our hydrodynamic simulations, we show that compressive forcing leads to significantly steeper density fluctuation spectra and consequently to fractal dimensions of the turbulent gas structures, that are significantly smaller compared to the usually adopted solenoidal forcing. We use Fourier analysis, $\Delta$-variance analysis, structure functions, the fractal mass size, box counting, and perimeter area methods to obtain fractal dimension estimates. We apply the $\Delta$-variance analysis to both our 3-dimensional data and to 2-dimensional projections, and the perimeter area method to projections and slices through the turbulent density structures supporting the result of a significantly smaller fractal dimension for compressive forcing compared to solenoidal forcing. Although compressive forcing yields significantly smaller fractal dimensions than solenoidal forcing, the estimated fractal dimensions are in the range $2.0\dots2.7$ consistent with observational estimates \citep[e.g.,][]{ElmegreenFalgarone1996,SanchezEtAl2007} We explain our numerical method, construction of solenoidal and compressive forcing fields and fractal analysis techniques in Section~\ref{sec:methods}. In Section~\ref{sec:results}, we show that our results are consistent with previous studies using solenoidal forcing, whereas compressive forcing yields much stronger density contrasts and consequently leads to significantly smaller fractal dimensions. In Section~\ref{sec:conclusions}, we summarize our conclusions. | \label{sec:conclusions} We have presented results of two high-resolution ($1024^3$ grid cells) hydrodynamic simulations of supersonic isothermal turbulence driven to rms Mach numbers $\mathcal{M}\!\sim\!5.5$. The first simulation uses the typically adopted solenoidal (divergence-free) forcing to excite turbulent motions, whereas the second one uses compressive (curl-free) forcing. We have shown that compressive forcing yields much stronger density contrasts compared to solenoidal forcing for the same rms Mach number. This implies that the turbulence production mechanism leaves a strong imprint on compressible turbulence statistics, especially altering the density statistics. Our results particularly suggests that the mixture of solenoidal and compressive modes of the turbulence forcing must be taken into account. We summarize our results as follows: \begin{itemize} \item The velocity Fourier spectra exhibit power laws in the inertial range for solenoidal and compressive forcing. The slopes obtained for both forcings are significantly steeper ($\sim\!1.9$) compared to the Kolmogorov slope ($5/3$), in agreement with previous studies \citep[e.g.,][]{KritsukEtAl2007,SchmidtEtAl2008} and in agreement with velocity dispersion to size relations inferred from observations \citep[e.g.,][]{Larson1981,FalgaronePugetPerault1992,HeyerBrunt2004,PadoanEtAl2006}. \item From the integral of the density fluctuation Fourier spectra and from the asymptotic behavior of the 2nd order density structure function, we obtained the standard deviation of the density distribution $\sigma_\rho$. Compressive forcing yields a standard deviation $\sim\!$ three times larger compared to solenoidal forcing, in agreement with the results found in our previous study analyzing density probability distribution functions \citep{FederrathKlessenSchmidt2008} and in agreement with the studies by \citet{PassotVazquez1998}, \citet{KritsukEtAl2007}, \citet{BeetzEtAl2008} and \citet{SchmidtEtAl2008}. \item The density fluctuation Fourier spectra are significantly steeper for compressive forcing in the inertial range compared to solenoidal forcing. Consistent results were obtained using complementary analysis methods, i.e., by comparing the $\Delta$-variances \citep{OssenkopfKripsStutzki2008a} and the 2nd order structure functions of the density field. Our estimates of density spectra for solenoidal forcing are in agreement with previous studies, e.g., the weakly magnetized super-Alfv\'enic supersonic MHD models by \citet{PadoanEtAl2004} and \citet{KowalLazarianBeresnyak2007}, and consistent with the hydrodynamic estimates by \citet{KritsukNormanPadoan2006}. Although a comparison with observational results must be regarded with caution due to systematic uncertainties, our results for solenoidal and compressive forcing are in the range of inferred scaling exponents by observations \citep[e.g.,][]{BenschStutzkiOssenkopf2001}. \item From the scaling of the density fluctuation Fourier spectra and the $\Delta$-variance applied to the 3-dimensional data and applied to 2-dimensional projections, we obtained fractal Hurst exponents following the analysis by \citet{StutzkiEtAl1998}. This implies fractal box counting and fractal perimeter area dimensions significantly smaller for compressive forcing compared to solenoidal forcing (see Table~1). \item We analyzed the density structure using the fractal mass size method as introduced by \citet{KritsukEtAl2007}. Compressive forcing yields a smaller fractal mass dimension compared to solenoidal forcing. The mass size method is, however, particularly sensitive to the temporal fluctuations of density peaks. Given the large uncertainties, our results using this method are roughly consistent with the estimates by \citet{KritsukEtAl2007} and \citet{KowalLazarian2007}. \item We analyzed the fractal density structure using the box counting method described in Section~\ref{sec:box-counting} and the perimeter area method (Section~\ref{sec:perimeter-area}) applied to projections and slices. We recover the significant differences between solenoidal and compressive forcing inferred from the density spectra and $\Delta$-variance analysis. However, the box counting dimension varies strongly with the defining density threshold. The perimeter area dimensions obtained from slices are roughly consistent with the computed perimeter area dimensions from the $\Delta$-variance given the systematic uncertainties (of order $\sim\!0.1$ for fractal dimension estimates) comparing different methods. The range of fractal dimensions obtained is consistent with the observations analyzed by \citet{ElmegreenFalgarone1996} suggesting an overall fractal dimension of interstellar clouds in the range $D\!\sim\!2.3\pm0.3$. \end{itemize} | 7 | 10 | 0710.1359 |
0710 | 0710.3454_arXiv.txt | {There are now four low mass X-ray binaries with black holes which show twin resonant-like HFQPOs. Similar QPOs might have been found in Sgr A$\sp *$. I review the power spectral density distributions of the three X-ray flares and the six NIR flares published for Sgr A$\sp *$ so far, in order to look for more similarities than just the frequencies between the microquasar black holes and Sgr A$\sp *$. The three X-ray flares of Sgr A$\sp *$ are re-analysed in an identical way and white noise probabilities from their power density distributions are given for the periods reported around $\sim$ 1100 s. Progress of the resonant theory using the anomalous orbital velocity effect is summarized. | % \label{sect:intro} Quasi-periodic oscillations (QPOs) are believed to arise from variations of the accretion flow around compact objects, i.e. white dwarfs, neutron stars and black holes. As far as black holes are concerned, Remillard and McClintock (2006) have compiled a list of a total of 40 sources which show up in galactic X-ray binaries, 20 of which show a dark companion with a mass largely exceeding the mass of a neutron star but apparently limited to about 18 solar masses. Out of the 20 sources with established dynamically measured masses 16 sources show low frequency QPOs, whereas high frequency QPOs (HFQPOs, $\nu$ $>$ 10 Hz) have been detected in 8 sources. | 7 | 10 | 0710.3454 |
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0710 | 0710.4274_arXiv.txt | We performed MHD simulations of very light bipolar jets with density contrasts down to $10^{-4}$ in axisymmetry, which were injected into a medium of constant density and evolved up to $200$ kpc ($200\: r_{\mathrm{j}}$) full length. These jets show weak and roundish bow shocks as well as broad cocoons and thermalize their kinetic energy very efficiently. We argue that very light jets are necessary to match low-frequency radio observations of radio lobes as well as the bow shocks seen in X-rays. Due to the slow propagation, the backflows and their turbulent interaction in the midplane are important for a realistic global appearance. | During the last years, simulations of extragalactic jets with reasonable resolution and realistic sizes became computationally feasible, which makes comparisons between simulated and observed properties possible \citep{Saxton2002,Zanni2003,Carvalho2005,KrauseVLJ2,ONeill2005}. Unfortunately, the direct physical variables and the observed properties are rather hard to link, which leaves simulations with a wide range of parameters. Simulations are mainly governed by the initial setup of the density ratio between jet and ambient gas, the Mach number and the magnetic field. If the magnetic field is not dynamically dominant (though important), the density contrast is the most dominant parameter, but may be one of the hardest to measure. The thermal jet pressure has turned out to be of little importance in the very light jet limit \citep{KrauseVLJ1}. As the (kinetic) power of a jet can be estimated from energies in X-ray bubbles, typical values of velocity, lifetime, jet radius and cluster gas densities indicate that density contrasts of $10^{-2}$ to $10^{-4}$ (or even lower) are necessary to describe real sources. Parameter studies support this further, if the global jet/cocoon/bow shock properties are compared. Thus, we concentrate on the very light jets with magnetic fields as another important ingredient. | 7 | 10 | 0710.4274 |
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0710 | 0710.3381_arXiv.txt | We present optical and X-ray data for the first object showing strong evidence for being a black hole in a globular cluster. We show the initial X-ray light curve and X-ray spectrum which led to the discovery that this is an extremely bright, highly variable source, and thus must be a black hole. We present the optical spectrum which unambiguously identifies the optical counterpart as a globular cluster, and which shows a strong, broad [O III] emission line, most likely coming from an outflow driven by the accreting source. | Since the early days of X-ray astronomy, there has been considerable debate over whether globular clusters contained black holes. With the discovery of Type I X-ray bursts from all globular clusters in the Milky Way with bright X-ray sources (starting with Grindlay et al. 1976), and their subsequent explanation as episodes of thermonuclear burning on the surfaces of neutron stars (Woosley \& Taam 1976; Swank et al. 1977), it became clear that there was no evidence for any accreting black holes in the Milky Way's globular cluster system. Intepretations of the observations have been taken in two directions. One is simply that given only 13 bright X-ray sources in the Milky Way's globular cluster system, it is not so unlikely for them all to have neutron star accretors, especially in light of the fact that about 10 times as many neutron stars as black holes are expected to be produced for most stellar initial mass functions. The alternative is that dynamical effects are responsible for ejecting black holes from globular clusters. Severe mass segregation is likely to take place for globular cluster black holes, as they should be many times heavier than all the other stars in the cluster. This can lead to the formation of a ``cluster within a cluster'' where the heaviest stars (i.e. the black holes) feel negligble effects from the other stars in the cluster, which in turn leads to a cluster with a short evaporation timescale (Spitzer 1969). Numerical calculations have found that this evaporation can be accelerated further due to binary processes (e.g Portegies Zwart \& McMillan 2000). Early results from the Chandra X-ray Observatory gave new hope that globular cluster black holes might be detectable, by opening up the window of looking in other galaxies. Previously, only ROSAT could resolve point sources in other galaxies, and its localization of sources was generally not good enough to allow for unique identification of optical counterparts. The first few years of Chandra observations revealed several extragalactic globular cluster X-ray sources brighter than the Eddington limit for a neutron star (e.g. Angelini et al. 2001; Kundu et al. 2002), but a globular cluster may contain multiple bright neutron stars (as does, for example M~15 in our own galaxy -- White \& Angelini 2001), and that the quality of X-ray spectra available from Chandra for even the brightest extragalactic sources is insufficient to make phenomenological determinations that a source has a black hole accretor. It was pointed out that only large amplitude variability could prove that we were seeing the emission from a single source, rather than multiple sources (Kalogera, King \& Rasio 2003). Furthermore, the optical catalogs used to identify globular cluster counterparts to X-ray sources have been predominantly photometric catalogs, sometimes made even without color selections being used to ensure that that the optical source in question truly is a globular cluster. Most studies done to date have focused on HST images of the central regions of elliptical galaxies with high specific frequencies of globular clusters. In these regions, and with the angular resolution of HST, contamination will be rare, especially if color cuts are used to ensure that the contribution of background quasars is minimized. In the halos of galaxies, the surface density of real globular clusters will drop, and contamination will be a more serious problem. In either case, when one is looking for conclusive proof that an object is a globular cluster black hole, spectroscopic confirmation that the object is a globular cluster is essential -- the fractional contamination of the X-ray sources due to background AGN will be more serious at very high fluxes, corresponding to luminosities above $10^{39}$ ergs/sec than it will at lower levels consistent bright neutron star accretors. Furthermore, the optical to X-ray ratios for globular cluster black holes near the Eddington limit and background quasars are quite similar. | We have shown for the first time clear evidence of a globular cluster black hole, on the basis of strong, highly variable X-ray emission from a source in a spectroscopically confirmed globular cluster. Based on the X-ray spectrum, the characteristic variability, and the [O III] emission in the optical spectrum, the black hole is most likely a stellar mass object accreting far faster than its Eddington rate. Given that it is difficult to develop a scenario in which a globular cluster could have both a stellar mass black hole in a binary and an intermediate mass black hole, this argues against the idea that all globular clusters contain intermediate mass black holes. This discovery also motivates future searches for quiescent stellar mass black holes in the Milky Way's globular clusters; these may be hiding among the X-ray sources currently classified as cataclysmic variable stars. Radio emission should be detectable only from quiescent stellar mass black holes, and should be the one feasible discriminant between the two classes of systems. | 7 | 10 | 0710.3381 |
0710 | 0710.4935_arXiv.txt | Intercluster filaments negligibly contribute to the weak lensing signal in general relativity (GR), $\gamma_{N}\sim 10^{-4}-10^{-3}$. In the context of relativistic modified Newtonian dynamics (MOND) introduced by Bekenstein, however, a single filament inclined by $\approx 45^\circ$ from the line of sight can cause substantial distortion of background sources pointing towards the filament's axis ($\kappa=\gamma=(1-A^{-1})/2\sim 0.01$); this is rigorous for infinitely long uniform filaments, but also qualitatively true for short filaments ($\sim 30$Mpc), and even in regions where the projected matter density of the filament is equal to zero. Since galaxies and galaxy clusters are generally embedded in filaments or are projected on such structures, this contribution complicates the interpretation of the weak lensing shear map in the context of MOND. While our analysis is of mainly theoretical interest providing order-of-magnitude estimates only, it seems safe to conclude that when modeling systems with anomalous weak lensing signals, e.g. the ``bullet cluster" of Clowe et al., the ``cosmic train wreck" of Abell 520 from Mahdavi et al., and the ``dark clusters" of Erben et al., {\it filamentary structures might contribute} in a significant and likely complex fashion. On the other hand, {\it our predictions of a (conceptual) difference in the weak lensing signal could, in principle, be used to falsify MOND/TeVeS} and its variations. | \label{intro} Without resorting to cold dark matter (CDM), the modified Newtonian dynamics (MOND) paradigm \citep{Mond3, mondnew} is known to reproduce galaxy scaling relations like the Tully-Fisher relation \citep{tully}, the Faber-Jackson law \citep{faber} and the fundamental plane \citep{fundamental}) as well as the rotation curves of individual galaxies over five decades in mass \citep{spiral1,mondref1,mondref2,mondref3,mondref4,mondref5,escape}. In particular, the recent kinematic analysis of tidal dwarf galaxies by \cite{debris} is very hard to explain within the classical CDM framework while it is in accordance with MOND \citep{tidal1,tidal2}. In addition, observations of a tight correlation between the mass profiles of baryonic matter and dark matter in relatively isolated (field) galaxies at all radii \citep{insight2,insight} are most often interpreted as supporting MOND. Nevertheless, in rich clusters of galaxies, the MOND prescription is not enough to explain the observed discrepancy between visible and dynamical mass \citep{neutrinos2,tevesfit,asymmetric}. At very large radii, the discrepancy is about a factor of $2$, meaning that there should be as much dark matter (mainly in the central parts) as observed baryons in MOND clusters. One solution is that neutrinos have a mass at the limit of detection, i.e. $\sim2$ eV, which can solve the bulk of the problem of the missing mass in galaxy clusters, but other issues remain \citep{group}. These $2$ eV neutrinos have also been invoked to fit the angular power spectrum of the cosmic microwave background (CMB) in relativistic MOND \citep{tevesneutrinocosmo}, and are thus part of the only consistent MOND cosmology presented so far. In the following, we will refer to this model as the MOND hot dark matter ($\mu$HDM) cosmology \citep{tevesfit}\footnotemark\footnotetext[6]{Note, however, that one could also switch to sterile neutrinos with masses of a few eV \citep[e.g.][]{sterile,maltoni1,maltoni2} and that massive (sterile) neutrinos are not indispensable within certain covariant formulations of modified gravity, e.g. the $V\Lambda$ model, which can mimic the effects of neutrinos in clusters and cosmology as well as the behavior of a cosmological constant \citep{vector}.}. On the other hand, strange features have recently been discovered in galaxy clusters, which are hard to explain, such as the ``dark matter core" devoid of galaxies at the center of the ``cosmic train wreck" cluster Abell 520 \citep{abell520} and others \citep{darkcluster,bullet}. Here, we consider the possibility that this kind of features could be due to the gravitational lensing effects generated by an intercluster filament in a universe based on tensor-vector-scalar gravity \citep[TeVeS;][]{teves}, one possible relativistic extension of MOND \citep[cf.][]{tv1,tv2,vector}. However, we are not performing a detailed lensing analysis of any particular cluster in the presence of filaments, but rather provide a proof of concept that the influence of filaments could be much less negligible in a MONDian universe than within the framework of general relativity (GR). Filaments are among the most prominent large-scale structures of the universe. From simulations in $\Lambda$CDM cosmology, we know that almost every two neighboring clusters are connected by a straight filament with a length of approximately $20-30$ Mpc \citep{LCDMfilament}. For instance, the dynamics of field galaxies, which are generally embedded in such filaments, as well as their weak lensing properties are persistently influenced by this kind of structures, generally encountering accelerations of about $0.01-0.1\times 10^{-10}$ m s$^{-2}$. Filaments also cover a fair fraction of the sky, much larger than the covering factor of galaxy clusters. Thus, there is a good chance that filaments might be superimposed with other objects on a given line of sight, hence affecting the analysis of observational data like, for example, weak lensing shear measurements. Such recent studies prompted us to investigate the possibility that, in the context of MOND, end-on filamentary structures could be responsible for creating anomalous features in reconstructions of weak lensing convergence maps such as the peculiar ``dark matter core" devoid of galaxies in Abell 520 \citep{abell520}. Short straight filaments are structures which, at the best, are partially virialized in two directions perpendicular to their axis. According to \cite{LCDMfilament}, a filament generally corresponds to an overdensity of about $10-30$, having a cigar-like shape. Furthermore, filamentary structures tend to have a low-density gradient along their axis and, in the perpendicular directions, they have a nearly uniform core which tapers to zero at larger radii, usually about $2-5$ times their core radius. Since filaments are typically much longer than their diameter, we shall approximately treat them as infinite uniform cylinders of radius $R_{f}=2.5$ $h^{-1}$ Mpc. Lacking a MOND/TeVeS structure formation $N$-body simulation (with or without substantially massive neutrinos), we shall adopt the naive assumption that filamentary structures have roughly the same properties in MOND and in CDM, which will be justified in \S \ref{app}. Deriving expressions for the TeVeS deflection angle and setting up a cosmological background, we conclude that the order of magnitude of the TeVeS lensing signal caused by filaments is compatible with that of the previously mentioned observed anomalous systems. In addition, we find that there is fundamental difference between GR and MOND/TeVeS for cylindrically symmetric lens geometries (see Fig. \ref{fig1}); in contrast to GR, the framework of MOND/TeVeS allows us to have image distortion and amplification effects where the projected matter density is equal to zero. As for a more realistic approach, we also consider a model where the filament has a fluctuating density profile perpendicular to its axis. Compared to the uniform model, we find that the lensing signal in this case is smaller, but still of the same order, taking into account that the filamentary structures may be inclined to the line of sight by rather small angles ($\theta\lesssim 20^{\circ}$). Finally, we demonstrate the impact of filaments onto the convergence map of other objects by considering superposition of such structures with a toy cluster along the line of sight. Again, our results show an additional contribution comparable to that of a single isolated filament. \begin{figure} \centering \includegraphics[trim=-20 0 0 0,width=\linewidth]{figures/Fig1.pdf} \caption{Light deflection by an infinitely elongated cylinder of constant mass density; the unperturbed photon traveling along the $z$-direction passes the filament at the distance $y$ (impact parameter) from the filament's axis and is deflected by the angle $\hat\alpha$. The line density of the filament is assumed to be constant, $\lambda = M/L = \rho \pi R_f^2$, where $\rho$ is the volume density and $R_f$ is the cylinder's radius.} \label{fig1} \end{figure} | \label{discussion} In this work, we have analyzed the gravitational lensing effect by filamentary structures in TeVeS, a relativistic formulation of the MOND paradigm. For this purpose, we have set up two different cosmological models in TeVeS: the so-called $\mu$HDM cosmology including massive neutrinos on the order of $2$ eV which have already been proposed as a remedy for the discrepancies between dynamical and visible mass on cluster scales \citep{neutrinos2} as well as for the CMB \citep{tevesneutrinocosmo}, and a simple minimal-matter cosmology accounting for a universe which is made up of baryons alone. Encouraged by several HDM simulations and the fact that filamentary structures are generic, we have assumed that the properties of such structures, i.e. their shape and relative densities, are similar in CDM and MOND/TeVeS scenarios independent of the particularly used cosmological background. Modeling these filaments as infinite uniform mass cylinders, we have derived analytic expressions for their lensing properties in MOND/TeVeS and Newtonian/GR gravity. Regardless of the actual used cosmological background, we have shown that TeVeS filaments can account for quite a substantial contribution to the weak lensing convergence and shear field, $\kappa \sim \gamma \sim 0.01$, as well as to the amplification bias, $A^{-1}\sim 1.02$, which is even true outside but close $(y\sim 2R_{f})$ to the projected ``edges" of the filament's matter density. Exploring a simple oscillating density model of a filament and its surrounding area, we have found that the lensing signal in this case is generally smaller, but can still be of the same order, taking into account that the filamentary structures may be inclined to the line of sight by rather small angles ($\theta\lesssim 20^{\circ}$). In addition, we find that there is fundamental difference between GR and MOND/TeVeS considering idealized cylindrically symmetric lens geometries: wherever the projected matter density is zero, there will be no distortion as well as no amplification effects, i.e. image and source will look exactly the same. In the context of MOND/TeVeS, however, this changes as one can have such effects in these regions. Finally, we have demonstrated the impact of filaments onto the convergence map of other objects by considering superposition with a toy cluster along the line of sight. Again, our results have shown an additional contribution comparable to that of a single isolated filament and that the contribution pattern of filaments can be generally quite complex. Here we have considered the lensing signal generated by single filaments alone. Simulating the cosmic web in a standard $\Lambda$CDM cosmology, \cite{dolag} have found a shear signal $\gamma\sim 0.01-0.02$ along filamentary structures, which seems quite similar to what MOND/TeVeS can do. Note, however, that this signal is entirely dominated by the simulation's galaxy clusters, with the filament's signal being much smaller, approximately on the order of $10^{-4}-10^{-3}$. Although our analysis is mainly of theoretical interest, the above result points to an interesting possibility concerning recent measurements of weak lensing shear maps. For instance, the weak shear signal in the ``dark matter peak" of Abell 520 \citep{abell520} is roughly at a level of $0.02$, % which is comparable to what filaments could produce in MOND/TeVeS, but not in Newtonian gravity (also cf. \cite{wedding}).% Therefore, we conclude that filamentary structures might actually be able to cause such anomalous lensing signals within the framework of MOND/TeVeS. In principle, the predicted difference in the weak lensing signal could also be used to test the validity of modified gravity. As several attempts to detect filaments by means of weak lensing methods have failed so far, e.g. the analysis of Abell $220$ and $223$ by \cite{dietrich}, this might already be a first hint to possible problems within MOND/TeVeS gravity. On the other hand, shear signals around $\gamma\sim 0.01$ are still rather small to be certainly detected by today's weak lensing observations, and lacking $N$-body structure formation simulations in the framework of MOND/TeVeS, we cannot even be sure about how filaments form and how they look like in a MONDian universe compared to the CDM case. Clearly, more investigation is needed to gain a better understanding about the impact of filamentary structures. | 7 | 10 | 0710.4935 |
0710 | 0710.3809_arXiv.txt | The bipolar morphology of the planetary nebula (PN) K 3-35 observed in radio-continuum images was modelled with 3D hydrodynamic simulations with the adaptive grid code {\sc yguaz\'u-a}. We find that the observed morphology of this PN can be reproduced considering a precessing jet evolving in a dense AGB circumstellar medium, given by a mass loss rate $\dot{M}_{csm}=5\times 10^{-5}~\mathrm{M_{\odot} yr^{-1}}$ and a terminal velocity $v_\mathrm{w}=10~\mathrm{km s^{-1}}$. Synthetic thermal radio-continuum maps were generated from numerical results for several frequencies. Comparing the maps and the total fluxes obtained from the simulations with the observational results, we find that a model of precessing dense jets, where each jet injects material into the surrounding CSM at a rate $\dot{M}_j=2.8\times 10^{-4}~\mathrm{M_{\odot}\ yr^{-1}}$ (equivalent to a density of $8 \times 10^{4}~\mathrm{cm}^{-3}$), a velocity of $1\,500~\mathrm{km~s^{-1}}$, a precession period of 100~yr, and a semi-aperture precession angle of 20\degr\ agrees well with the observations. | \label{sec:intro} The evolution between the end of the asymptotic giant branch (AGB) and the planetary nebula (PN) phases has for a long time been a poorly understood link in the late stages of intermediate-mass stars (1 -- 8 M$_\odot$). It is in this transition phase where the fast stellar wind of the emerging PN interacts with the slow wind from its precursor AGB star (Kwok, Purton, \& Fitzgerald 1978), shaping the final morphology of the PN. Other ingredients that can contribute to the final shape of a PN are the presence of interacting binary stars (e.g. Morris 1987; Balick \& Frank 2002; Soker \& Bisker 2006) or magnetic fields (e.g. Garc\'\i a-Segura \& L\'opez 2000; Blackman et al. 2001; Garc\'\i a-Segura, L\'opez \& Franco 2005). Thus, the study of objects in this transition phase can give us important clues about the physical mechanisms responsible for the different morphologies observed in PNe. K 3-35 is a very young PN with a characteristic S-shaped emission morphology that suggests the presence of precessing bipolar jets (Aaquist \& Kwok 1989; Aaquist 1993; Miranda et al. 2000; 2001). The detection of OH and H$_2$O maser emission as well as the presence of CO and HCO$^+$ emission suggest that K~3-35 departed from the proto-PN phase only a few decades ago (Miranda et al. 2001; Tafoya et al. 2007). Miranda et al. (2001) estimate a dynamical age of $\leq$ 15 years for the ionised inner core, which expands at $\sim$25 km~s$^{-1}$. For the jets, assuming a modest jet velocity of $\sim$100 km~s$^{-1}$, an age of $\sim$800 years is obtained. Therefore, the jet formation in K~3-35 occurred during the proto-PN phase. In this paper we show that the radio morphology of K 3-35 can be explained by a precessing jet model. | Several 3D hydrodynamical simulations were carried out employing the adaptive grid code {\sc yguaz\'u-a}, in order to model both the morphology and the thermal radio emission of the planetary nebula K 3-35. After analysing our results, we found that the bipolar structure of this PN can be described as the result of the interaction of a dense jet (with an initial number density of $8\sim 10^4$~cm$^{-3}$ or $\dot{M}_j=2.8\times 10^{-4} M_{\odot} yr^{-1}$) moving into a dense environment, previously swept up by the AGB wind of the central star. The `S' morphology shown by K~3-35 in 8.3 GHz radio-continuum images \citep{lfm98,lfm01,gomez03} can be reproduced if the modelled jet precesses with a period of 100~yrs on a cone with a half-opening angle of 20\degr. For an integration time of 40~yrs, the simulated jet has a total length of $1.8\times 10^{17}$~cm, which is equivalent to 2.4\arcsec \citep{lfm01,gomez03}considering an estimated distance of 5 kpc to K 3-35 and also the orientation of this object \citep{uscanga07}. This time is almost 20 times smaller than the one given by \citet{lfm01} for the jet, where a slower velocity was assumed. However, it is only $2.7$ times larger than the dynamical age of the inner core \citep[also in][]{lfm01}, suggesting that the two events are more related than previously thought, and further supporting the idea that K 3-35 is a young object. Synthetic radio-continuum maps were generated from our numerical results. These maps show that the predicted morphologies and fluxes are in reasonable agreement with the observations. At an integration time of 40~yrs, the obtained spectral index is the one of optically thin emission. For an integration time of 50~yrs, the observed change of the spectral index with frequency (Aaquist 1993) is also reproduced by our simulation. A direct comparison between the observational and numerical results is given in Fig. \ref{fcomp}, where we show the observed and synthetic (for an integration time of 40~yrs) 8.3 GHz radio maps. We must note that the values for the velocity and mass loss rate employed in the simulation for the jet and the CSM seem to be rather high. It is difficult for AGB and post-AGB starts to launch jets at velocities of the order of $1\,500\,\mathrm{km~s}^{-1}$ (although Riera et al. 2003 have reported this kind of velocity for the outflows of the PPN 3-1475). Besides, the mass injection is also a bit extreme, in 40 years both jets have injected $0.02~\mathrm{M}_{\odot}$ into the surrounding CSM. This scenario might be explained in terms of a binary system, if the jet is produced by a companion accreting material (at a rate about ten times higher than $\dot{M}_j$) from a massive star (the AGB progenitor that produced the density distribution of the circumstellar material). The primary in the last 40 years has lost $0.2~\mathrm{M}_{\odot}$. This means that it started with an envelope containing this mass, which appears to correspond to an AGB star rather than a post-AGB star (even though at the present time the star has already evolved to the planetary nebula stage). Furthermore, the CSM is very dense. With the parameters employed, in a radius of $10^{17}$~cm, a mass of 1 M$_{\odot}$ is contained, implying a massive PN. There is observational evidence that favours a very dense surrounding CSM. However, the total mass derived from HCO$^+$ observations (Tafoya et al. 2007) is quite low ($\sim 0.017~$M$_\odot$), so that the molecular emission from this molecule would be confined to a small, possibly shock excited volume, in order to be consistent with our much higher mass CSM. These values imply that this scenario would be plausible if it is a short lived event. Clearly, a better determination of these parameters could be done with proper motion studies and better distance determinations. Notwithstanding the extreme values for the employed parameters, it is important to note that a simple model of a precessing dense jet moving into an also dense CSM, successfully reproduces the observed morphology of the PN K 3-35, obtaining a qualitatively and quantitatively good agreement between the model predictions and the observations, although this scenario would be a short lived event. \begin{figure} \includegraphics[width=\columnwidth]{fig4.eps} \caption{Comparison between observational and numerical results. The top panel shows the 8.3~GHz radio continuum image of K 3-35 \citep{lfm01,gomez03}. The bottom panel shows the tilted synthetic radio emission map at the same frequency, obtained for an integration time of 40~yr. An inclination of 36\degr between the precession axis and the plane of the sky was considered. The simulated map was smoothed with the observed beam of $0.2\arcsec\times 0.2\arcsec$ and it is shown in the same scale that the observed one. The contours corresponds to the levels 0.18, 0.31, 0.56, 1.0, 1.78, 3.16, 5.62, and 10.00 mJy beam$^{-1}$. The arrows indicate the jet direction (projected on the plane of the sky) for integrations times of 0 and 40~yr. The last one is coincident with the P.A. measured by \citet{lfm01}.} \label{fcomp} \end{figure} | 7 | 10 | 0710.3809 |
0710 | 0710.2722_arXiv.txt | We present the results of the multi-frequency observations of radio outburst of the microquasar Cyg X-3 in February and March 2006 with the Nobeyama 45-m telescope, the Nobeyama Millimeter Array, and the Yamaguchi 32-m telescope. Since the prediction of a flare by RATAN-600, the source has been monitored from Jan 27 (UT) with these radio telescopes. At the eighteenth day after the quench of the activity, successive flares exceeding 1 Jy were observed successfully. The time scale of the variability in the active phase is presumably shorter in higher frequency bands. We also present the result of a follow-up VLBI observation at 8.4 GHz with the Japanese VLBI Network (JVN) 2.6 days after the first rise. The VLBI image exhibits a single core with a size of $<8$ mas (80 AU). The observed image was almost stable, although the core showed rapid variation in flux density. No jet structure was seen at a sensitivity of $T_b = 7.5\times 10^5$ K\@. | Cyg X-3 is a famous X-ray binary including a black hole candidate (e.g., \cite{schalinski1998}). This object is classified as a microquasar due to its bipolar relativistic jet accompanied by radio flares. Because it is located on the Galactic plane at a distance of about 10 kpc (e.g., \cite{predehl2000}) and obscured by intervening interstellar matter, it has been observed mainly in radio and X-ray. Its giant radio flares have been observed once every several years since its initial discovery \citep{gregory1972,braes1972}. The peak flux densities in the radio flares have often increased up to levels of 10 Jy or more at centimeter wave (e.g., \cite{waltman1994}). The radio emission seems to be correlated with hard X-ray emission, and not with soft X-ray emission \citep{mccollough1999}. Although the radio emission arises through synchrotron process of relativistic electrons in the jet \citep{hjellming1988}, the millimeter behavior during the flares is not yet established. An observation at a shorter wavelength and with a higher time resolution is desirable to understand the mechanism of the flares. The quenched state of Cyg X-3, in which the radio emission is suppressed below 1 mJy, is a possible precursor of flares (e.g., \cite{waltman1994}). In January 2006, this quenched state was detected in monitoring observations with the RATAN-600 radio telescope \citep{trushkin2006}. The source has been monitored from MJD$=53762$ (Jan 27 2006 in UT) with the Nobeyama 45-m radio telescope (NRO45), the Nobeyama Millimeter Array (NMA), and the Yamaguchi 32-m radio telescope (YT32). We detected the initial state, or rising phase, of the radio flare of Cyg X-3 at MJD$= 53768$ (February 2 2006) and observed successive flares exceeding 1 Jy \citep{tsuboi2006}, which turned out to be the beginning of an active phase lasting more than 40 days . In this paper, radio observations with NRO45, NMA, YT32, and the Japanese VLBI Network (JVN) are reported. The observation procedures are summarized in section 2. The light curves and the spectral evolution observed with NRO45, NMA and YT32 are shown in section 3, together with the result of JVN\@. Detailed discussion based on these observations will be published as separate papers. | 7 | 10 | 0710.2722 |
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0710 | 0710.0727_arXiv.txt | { Gamma-ray burst are thought to be produced by highly relativistic outflows. Although upper and lower limits for the outflow initial Lorentz factor $\Gamma_0$ are available, observational efforts to derive a direct determination of $\Gamma_0$ have so far failed or provided ambiguous results. As a matter of fact, the shape of the early-time afterglow light curve is strongly sensitive on $\Gamma_0$ which determines the time of the afterglow peak, i.e. when the outflow and the shocked circumburst material share a comparable amount of energy. We now comment early-time observations of the near-infrared afterglows of GRB\,060418 and GRB\,060607A performed by the REM robotic telescope. For both events, the afterglow peak was singled out and allowed us to determine the initial fireball Lorentz, $\Gamma_0\sim400$. | The early stages of gamma-ray burst (GRB) afterglow light curves display a rich variety of phenomena at all wavelengths and contain significant information which may allow determining the physical properties of the emitting fireball. The launch of the \textit{Swift} satellite \citep{Neil04}, combined with the development of fast-slewing ground-based telescopes, has hugely improved the sampling of early GRB afterglow light curves. Since many processes work in the early afterglow, it is often difficult to model them well enough to be able to determine the fireball characteristics. The simplest case is a light curve shaped by the forward shock only. This case is particularly interesting because, while the late-time light curve is independent of the initial conditions (the so-called self-similar solution), the time at which the afterglow peaks depends on the original fireball Lorentz factor $\Gamma$, thus allowing a direct measurement of this fundamental parameter \citep{SP99}. The short variability timescales, coupled with the nonthermal GRB spectra, indeed imply that the sources emitting GRBs have a highly relativistic motion \citep{Ruderman75,Fenimore93,Piran00,Lith01}, to avoid suppression of the high-energy photons due to pair production. This argument, however, can only set a lower limit to the fireball Lorentz factor. Late-time measurements (weeks to months after the GRB) have shown $\Gamma \sim \mbox{a few}$ \citep{Frail97,Taylor05}, but a direct measure of the initial value (when $\Gamma$ is expected to be $\sim 100$ or more) is still lacking. We present here the NIR early light curves of the GRB\,060418 and GRB\,060607A afterglows observed with the REM robotic telescope\footnote{\url{http://www.rem.inaf.it}} \citep{Zerb01,Chinc03} located in La Silla (Chile). These light curves show the onset of the afterglow and its decay at NIR wavelengths as simply predicted by the fireball forward shock model, without the presence of flares or other peculiar features. A detailed discussion of these data has also been reported by \citet{Mol07} and \citet{Mal07}. \begin{figure} \centering \includegraphics[width=11cm]{CovinoS_2007_01_fig01.ps} \\ \centering\includegraphics[width=11cm]{CovinoS_2007_01_fig02.ps} \caption{NIR and X-ray light curves of GRB\,060418 (upper panel) and GRB\,060607A (lower panel). The dotted lines show the models of the NIR data using the smoothly broken power law (see Sect. \ref{modelling}). For GRB\,060418 the dashed line shows the best-fit to the X-ray data.\label{fig:lcs}} \end{figure} | The REM discovery of the afterglow onset has demonstrated once again the richness and variety of physical processes occurring in the early afterglow stages. The very fast response observations presented here provide crucial information on the GRB fireball parameters, most importantly its initial Lorentz factor. This is the first time that $\Gamma(t_{\rm peak})$ is directly measured from the observations of a GRB. The measured $\Gamma_0$ value is well within the range ($50 \la \Gamma_{0} \la 1000$) envisaged by the standard fireball model \citep{Piran00,Guetta01,Alicia02,Meszaros}. It is also in agreement with existing measured lower limits \citep{Lith01,zhang06}. Using $\Gamma_0=400$ we can also derive other fundamental quantities characterising the fireball of the two bursts. In particular, the isotropic-equivalent baryonic load of the fireball is $M_{\rm fb} = E/(\Gamma_0 c^2) \approx 7\times 10^{-4}\,M_\odot$, and the deceleration radius is $R_{\rm dec} \approx 2 c t_{\rm peak} [\Gamma(t_{\rm peak})]^2/(1+z) \approx 10^{17}$~cm. This is much larger than the scale of $\sim 10^{15}$~cm where the internal shocks are believed to power the prompt emission \citep{MesRees97,Rees94}, thus providing further evidence for a different origin of the prompt and afterglow stages of the GRB. | 7 | 10 | 0710.0727 |
0710 | 0710.0382_arXiv.txt | % We present results from a 150 ksec {\it Suzaku} observation of the Seyfert 1.5 NGC 3516 in October 2005. The source was in a relatively highly absorbed state. Our best-fit model is consistent with the presence of a low-ionization absorber which has a column density near 5 $\times$ 10$^{22}$ cm$^{-2}$ and covers most of the X-ray continuum source (covering fraction 96--100$\%$). A high-ionization absorbing component, which yields a narrow absorption feature consistent with Fe K {\sc XXVI}, is confirmed. A relativistically broadened Fe K$\alpha$ line is required in all fits, even after the complex absorption is taken into account; an additional partial-covering component is an inadequate substitute for the continuum curvature associated with the broad Fe line. A narrow Fe K$\alpha$ emission line has a velocity width consistent with the Broad Line Region. The low-ionization absorber may be responsible for producing the narrow Fe K$\alpha$ line, though a contribution from additional material out of the line of sight is possible. We include in our model soft band emission lines from He- and H-like ions of N, O, Ne and Mg, consistent with photo-ionization, though a small contribution from collisionally-ionized emission is possible. | The 6.4 keV Fe K$\alpha$ emission line has long been known to be an important diagnostic of the material accreting onto supermassive black holes. The Compton reflection hump, frequently seen in Seyfert spectra above $\sim$7~keV and peaking near 20--30 keV (Pounds et al.\ 1990), indicates that Seyferts' Fe K lines may have an origin in optically thick material, such as the accretion disk. Observations with {\it ASCA} indicated that many Fe K$\alpha$ lines were broad (FWHM velocities up to $\sim$0.3$c$) and asymmetrically skewed towards lower energies, implying an origin in the inner accretion disk; the line profile is sculpted by gravitational redshifting and relativistic Doppler effects (e.g., Tanaka et al.\ 1995, Fabian et al.\ 2002). However, {\it XMM-Newton} and {\it Chandra} observations have been revealing a more complex picture. A narrow Fe K component (FWHM velocities $\sim$5000 km s$^{-1}$ or less) appears to be quite common; these lines' widths suggest emission from distant material, such as the outer accretion disk, the optical/UV Broad Line Region (BLR) or the molecular torus. Spectral observations in which the broad and narrow components are deconvolved are thus a prerequisite for using the Fe K line as a tracer of the geometry of the emitting gas. At the same time, there is strong evidence from X-ray and UV grating observations for the presence of ionized material in the inner regions of a large fraction of AGN (e.g., Blustin et al.\ 2005; McKernan et al.\ 2007). High-resolution spectroscopy shows the gas is usually outflowing from the nucleus; typical velocities are $\sim$ a few hundred km s$^{-1}$. Absorption due to a broad range of ionic species is commonly seen; and for many sources, there is evidence for several different photo-ionized absorbing components, as opposed to a single absorber, along the line of sight. In the Fe K bandpass, some Seyferts show evidence for absorption by H- or He-like Fe, indicating a zone of highly-ionized absorbing material (e.g., NGC 3783, Reeves et al.\ 2004). Cold absorbing gas, with line of sight columns in excess of the Galactic value, are routinely observed in Seyfert 2 AGN in accordance with unification schemes (Urry, Padovani 1995), and have also been reported in some Seyfert 1 AGN. Importantly, variations in column density on timescales from hours to years have been observed in both Seyfert 1 AGN (e.g., in I Zw 1, Gallo et al.\ 2007; see also Lamer et al.\ 2003, Puccetti et al.\ 2004) and Seyfert 2 AGN (Risaliti et al.\ 2002; Risaliti et al.\ 2005), suggesting that the absorbing circumnuclear material is not homogeneous in either Seyfert type, has a high tranverse velocity and occurs over a range of length scales. NGC 3516 is a well-studied, nearby (z = 0.008836; Keel 1996) Seyfert 1.5 AGN that can display strong 2--10 keV flux variability on timescales of hours to years (e.g., Markowitz, Edelson 2004). Previous X-ray spectroscopic observations of NGC 3516 e.g., Nandra et al.\ (1997) using {\it ASCA}, have indicated the presence of a broad Fe K line, and this source is known to also contain complex and ionized absorption. Numerous UV absorption lines, including N {\sc V}, C {\sc IV} and Si {\sc IV}, were observed with the {\it International Ultraviolet Explorer} (Ulrich, Boisson 1983); absorption line strengths vary on timescales as short as weeks as the absorber responds to variations in the ionizing flux (e.g., Voit et al.\ 1987). {\it Hubble Space Telescope} observations have revealed that this component of absorbing gas (henceforth called the ``UV absorber'') may consist of several distinct kinematic components (Crenshaw et al.\ 1998). X-ray spectra of NGC 3516 can exhibit evidence for large columns ($\gtrsim$10$^{22}$ cm$^{-2}$) of absorbing gas (e.g., Kolman et al.\ 1993), and the X-ray absorbers can display variability on timescales of years (e.g., Mathur et al.\ 1997). Using {\it Chandra} gratings data from observations in April 2001 and November 2001, Turner et al.\ (2005) observed K-shell absorption lines due to H- like Mg, Si and S, and He-like Si, evidence for a highly-ionized absorber, likely with column density $\gtrsim$10$^{22}$ cm$^{-2}$, outflowing at $\sim$1100 km s$^{-1}$. Simultaneous with these {\it Chandra} observations in 2001 were two {\it XMM-Newton} observations. Turner et al.\ (2005) modeled the continuum curvature of the two {\it XMM-Newton} EPIC spectra by including a partial covering, mildly-ionized absorber; the column density was $\sim$2.5 $\times$ 10$^{23}$ cm$^{-2}$, with a covering fraction of $\sim$50$\%$. However, the formal requirement for the broad Fe line was reduced, leading to uncertainty as to whether the broad Fe line really existed in NGC 3516. Spectral fitting using an instrument with a wide bandpass is thus necessary to remove such model degeneracies. In this paper, we report on an observation of NGC 3516 made with the {\it Suzaku} observatory in October 2005. The combination of the X-ray Imaging Spectrometer (XIS) CCD and the Hard X-ray Detector (HXD) instruments have yielded a broadband spectrum covering 0.3 to 76 keV, allowing us to deconvolve the various broadband emitting and absorbing components. Furthermore, the exceptional response of the XIS CCD and high signal-to-noise ratio of this observation have allowed us to study narrow emission lines in great detail. $\S$2 gives a brief overview of the {\it Suzaku} observatory, and describes the observation and data reduction. $\S$3 describes fits to the time-averaged spectrum. Variability analysis is briefly discussed in $\S$4. Flux-resolved spectral fits are discussed in $\S$5. In $\S$6, we describe a search for narrow red- or blue-shifted lines in the Fe K bandpass. The results are discussed in $\S$7, and a brief summary is given in $\S$8. | During the late 1990's, NGC 3516 typically displayed a 2--10 keV flux of $\sim$4--6 $\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$ (e.g., Markowitz, Edelson 2004). During the 2001 {\it XMM-Newton}/{\it Chandra} observations, however, the observed 2--10 keV flux was much lower: 1.6--2.3 $\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$ (Turner et al.\ 2005). Table 4 lists the inferred absorption-corrected 2--10 keV nuclear fluxes from the {\it XMM-Newton} observations, as well as during the 2005 {\it Suzaku} observation. In addition, Figure 10 shows the unfolded observed spectra for the {\it Suzaku} XIS and the 2001 {\it XMM-Newton} EPIC-pn data (see Turner et al.\ 2005 for details regarding the {\it XMM-Newton} data). The {\it Suzaku} observation apparently caught the source in a similar low level of nuclear flux as the 2001 observations. The observed 0.5--2.0 keV flux during the {\it Suzaku} observation, however, was $\sim$2--3 times lower than during the {\it XMM-Newton} observations, indicating that the source was still heavily obscured, and confirming that the complex absorption in this source cannot be ignored when fitting the broadband spectrum and modeling diskline components. \subsection{Power-law Components} The primary power-law observed in the hard X-rays is likely emission from a hot corona very close to the supermassive black hole, as seen in all Seyferts. The nature of the soft power-law component, however, is not as clear. It could represent nuclear emission scattered off optically-thin material, e.g., in the optical/UV Narrow Line Region (NLR). In the baseline model, the normalization of the soft power-law relative to that of the primary power-law was 4.2$\pm$0.4$\%$. Assuming a covering fraction of unity, this ratio is equal to the optical depth of the scattering material, indicating a column density of roughly 5 $\times$ 10$^{22}$ cm$^{-2}$, consistent with this notion, though the column density is somewhat too high to likely be associated with the NLR. It is interesting to note that this column density is similar to that obtained for the high-ionization absorber, suggesting the possibility that this absorbing component could be associated with a zone of scattering. Using {\it Chandra}-ACIS, George et al.\ (2002) found the extended circumnuclear gas to have a 0.4--2.0 keV flux of roughly 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$. However, that emission was studied over an annular extraction region 3$\arcsec$ to 10$\arcsec$ (0.6--1.8 kpc), and so that flux value is likely a lower limit to the 0.4--2.0 keV flux that Suzaku would observe (given the XRTs' 2$\arcmin$ HPD). In our baseline model, we found an unabsorbed 0.4--2.0 keV flux of 1 $\times$ 10$^{-12}$ erg cm$^{-2}$ s$^{-1}$, consistent with the notion that the soft power-law is scattered emission. In this case, the decrease in observed 0.5--2.0 keV flux from 2001 to 2005 could potentially be explained by the scattered emission responding to a recent decrease in nuclear continuum flux. However, this scenario would require the bulk of the scattered emission to lie within at most a few light years of the black hole, and the nuclear flux would have had to decrease between 2001 and 2005 (when the source was not observed by any major X-ray mission in the 2--10 keV band) then return to 2001 levels by the {\it Suzaku} observation. Alternatively, the soft power-law could be unobscured, ``leaked'' nuclear emission as part of a partial covering scenario. In this case, the primary absorber would obscure 96$\%$ of the sky as seen from the nuclear continuum source. The lack of significant variability in the 0.3--1.0 keV band could argue for the soft power-law to originate in scattered emission, since we might expect to observed variability of the same amplitude as the 2--10 keV band only if the soft power-law were leaked nuclear emission. However, this is far from certain, as the 0.3--1.0 keV band had a low count rate and the presence of the soft emission lines in the XIS spectrum could contribute to dilution of intrinsic variability in the soft power-law. A broadband observation spanning a larger observed flux range is needed to clarify this issue. The soft power-law could of course represent a blend of scattered emission plus leaked nuclear emission. We therefore conclude that the primary absorber has a covering fraction between 96--100$\%$. \subsection{Complex Absorption} We detect two zones of absorption: in addition to the primary absorber, which has a covering fraction of 96--100$\%$, there is the high-ionization absorber, which is assumed here to have a covering fraction of unity. The high-ionization absorber is potentially the same as that reported by Turner et al.\ (2005); we find a column density $N_{\rm H}$ of 4.0$^{+4.6}_{-3.1}$ $\times$ 10$^{22}$ cm$^{-2}$, consistent with the column density of 2 $\times$ 10$^{22}$ cm$^{-2}$ used by Turner et al.\ (2005), although we use a slightly higher ionization parameter (see Turner et al.\ 2007). Previous studies of NGC 3516, such as Netzer et al.\ (2002), have discussed in detail the UV absorber, responsible for H Ly$\alpha$, C {\sc IV} and N {\sc V} absorption features in {\it Hubble Space Telescope} spectra (Kraemer et al.\ 2002). In the X-ray band, discrete features associated with Mg {\sc VII--IX} and Si {\sc VII--IX} are expected from this component, but with the CCD resolution and with calibration-related artifacts near 1.7--1.8 keV in the XIS, such features are not detected by {\it Suzaku}. {\it Suzaku} has found the primary absorber of the hard X-ray continuum to be lowly-ionized (log($\xi$) = 0.3$\pm$0.1 erg cm s$^{-1}$), with a column density $N_{\rm H}$ of 5.5$\pm$0.2 $\times$ 10$^{22}$ cm$^{-2}$. It is possible that it is the same absorber that Turner et al.\ (2005) designated as the ``heavy'' partial-covering absorber, though we use a somewhat lower ionization parameter (see Turner et al.\ 2007). A new observation of NGC 3516 with {\it XMM-Newton} in 2006 October showed that the source had returned to similar $>$6 keV brightness and similar obscuration levels ($N_{\rm H}$ $\sim$ 2 $\times$ 10$^{23}$ cm$^{-2}$; covering fraction $\sim$ 45$\%$) as during the 2001 observations (Turner et al.\ 2007). Long-term changes in the covering fraction of the heavy absorber could explain the bulk of the spectral variability changes between the 2001, 2005 and 2006 observations. In this case, the column density has decreased by a factor of 4.5, while the covering fraction has increased from $\sim$40--60$\%$ to 96--100$\%$, from 2001 to 2005, and subsequently returned to approximately the same levels as in 2001 within the next 12 months. However, because we have not actually observed the entire eclipse associated with a specific, discrete blob of absorbing gas traversing the line of sight, it is not clear whether the covering fractions derived are associated with single, large blobs partially blocking the line of sight to the X-ray continuum source, or if the absorber consists of numerous, discrete blobs or has a filamentary or patchy structure. On the other hand, given the gaps between the 2001 and 2006 {\it XMM-Newton} and 2005 {\it Suzaku} observations, it is certainly plausible that these observations could have caught independent, discrete blobs or filaments with differing column densities and differing physical sizes and/or radial distances lying on the line of sight. To estimate the distance $r$ between the central black hole and the absorbing gas, we can use a definition of the ionization parameter $\xi$ = $L_{1-1000 {\rm Ryd}}$/($n$$r^2$), where $n$ is the number density. $L_{1-1000 {\rm Ryd}}$ is the 1--1000 Ryd illuminating continuum luminosity, and the value of the ionization parameter is taken to be the value of 2 erg cm s$^{-1}$ measured above. We estimate the maximum possible distance to the material by assuming that the thickness $\Delta$$r$ must be less than the distance $r$. The column density $N_{\rm H}$ = $n$$\Delta$$r$, yielding the upper limit $r$ $<$ $L_{1-1000 {\rm Ryd}}$/($N_{\rm H}$$\xi$). We estimate the 1-1000 Ryd flux from the baseline model to be 9.9 $\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$, which corresponds to $L_{1-1000 {\rm Ryd}} = 1.7 \times 10^{43}$ erg s$^{-1}$ (assuming $H_{\rm o}$ = 70 km s$^{-1}$ Mpc$^{-1}$ and $\Lambda_{\rm o}$ = 0.73). $r$ is thus $<$2 $\times$ 10$^{20}$ cm (180 light-years), a very loose upper limit encompassing both distances associated with the BLR ($\sim$10 light-days; Peterson et al.\ 2004) and a possible cold molecular torus at a 1 pc radius. Variability in the absorber properties between 2001--2005 and 2005--2006 thus imply a radial distance of at most a few light years in the case of clouds traversing the line of sight to the nucleus. In addition, in the case of a partial covering scenario, it is plausible that the absorber's size could be of the same order as that of the X-ray continuum source. The absorbing material could thus be e.g., associated with the base of an outflow or from dense clouds associated with magnetohydrodynamic disk turbulence (e.g., Emmering et al.\ 1992). NGC 3516's transition from an unobscured source to a moderately-obscured source in a 4 year span presents a challenge to standard Seyfert 1/2 unification schemes. If the obscuration in NGC 3516 is associated with an equatorial molecular torus usually invoked in Seyfert 1/2 unification schemes, then it is possible that during the {\it Suzaku} observation, the inner edge of the torus could have intersected the line of sight, but given NGC 3516's classification as a Seyfert 1, this could only occur if the torus opening angle were extremely small and the torus were not azimuthally symmetric. Alternatively, variations in column density and/or covering fraction could be due to fine structure in large-distance (tens of pc), non-equatorial filaments that traverse the line of sight (e.g., Malkan et al.\ 1998). \subsection{Fe K Emission Components and Compton Reflection} We have deconvolved the broadband emitting components, and determined that 1) the existence of the broad Fe line is robust in that it was required in all models for an adequate fit, and 2) a partial covering component could not mimic the curvature associated with a relativistic broad line. We note, for instance, that if we remove the diskline components from the baseline model and refit, not only is the fit worse ($\chi^2$ increases by over 170), but the value of $R$ becomes $\sim$ 3.2. This value is incompatible with the observed $EW$ of the narrow line unless the Fe abundance is extremely sub-solar ($\lesssim$0.3; see below). The best-fit disk inclination was typically $\lesssim$25$\arcdeg$. The inner radius was typically $\lesssim$5$R_{\rm g}$. The line energy was seen to be consistent with neutral to mildly-ionized Fe (up to Fe $\sim$ {\sc XX}; Kallman et al.\ 2004). The equivalent width with respect the primary continuum was 287$^{+49}_{-24}$ eV, consistent with the value of 431$^{+193}_{-172}$ eV obtained by Turner et al.\ (2005) for the April 2001 {\it XMM-Newton} observation, where the spectrum could be fit with a diskline component in addition to the complex absorbing components. The line energy of the narrow Fe K$\alpha$ line was also consistent with emission from neutral Fe. The intensity of the narrow line during the {\it Suzaku} observation is roughly 40$\%$ higher than that measured during the 2001 {\it XMM-Newton} and {\it Chandra} observations (Turner et al.\ 2002), possibly indicating that a substantial fraction of the Fe K line photons originate in a region $\lesssim$5 lt.-yr.\ in size. We measured a FWHM velocity line width for the narrow Fe K$\alpha$ line of $<$ 4900 km s$^{-1}$ (99$\%$ confidence level for two interesting parameters). This velocity does not rule out an origin in the BLR; Peterson et al.\ (2004) reported FWHM velocities for the H$\alpha$ and H$\beta$ lines of 4770$\pm$893 and 3353$\pm$310 km s$^{-1}$, respectively. However, we also cannot exclude a contribution from an origin in the putative molecular torus; there could potentially be a very narrow line component with FWHM velocity $\sim$ a few hundred km s$^{-1}$, but the XIS would be unable to separate it from the relatively broader line component. It is possible that the same material that absorbs the hard X-rays along the line of sight is responsible for producing the narrow Fe line. The material producing the Fe line cannot have a column substantially less than 10$^{\sim 22}$ cm$^{-2}$ or else there would be insufficient optical depth to produce a prominent Fe K line. The primary absorber, with its column density of 5.5 $\times$ 10$^{22}$ cm$^{-2}$ and low ionization state, is thus a plausible candidate for the narrow line origin. As an estimate of the Fe K$\alpha$ equivalent width expected in this case, we can assume an origin in optically-thin gas which completely surrounds a single X-ray continuum source and is uniform in column density, and use the following equation: \begin{equation} EW_{\rm calc} = f_{\rm c} \omega f_{\rm K\alpha} A \frac{\int^{\infty}_{E_{\rm K edge}}P(E) \sigma_{\rm ph}(E) N_{\rm H} dE}{P(E_{\rm line})} \end{equation} Emission is assumed to be isotropic. Here, $f_{\rm c}$ is the covering fraction, initially assumed to be 1.0. $\omega$ is the fluorescent yield, 0.34 (Kallman et al.\ 2004). $f_{\rm K\alpha}$ is the fraction of photons that go into the K$\alpha$ line as opposed to the K$\beta$ line; this is 0.89 for Fe {\sc I}. $A$ is the number abundance relative to hydrogen. We assumed solar abundances, using Lodders (2003). $P(E)$ is the spectrum of the illuminating continuum at energy $E$; $E_{\rm line}$ is the K$\alpha$ emission line energy. $\sigma_{\rm ph}(E)$ is the photo-ionization cross section assuming absorption by K-shell electrons only (Veigele 1973\footnote{http://www.pa.uky.edu/$\sim$verner/photo.html}). For $N_{\rm H}$ = 5.5 $\times$ 10$^{22}$ cm$^{-2}$, $EW_{\rm calc}$ = 29 eV, substantially lower than the observed $EW$ of 123$\pm$7 eV. We conclude that it is possible for the primary absorber to contribute to the observed line $EW$, but there is also likely a contribution from some other (non-continuum absorbing) material lying out of the line of sight, likely with column densities 10$^{\sim 23}$ cm$^{-2}$ (e.g., Matt et al.\ 2002). For instance, if the putative cold molecular torus does not intersect the line of sight, it could contribute to the observed $EW$. The 13$\%$ upper limit to ratio of the Compton shoulder/ narrow Fe K$\alpha$ core intensity was a 90$\%$ confidence limit only, and does not exclude at high confidence the possibility of Compton-thick material out of the line of sight. An additional possibility is that the material emitting the bulk of the line photons could be responding to a continuum flux that was higher in the past. For instance, if the putative molecular torus is located $\sim$ a pc or so from the black hole, the torus will yield a line $EW$ corresponding to the continuum flux averaged over the past few years. This situation is plausible for NGC 3516, as the 2--10 keV flux of NGC 3516 during $\sim$1998--2001 (Markowitz, Edelson 2004) was a factor of $\sim$1.5--2 times brighter than during 2005. We now turn our attention to properties of the Compton reflection continuum. {\it Suzaku} has observed other Seyferts to display reduced levels of variability in the PIN band compared to the 2--10 keV band, e.g., in MCG--6-30-15 (Miniutti et al.\ 2007). This behavior is thought to be caused by the presence of the relatively non-varying Compton reflection hump, which dilutes the observed $>$10 keV variability of the power-law component. Gravitational light-bending in the region of strong gravity has been invoked to explain the relative constancy of the reflection spectrum (Compton hump and Fe K diskline component) despite large variations in the coronal power-law flux in MCG--6-30-15, for instance (Miniutti et al.\ 2007). In the case of NGC 3516, the observed fractional variability amplitudes for the 2--10 and 12--76 keV bands were $F_{\rm var,2-10}$ = 9.2$\pm$0.3$\%$ and $F_{\rm var,12-76}$ $<$ 4.4$\%$, respectively. These measurements allow us to rule out the possibility that the Compton hump varies in concert with the power-law, since the variability amplitudes would be consistent in that case. The primary power-law and Compton hump contribute 44$\%$ and 56$\%$, respectively, of the total 12--76 keV flux. In the case of a constant Compton hump and variable power-law, $F_{\rm var,12-76}$ would then be equal to $F_{\rm var,2-10}$ / 2.25, or roughly 4.1$\%$. This is roughly consistent with the observed upper limit on $F_{\rm var,12-76}$, suggesting that the reflection component varies less strongly than the primary power-law over the course of the observation. To verify this, however, we would need to observe NGC 3516 over a larger % X-ray flux range than in the current {\it Suzaku} observation to potentially observe any significant variability in the PIN band. Finally, we discuss the origin of the material that gives rise to the observed Compton reflection hump. The primary and high-ionization absorbers lack the necessary column density, and are excluded. We next consider an origin in the same material that yields either the broad or narrow Fe lines. George, Fabian (1991) calculated that $R$ = 1 corresponds to an observed Fe K$\alpha$ line $EW$ (relative to a primary continuum with a photon index of 1.9) of 140 eV for neutral Fe, assuming an inclination angle of 25$\arcdeg$. However, George, Fabian (1991) used the elemental abundances of Morrison, McCammon (1983), where the Fe number abundance relative to hydrogen was $A_{\rm Fe}$ = 3.3 $\times$ 10$^{-5}$. More recent papers have slightly lower values of $A_{\rm Fe}$, 2.7--3.0 $\times$ 10$^{-5}$ (Lodders 2003; Wilms et al.\ 2000). The expected equivalent width corresponding to $R$ = 1 is thus 115--125 eV. In our baseline model, we found a best-fit value of $R$ = 1.7, which corresponds to an expected line $EW$ (relative to the primary continuum) of 200--215 eV, a value in between the observed $EW$s of the broad line (287 eV in the baseline model) and the narrow line (123 eV). It is thus not clear from this measurement alone whether the total Compton reflection continuum is associated with the broad line (disk), narrow line (a distant origin), or both. That is, while is it a possibility that at least some portion of the Compton reflection component is associated with the broad Fe K component, we cannot exclude the possibility that the narrow line contributes as well and that there is reflection off cold, distant material. For example, in $\S$3.3, we demonstrated that a model wherein there existed both blurred reflection from an ionized disk plus reflection from cold, distant material, such as the molecular torus, provided a good fit to the data. In addition, we demonstrated in this section that the observed $EW$ of the narrow Fe K line means we cannot rule out a contribution to the narrow Fe line, and to the reflection continuum as well, from Compton-thick material out of the line of sight. | 7 | 10 | 0710.0382 |
0710 | 0710.2343.txt | This paper analyses the radio properties of a subsample of optically obscured ($R\ge 25.5$) galaxies observed at 24$\mu$m by the {\it Spitzer Space Telescope} within the First Look Survey. 96 $F_{24\mu\rm m}\ge 0.35$~mJy objects out of 510 are found to have a radio counterpart at 1.4~GHz, 610~MHz or at both frequencies respectively down to $\sim 40\mu$Jy and $\sim 200\mu$Jy. IRAC photometry sets the majority of them in the redshift interval $z\simeq [1-3]$ and allows for a broad distinction between AGN-dominated galaxies ($\sim 47$\% of the radio-identified sample) and systems powered by intense star-formation ($\sim 13$\%), the remaining objects being impossible to classify. The percentage of radio identifications is a strong function of 24$\mu$m flux: almost all sources brighter than $F_{24\mu\rm m}\sim 2$~mJy are endowed with a radio flux at both 1.4~GHz and 610~MHz, while this fraction drastically decreases by lowering the 24$\mu$m flux level. The radio number counts at both radio frequencies suggest that the physical process(es) responsible for radio activity in these objects have a common origin regardless of whether the source shows mid-IR emission compatible with being an obscured AGN or a star-forming galaxy. We also find that both candidate AGN and star-forming systems follow (although with a large scatter) the relationship between 1.4~GHz and 24$\mu$m fluxes reported by Appleton et al. (2004) which identifies sources undergoing intense star formation activity. However, a more scattered relation is observed between 24$\mu$m and 610~MHz fluxes. On the other hand, the inferred radio spectral indices $\alpha$ indicate that a large fraction of objects in our sample ($\sim 60$\% of all galaxies with estimated $\alpha$) may belong to the population of Ultra Steep Spectrum (USS) Sources, typically 'frustrated' radio-loud AGN. We interpret our findings as a strong indication for concurrent AGN and star-forming activity, whereby the 1.4~GHz flux is of thermal origin, while that at 610~GHz mainly stems from the nuclear source. | The advent of the {\it Spitzer Space Telescope} has marked a fundamental milestone in our understanding of the assembly history of massive spheroidal galaxies, one of the major issues for galaxy formation models. The unprecedented sensitivity of the Multiband Imaging Photometer for {\it Spitzer} (MIPS) at 24$\mu$m has in fact for the first time allowed the detection at high redshifts of a population of Luminous and UltraLuminous Infrared Galaxies (LIRGs; ULIRGs) with huge infrared luminosities ($L_{\rm IR}>10^{11}L_{\odot}$). Such sources are underluminous at rest frame optical and UV wavelengths because they are reprocessing and radiating much of their energy in the IR (e.g. Sanders \& Mirabel 1996). As a consequence, LIRGs and ULIRGs at high redshifts had been missed so far either due to their extreme optical faintness or because previous infrared missions such as the InfraRed Astronomical Satellite (IRAS) or the Infrared Space Observatory (ISO) did not have enough sensitivity to push the observations beyond $z \sim 1$. Recent studies have shown that these objects, while relatively rare in the local universe (e.g. Sanders \& Mirabel 1996), become an increasingly significant population at higher redshifts (e.g. Le Floc'h et al. 2004; 2005; Lonsdale et al. 2004; Caputi et al. 2006; 2007) and likely dominate the luminosity density at $z>1$ (see e.g. Dole et al. 2006). Their space density, found to range between $10^{-4}$ and a few $10^{-5}$~Mpc$^{-3}$ according to the selection criteria adopted by different studies (see e.g. Caputi et al 2007; Daddi et al. 2007a, Magliocchetti et al. 2007a), is a factor of 10 to 100 higher than that of optically selected quasars in the same redshift range (e.g. Porciani, Magliocchetti \& Norberg 2004). Furthermore, clustering studies (e.g. Magliocchetti et al. 2007; 2007a; Farrah et al. 2006) prove that, at variance with their local counterparts, LIRGs and ULIRGs at $z\sim 2$ are associated with extremely massive ($M\simgt 10^{13} M_{\odot}$, where $M$ here refers to the dark matter) structures, only second to those which locally host very rich clusters of galaxies. Given their properties, it then appears clear that these sources represent a fundamental phase in the build up of massive galactic bulges, and in the growth of their supermassive black holes. Emission line diagnostics for bright ($F_{24\mu \rm m}\simgt 1$~mJy) mid-IR samples of LIRGs and ULIRGs at $z\sim 2$ in the near and mid-IR spectral regimes (see e.g. Yan et al. 2005; 2007; Weedman et al. 2006; 2006a; Brand et al. 2007; Martinez-Sansigre et al. 2006a) have shown these sources to be a mixture of obscured type1-type2 AGN and systems undergoing intense star formation activity. These findings are confirmed by photometric follow up mainly undertaken in the mid-IR and X-ray (both soft and hard) bands which also prove that the fraction of galaxies dominated by a contribution of AGN origin is drastically reduced at faint mid-IR fluxes (e.g. Brand et al. 2006; Weedman et al. 2006; Treister et al. 2006; Magliocchetti et al. 2007). Unfortunately, the exact proportion of AGN vs starforming dominated galaxies is still undetermined. This separation is further complicated by the existence of a noticeable number of mixed systems where both star formation and AGN activity significantly contribute to the IR emission. For instance, Daddi et al. (2007a) find that about 20\% of 24$\mu$m-selected galaxies in the GOODS sample show a mid-IR excess which is not possible to reconcile with pure star-forming activity. Such a fraction increases to $\sim 50-60$\% at the highest (stellar) masses probed by their study. Clearly, understanding how the mid-IR sources divide between starbursts, AGN and composite systems is now the next essential step in order understand the relationships amongst the formation and evolution of stars, galaxies and massive black holes powering AGNs within dusty environments and more generally within massive systems observed at the peak of their activity. This paper approaches the study of the population of optically faint luminous infrared galaxies from the point of view of their multifrequency radio emission. Diagnostics based on the radio signal steming from these sources can in fact provide precious information on the process(es) which are actively taking place within such systems. Enhanced radio activity can stem from supernova remnants associated with regions which are vigorously forming stars, or originate from nuclear activity (AGN-dominated sources). These two processes determine a rather different spectral behaviour at radio wavelengths: star-forming systems are in general characterized by radio spectra which feature power-law shapes with slope (hereafter called {\it radio spectral index} $\alpha$, defined as $F\propto \nu^{-\alpha}$, with $F$ radio flux and $\nu$ radio frequency) of the order of $\sim$0.7-0.8, while typical radio-loud quasars exhibit values for $\alpha$ between 0 and 0.5 even though, especially at high redshifts, there is a non negligible population of radio-loud sources with very high, $\alpha>1$, values (see e.g. De Breuck et al. 2000). Radio counterparts to the 510 optically faint mid-IR selected sources drawn from the whole First Look Survey sample (Fadda et al. 2006) have been searched in the overlapping region between MIPS and IRAC observations (Magliocchetti et al. 2007). Despite not having direct redshift estimates except for a handful of cases (e.g. Weedman et al. 2006; Yan et al. 2005; 2007; Martinez-Sansigre et al. 2006a), mid-IR photometry indicates that the overwhelming majority of such sources reside at redshifts $1.7\simlt z\simlt 2.5$ (\S~4; see also Brand et al. 2007; Houck et al. 2005; Weedman et al. 2006a).\\ The First Look Survey region provides an excellent laboratory for investigations of the radio properties of dusty galaxies set at redshifts $z\sim 2$ as its area has been observed at a number of radio frequencies down to very low flux densities (Condon et al. 2003; Morganti et al. 2004; Garn et al. 2007), therefore maximizing our chances of finding radio emitting galaxies. The layout of the paper is as follows. In \S2 we introduce the parent (mid-IR and radio) catalogues, while in \S3 we present the matching procedure leading to the sample of optically obscured {\it Spitzer}-selected sources with a radio counterpart at 1.4~GHz and/or 610~MHz. \S4 uses IRAC photometry to provide some information on the typology of such objects (i.e. whether mainly powered by an obscured AGN or by a starburst) and also -- where possible -- to assign them to a redshift interval. \S5 presents the results on the radio number counts both at 1.4~GHz and 610~MHz (\S5.1) and on the relationship between radio and 24$\mu$m emission (\S5.2). \S6 discusses our findings on the 1.4~GHz vs 610~MHz radio spectral indices for the objects in our sample, while \S7 summarizes our conclusions. | This paper has presented an analysis of the radio properties of a subsample of optically faint ($R\simgt 25.5$), 24$\mu$m-selected galaxies observed by {\it Spitzer} in the FLS (Magliocchetti et al. 2007). These objects have been cross-correlated with a number of radio catalogues which cover the same region of the sky, namely that of Condon et al. (2003) which probes 1.4~GHz fluxes brighter than $\sim 100 \mu$Jy, that of Garn et al. (2007) -- which probes 610~MHz fluxes brighter than $\sim 200 \mu$Jy -- and, on a smaller portion of the sky, that of Morganti et al. (2004) which reaches 1.4~GHz fluxes as faint as $\sim 40\mu$Jy. 70 optically faint {\it Spitzer} sources have been identified in the Condon et al. (2003) catalogue, 33 in the Morganti et al. (2004) dataset, while 52 are found in the survey performed by Garn et al. (2007). After performing a number of corrections to account for multiple identifications, sources erroneously split in the original {\it Spitzer} catalogue into different components and mid-IR objects with real radio counterparts at one of the two radio frequencies which were further away than the allowed ($10^{\prime\prime}$) matching radius, we end up with a sample of 96 radio-identified, optically faint, mid-IR emitting sources, 45 of which have an identification at both 1.4~GHz and 610~MHz. The fraction of radio identifications is a strong function of 24$\mu$m flux: almost all sources brighter than $F_{24\mu\rm m}\sim 2$~mJy are endowed with a radio flux at both 1.4~GHz and 610~MHz, while this fraction drastically decreases by lowering the flux level. IRAC photometry for all those sources which also have detected fluxes in at least one of the four 8$\mu$m, 5.8$\mu$m, 4.5$\mu$m and 3.6$\mu$m channels (64 out of 96), allows to classify them into two categories: obscured AGN (45 sources) and systems mainly powered by starformation activity, (SB, 13 objects). We also find five low-z (i.e. $z\simlt 0.5$ m51-type) interlopers, while the remaining 33 sources are unclassified. Furthermore, with the help of IRAC photometry it was possible to assign broad redshift intervals to all those sources (mostly AGN) which presented in the lowest 3.6$\mu$m and $4.5\mu$m wavelength channels of IRAC a 'bump' compatible with being produced by an evolved (old) stellar population. The majority ($\sim 66$\%) of these galaxies reside at redshifts $z\simgt 1$, in agreement with other studies mainly based on mid-IR and near-IR spectroscopy of optically faint, 24$\mu$m-selected galaxies (i.e. Weedman et al. 2006; Yan et al. 2005; 2007; Brand et al. 2007). We stress that this inferred fraction can only be considered as a lower limit to the real portion of faint {\it Spitzer} sources set at high redshifts. In fact, because of their characteristic spectral properties in the mid-IR regime, we expect the majority of star forming systems (too faint at the IRAC frequencies to have measurable 3.6$\mu$m and/or 4.5$\mu$m fluxes) to be indeed located in the $z\sim [1.6-2.7]$ redshift range. A small fraction of objects in our sample present radio morphologies such as jets and/or lobes compatible with them being identified as radio-loud AGNs. Interestingly enough, we find that for most of these few extended radio sources the mid-IR emission is associated to such peripheral regions rather than steming from the centre of radio activity, generally coinciding with the location of the AGN. However, the majority of the objects in our sample present unresolved radio images. A compared analysis of the radio number counts for optically obscured {\it Spitzer} sources indicates that the $\Delta N/\Delta F$ as estimated at 1.4~GHz can only be reconciled with what found at 610~MHz if the population under investigation is endowed with an average value for the radio spectral index (defined as $F_R\propto \nu_R^{-\alpha}$, where $F_R$ is the radio flux and $\nu_R$ the generic radio frequency) $\left<\alpha\right>\simgt 1$. Classical, 'flat spectrum' radio sources can confidently be excluded as the typical objects constituting our sample. Furthermore, we have found that the radio number counts of sources classified as AGN and of those identified as starburst galaxies are quite similar, evidence which suggests that radio emission in $z\sim 2$ {\it Spitzer} galaxies originates from similar process(es), despite of the different mid-IR emission. Direct investigations of the relation between 24$\mu$m and 1.4~GHz fluxes show that the overwhelming majority of those galaxies not excluded from our analysis because morphologically classified as radio-loud AGN follow the relationship identified for $0\simlt z\simlt 1$ star-forming objects by Appleton et al. (2004), although with a large scatter. This happens regardless of whether the galaxy has been classified as an AGN or a starforming system on the basis of its mid-IR colours. The distribution of 24$\mu$m vs 610~MHz fluxes is instead found to be more scattered. The majority of these radio-identified objects (26, a figure which rises to 34 if one also includes sources with estimated lower limits on $\alpha$) present very steep, $\alpha> 1$ radio spectral indices, some galaxies being endowed with $\alpha$'s as high as 2.5. This excess of galaxies with large $\alpha$ values is statistically significant as it corresponds to 60 per cent of our sample, to be compared with the 43 per cent found by considering {\it the whole population} of faint radio objects with an identification at both 610~MHz and 1.4~GHz.\\ Such very high figures for $\alpha$ would identify the corresponding sources as Ultra Steep Spectrum galaxies, generally high redshift radio-loud AGN. However this is in striking disagreement with what found for the relation between 24$\mu$m and 1.4~GHz fluxes for the very same objects, relation which would explain their (1.4~GHz) radio emission as mainly due to processes connected with star forming activity. A natural explanation to the above issues could be found by assuming that AGN and star-formation activity are concomitant in the majority of $z\sim 2$, {\it Spitzer} sources, at least in those which present enhanced radio emission. The radio signal steming from these systems would then simply be the combination of two components: a shallower one -- dominating the spectrum at 1.4~GHz -- due to processes connected with star formation, and a steeper one -- being responsible for most of the 610~MHz signal -- connected with AGN activity. This framework could then also explain why the 1.4~GHz emission in our sources follows that of star forming systems, while the same does not seem to happen in the case of 610~MHz fluxes. Various explanations can be found in the literature to account for the presence of USS: from slow adiabatic expansion losses in high-density environments (e.g. Klamer et al. 2006) to the scattering between CMB photons and relativistic electrons at $z\sim 2$ (e.g. Martinez-Sansigre et al. 2006). However, the results presented in this work might provide an alternative scenario. In fact, they suggest that high values for $\alpha$ might be due to the concomitant presence within the same systems of an AGN and of a star forming region: the AGN expansion would then be halted by the encounter with the cooler/denser sites in which star formation takes place. This would determine the 'strangling' of the AGN, causing its radio spectrum at low radio frequencies to steepen. Clearly, more theoretical work is needed in order to quantify the above issue and we are planning to present it in a forthcoming paper. For the time being, we note that intense star forming activity within a high redshift galaxy host of a USS has been recently reported by Hatch et al. (2007). From a more observational point of view, the 'ultimate truth' on these sources could only come from very high resolution (and, given their faintness, sensitivity) measurements, capable to clearly disentangle the emission related to the AGN to that associated to star forming regions. The advent of instruments such as ALMA will then provide the answers we need. | 7 | 10 | 0710.2343 |
0710 | 0710.2452_arXiv.txt | We present highlights and an overview of 20 {\em FUSE} and {\em HST}/STIS observations of the bright symbiotic binary \object{EG And}. The main motivation behind this work is to obtain spatially-resolved information on an evolved giant star in order to understand the mass-loss processes at work in these objects. The system consists of a low-luminosity white dwarf and a mass-losing, non-dusty M2 giant. The ultraviolet observations follow the white dwarf continuum through periodic gradual occultations by the wind and chromosphere of the giant, providing a unique diagnosis of the circumstellar gas in absorption. Unocculted spectra display high ionization features, such as the \ion{O}{6} resonance doublet which is present as a variable ({\em hourly} time-scales), broad wind profile, which diagnose the hot gas close to the dwarf component. Spectra observed at stages of partial occultation display a host of low-ionization, narrow, absorption lines, with transitions observed from lower energy levels up to $\sim$5 eV above ground. This absorption is due to chromospheric/wind material, with most lines due to transitions of \ion{Si}{2}, \ion{P}{2}, \ion{N}{1}, \ion{Fe}{2} and \ion{Ni}{2}, as well as heavily damped \ion{H}{1} Lyman series features. No molecular features are observed in the wind acceleration region despite the sensitivity of {\em FUSE} to H$_2$. From analysis of the ultraviolet dataset, as well as optical data, we find that the dwarf radiation does not dominate the wind acceleration region of the giant, and that observed thermal and dynamic wind properties are most likely representative of isolated red giants. | 7 | 10 | 0710.2452 |
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0710 | 0710.1876_arXiv.txt | {During the period 1966.5--2006.2 the 15GHz and 8GHz lightcurves of 3C454.3 (z=0.859) show a quasi-periodicity of $\sim$12.8 yr ($\sim$6.9 yr in the rest frame of the source) with a double-bump structure. This periodic behaviour is interpreted in terms of a rotating double-jet model in which the two jets are created from the black holes in a binary system and rotate with the period of the orbital motion. The periodic variations in the radio fluxes of 3C454.3 are suggested to be mainly due to the lighthouse effects (or the variation in Doppler boosting) of the precessing jets which are caused by the orbital motion. In addition, variations in the mass-flow rates accreting onto the black holes may be also involved.\\ | In recent years many works have been devoted to the study of the periodic properties discovered in extragalactic sources, especially in Blazars. These include the periodic (or quasi-periodic) variations of flux density (or luminosity) in optical bands (for example, for OJ287: Sillanp{\"a}{\"a} et al. \cite{Si88}, Lehto \& Valtonen \cite{Le96}, Valtonen et al. \cite{Va99}, Villata et al. \cite{Vi98}, Valtaoja et al. \cite{Va00}, Valtonen \cite{Va06a}, \cite{Va06b}; AO\,0235+16: Raiteri et al. \cite{Ra01}); periodic variations of flux density in radio bands (for example, for AO0235+16: Raiteri et al. \cite{Ra01}, 3C454.3: Ciaramella et al. \cite{Ci04}, Kudryavtseva \& Pyatunina \cite{Ku06}); the periodic changes of the position angle of the ejection of superluminal knots (for example, Qian et al. \cite{Qi01}, Britzen et al. \cite{Br01}, Abraham \& Carrara \cite{Ab00}, Lobanov \& Roland \cite{Lo05}); periodic changes of the trajectory of superluminal knots (for example, for 3C345: Qian et al. \cite{Qi91}, Lobanov \& Roland \cite{Lo05}, Klare et al. \cite{Kl05}), etc. Up to now, the most reliable periodic phenomenon observed in blazars may be the 12 year quasi-periodic optical variations of the BL Lac object OJ287 which has been convincingly determined in the more than 100 year long records of optical observations since 1890.\\ In the interpretations of these periodic phenomena various binary balck hole models have been invoked. For the periodic optical variations in OJ287 all the proposed models invoke a binary system with two supermassive black holes. But the mechanisms for the production of the optical outbursts can be divided into two classes: accretion models and lighthouse models. \begin{itemize} \item Accretion models\\ These models (Lehto \& Valtonen \cite{Le96}, Valtonen \cite{Va06a}, \cite{Va06b}) propose that the optical periodicity is caused by a precessing binary black hole system of two supermassive black holes with a large eccentric orbit. The observed periodic 'superflares' with double peaks are suggested to be due to the crossings of the secondary black hole into the accretion disk of the primary black hole during the pericenter passage. Thus in this class of models the optical emission is thermal (bremsstrahlung) and the observed flux increase reflects the enhanced accretion related to the disk crossings and the general enhancement of the accretion rate during the pericenter passage. Doppler boosting is not taken into account. Detailed models for a binary black hole system have been proposed to fit the past optical records for OJ287 and predict the future optical events. If the timing of the optical activity (superflares) in OJ287 can be accurately predicted, then it is possible to calculate or determine the parameters of the binary black hole system (such as the masses of the black holes, the orbital parameters etc.). \item Lighthouse models\\ The periodic optical flares observed in OJ287 are suggested to be caused by the change of the orientation of the realtivistically moving emission regions with respect to the line of sight, resulting in an increase of the Doppler boosting factor (Camenzind \& Krockenberg \cite{Ca92}, Katz \cite{Ka97}, Villata et al. et al. \cite{Vi98}, Qian et al. \cite{Qi01}) and the apparent optical flux density (${\propto}{{\delta}^3}$). Katz (\cite{Ka97}) proposed that in a bianry black hole system the companion exerts a torque on the accretion disk of the primary black hole and results in the precessing of the relativistic jet from the primary sweeping across the line of sight and the periodic variation of the Doppler factor and thus the apparent flux density. Villata et al. (\cite{Vi98}) suggests that the two black holes of the binary system both create relativistic jets which are bent significantly in different directions. In the course of the binary's orbit motion, the directions of the bent parts of the jets from the two black holes rotate with the orbital period, resulting in periodic double-peak flares. In this class of models only the variation of the Doppler boosting factor plays the role, not considering any change in accretion in the central disk-hole systems. \item Alternative models\\ Valtaoja et al. (\cite{Va00}) have argued against both the accretion models and lighthouse models for OJ287. They found that in OJ287 the first optical flare of the double-structure is thermal (non-polarized) and lacks radio counterpart, but the second one is (polarized) synchrotron emission, having a radio counterpart. They proposed an alternative model in which the first optical flare is caused by the disk-crossing during the pericenter passage, having a regular period. And the second one occurs about one year later in the relativistic jet with an association with a radio outburst. For checking this model optical polarization measurements and VLBI studies of the optical-radio association are significant. \end{itemize} In this paper we will discuss the possible existence of a $\sim$13 year periodicity in the radio lightcurve of the optically violently variable (OVV) quasar 3C454.3 at 8GHz and 14.5GHz in which two broad peaks (or bumps) are observed during one period. We propose that the periodicity can be interpreted in the frame of a binary black hole model in which two jets from the two black holes rotate with the period ($\sim$13 year) of the orbital motion. Model-fits to the lightcurves are given. It is shown that both the periodic Doppler boosting effect (lighthouse effect) and the variability of the accretion activity (i.e. the mass-energy inflow into the jets) should be taken into account in order to fully explain the lightcurves. We also discuss some alternative models and future observations for further studying the periodic phenomenon in 3C454.3. | \begin{itemize} \item (1) the supposed model is not unique. In addition to the double-jet models proposed in this paper, models with a single jet are also plausible. But in this case, the two emitting components could be assumed to be situated at different positio ns in the curved jet and their rotation would produce the Doppler boosting profiles. In this paper we do not invoke a specific model for a binary black hole system. \item (2) In our model-fittings, multi-parameters can be chosen. When parameters were chosen we consider mainly three ingredients: (a) the width of the two radio bumps; (2) the difference of Doppler factor at flaring and quiescent epochs should not be too large in order to avoid a huge difference in timescale of varibility at different epochs (Valtaoja et al. \cite{Va00}); \item (3) We should consider both effects due to Doppler boosting and the energy transfer into the jets (or accretion rates). In our models, the Doppler boosting determine the 'profiles' of the possible activity in the radio and optical bands. If the intrinsic (in the rest frame of the flows) strengths are stable, then the apparent activity follows the pattern of the Doppler boosting. This seems correct for the observed radio variations, for example during 1978--1995. \item (4) As for the optical variations, they do not follow the Doppler boosting profiles derived from our models. This could imply we should adopt different parameters for the optical variations, i.e., the optically emitting regions could have orientations different from the radio regions. Moreover, optical flux variations could be largely determined by the variation in accretion rate and their much shorter timescales may be due to mechanisms for acceleartion of electrons and radiative losses. The optical peak at 2005.4 and radio peak at 2006.2, which both deviate the Doppler boosting pattern maximum, further indicate the necessity to consider the effects both from Doppler boosting, accretion and transfer of enery into the jets. \item (5) The radio periodic behaviour still needs to be tested, because the 2005 optical flare has a quite peculiar mm-radio outbursts, which have much narrower profiles than before, very different from the 1981 and 1994 radio flare's profiles. This may imply that the radio periodic behaviour could be a short-term phenomenon. The optical-radio relation is still important to be tested and thus could obtain some information about the combination of the black hole disk system and the Doppler boosting effect. \item (6) Relationship between mm- and cm-outburtsts needs to be further studied in order to study the evolution of the outbursts. \item (7) At present, we cannot predict whether the periodic behaviour in 3C454.3 in radio and optical bands can hold long, and this should be observationally tested during the next period (2007--2020). Both optical and radio monitoring (intensity and polarization) are required. \end{itemize} | 7 | 10 | 0710.1876 |
0710 | 0710.3042_arXiv.txt | {We show that future Ultra-High Energy Cosmic Ray samples should be able to distinguish whether the sources of UHECRs are hosted by galaxy clusters or ordinary galaxies, or whether the sources are uncorrelated with the large-scale structure of the universe. Moreover, this is true independently of arrival direction uncertainty due to magnetic deflection or measurement error. The reason for this is the simple property that the strength of large-scale clustering for extragalactic sources depends on their mass, with more massive objects, such as galaxy clusters, clustering more strongly than lower mass objects, such as ordinary galaxies. } \begin{document} | \label{sec:intro} Identifying the sources of ultrahigh energy cosmic rays (UHECRs, here $E>10^{19}$eV $\equiv 10$ EeV) is complicated by the deflection they presumably experience in Galactic and extragalactic magnetic fields, as well as their relatively poor arrival direction determinations, typically $\sim 1^{\circ}$. Arrival directions of most UHECRs are thus not known well enough to match their positions with specific astrophysical objects. However, there is also useful information in the clustering of UHECRs on large scales, where $\sim$ few degree uncertainties in position become unimportant. The clustering of galaxies in the universe is typically quantified by the two-point correlation function or its analog in Fourier space, the power spectrum. The two-point correlation function $\xi(r)$ of any class of objects (e.g., galaxies of a certain luminosity or color) is defined as the excess number of pairs of such objects at physical separation $r$ over that expected for a random (Poisson) distribution. In Cold Dark Matter models, the large-scale amplitude of $\xi(r)$ (usually referred to as the bias) of a population of objects depends only on their mass, with more massive objects, such as clusters of galaxies, clustering more strongly than less massive objects, such as ordinary galaxies \cite{mo_white_96,sheth_tormen_99,berlind_etal_06b}. The large-scale bias of a UHECR sample is therefore a robust and informative measure of the clustering properties of the source. We cannot measure physical separations for pairs involving UHECRs because they do not have measured redshifts. However, we can measure the angular correlation function $\omega(\theta)$. As is the case for $\xi(r)$, the large-scale amplitude of $\omega(\theta)$ for a UHECR sample depends on the nature of the astrophysical source. However, it also depends on the depth of the sample because deeper samples mix more physically uncorrelated pairs and thus show weaker angular clustering. In order to access the information in the large-scale angular clustering of UHECRs, we must therefore know the depth of our UHECR sample. In this paper, we demonstrate what can be learned from the large-scale angular clustering of UHECRs, we estimate what kind of sample is needed to do this analysis, and we show how to deal with the unknown depth of a UHECR sample, using the GZK effect. | 7 | 10 | 0710.3042 |
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0710 | 0710.4986_arXiv.txt | We calculate axisymmetric toroidal modes of magnetized neutron stars with a solid crust. We assume the interior of the star is threaded by a poloidal magnetic field that is continuous at the surface with the outside dipole field whose strength $B_p$ at the magnetic pole is $B_p\sim10^{16}$G. Since separation of variables is not possible for oscillations of magnetized stars, we employ finite series expansions of the perturbations using spherical harmonic functions to represent the angular dependence of the oscillation modes. For $B_p\sim 10^{16}$G, we find distinct mode sequences, in each of which the oscillation frequency of the toroidal mode slowly increases as the number of radial nodes of the eigenfunction increases. The frequency spectrum of the toroidal modes for $B_p\sim10^{16}$G is largely different from that of the crustal toroidal modes of the non-magnetized model, although the frequency ranges are overlapped each other. This suggests that an interpretation of the observed QPOs based on the magnetic toroidal modes may be possible if the field strength of the star is as strong as $B_p\sim10^{16}$G. | Recent discovery of quasi-periodic oscillations (QPOs) of magnetar candidates is one of the observational manifestations of global oscillations of neutron stars. Israel et al (2005) detected QPOs of frequencies $\sim18$, $\sim 30$ and $\sim$92.5Hz in the tail of the SGR 1806-20 hyperflare observed December 2004, and suggested that the 30Hz and 92.5Hz QPOs could be caused by seismic vibrations of the neutron star crust (see, e.g., Duncan 1998). Later on, in the hyperflare of SGR 1900+14 detected August 1998, Strohmayer \& Watts (2005) found QPOs of frequencies 28, 53.5, 84, and 155 Hz, and claimed that the QPOs could be identified with the low $l$ fundamental toroidal torsional modes of the solid crust of the neutron star. These recent discoveries of QPOs in the giant flares of Soft Gamma-Ray Repeaters SGR 1806-20 (Israel et al 2005, Watts \& Strohmayer 2006, Strohmayer \& Watts 2006) and SGR 1900+14 (Strohmayer \& Watts 2005) have made promising asteroseismology for magnetars, neutron stars with an extremely strong magnetic field (see, e.g., Woods \& Thompson 2006 for a review of SGRs). It is currently common to identify these QPOs with seismic vibrations caused by crustal toroidal modes of the neutron stars, since the frequency range of the modes overlap that of the observed QPOs and from the energetics point of view the crustal toroidal modes would be most easily excited to observable amplitudes by spending a least amount of available energies, for example, those released in magnetic field restructuring (e.g., Duncan 1998). Although the interpretation based on crustal torsional modes looks promising, we need detailed theoretical analyses of oscillations of magnetized neutron stars so that we could get information of physical conditions of the stars through the confrontation between theoretical modelings and observations. This is particularly true for high frequency QPOs (e.g., 625Hz QPO, Watts \& Strohmayer 2006; 1835Hz QPO and less significant QPOs at 720 and 2384 Hz in SGR 1806-20, Strohmayer \& Watts 2006), since there exist classes of modes other than the crustal toroidal modes that can generate the frequencies observed. The presence of a magnetic field makes it possible for toroidal modes to exist in a fluid star even without rotation as does the presence of the shear modulus in the solid crust. In this paper we are interested in axisymmetric toroidal modes since axisymmetirc toroidal and spheroidal modes are decoupled for a poloidal field when the star is non-rotating. For non-axisymmetric modes, the toroidal and spheroidal components are coupled even without rotation and hence the modal analyses of magnetized stars would be much more complicated. Theoretical calculations of toroidal modes of strongly magnetized neutron stars have been carried out by several authors, including Piro (2005), Glampedakis et al (2006), Sotani et al (2006, 2007), and Lee (2007). The analyses by Piro (2005), Glampedakis et al (2006), and Lee (2007) assume Newtonian gravity, while those by Sotani et al (2006, 2007) use general relativistic formulation. Although the studies by Piro (2006) and Lee (2007) ignore the effects of magnetic fields in the fluid core, those by Glampedakis et al (2006) and Sotani et al (2006, 2007) consider magnetic waves propagating in the fluid core, assuming the core is threaded by a magnetic field of substantial strength. Besides the differences mentioned above, most of the authors except Lee (2007) represent the angular dependence of the oscillations by a single spherical harmonic function $Y_l^m(\theta,\phi)$. Since the shear modulus in the crust dominates the magnetic pressure in most parts of the crustal regions for a dipole field of strength $B_p\ltsim 10^{15}$G, this treatment may be justified so long as the crustal toroidal modes are well decoupled from the fluid core. But, if the torsional waves in the crust are strongly coupled with magnetic waves in the core, the treatment may not be justified because the angular dependence of the magnetic waves in the fluid core cannot be correctly represented by a single spherical harmonics. For example, the analysis by Reese, Fincon, \& Rieutord (2004) of toroidal modes in a fluid shell have employed finite series expansions of long length for the perturbations. In this paper, using the method of series expansions of perturbations we calculate toroidal modes of a strongly magnetized neutron star having a fluid core and a solid crust, where the entire interior is assumed to be threaded by a poloidal magnetic field. We employ two different sets of oscillation equations, one for fluid regions and the other for the solid crust, and solutions in the solid and fluid regions are matched at the interfaces between them to obtain an entire solution of a mode. The method of calculation we employ is presented in \S 2, and the numerical results are given in \S 3, and conclusions are in \S 4. | In this paper, we have calculated toroidal modes of magnetized stars with a solid crust, where the entire interior of the star is assumed to be threaded by a poloidal magnetic field that is continuous at the stellar surface to the outside dipole field. We find distinct mode sequences of the toroidal modes, in each of which the mode frequency remains rather constant and only slowly increases as the radial order of the modes increases. In the presence of a solid crust, the frequency separation between the lowest and second lowest frequency mode sequences is approximately given by $\Delta\omega$, but that between the second and third lowest frequency mode sequences by $\Delta\omega/2$, where the frequency separation $\Delta\omega$ is roughly proportional to the field strength $B_p$. This frequency pattern of the low frequency sequences is different from that found for the model without the crust, for which the frequency separation between the sequential sequences of low frequency modes is given by $\Delta\omega/2$. We also find that for the equation of state we use, $\Delta\omega$ is larger for smaller $M$. The eigenfunction $\xi_\phi$ of the modes belonging to the lowest frequency sequence have much larger amplitudes than $B^\prime_\phi$ and can penetrate into the solid crust. On the other hand, $\xi_\phi$ of the modes belonging to the higher frequency sequences is much smaller than $B^\prime_\phi$, which is well confined into the fluid core and does not have any substantial amplitudes in the solid crust. The frequency ranges of the toroidal modes we find for the magnetized neutron star with $B_p\sim10^{16}$G overlap the QPO frequencies found for the magnetar candidates, SGR 1806-20 and SGR 1900+14. This suggests that we may interpret the observed QPOs based on the magnetic toroidal modes, and that detailed comparisons between observed frequency spectra and theoretical calculations make it relevant to infer physical parameters of the magnetar candidates, such as the equation of state and the strength of the magnetic field. But, we think it worth pointing out that except for the modes belonging to the lowest frequency sequence, the magnetic perturbations, which have much larger amplitude than $\xi_\phi$, are well confined in the fluid core and do not have substantial amplitudes in the solid crust, which make it difficult for the modes to be directly observable. If the magnetar candidates do no have a crust, the problem of observability could be avoided. In this case, however, we have to use more realistic surface boundary conditions than $\pmb{B}^\prime=0$ used in the present calculations. Note that, although the existence of a solid crust affects the frequency pattern of the low frequency mode sequences, the frequency range of the magnetic modes itself is not very much dependent on the presence or absence of a solid crust unless the crust is extremely thick. We have tried to find toroidal modes well confined in the solid crust or core toroidal modes that are in resonance with the crustal toroidal modes, but failed. We also find it extremely difficult to identify distinct toroidal mode sequences for magnetic fields weaker than $B_p\ltsim10^{15}$G when a solid crust is included in the models. It is therefore not clear whether the toroidal mode sequences of the kind we find in the present paper can survive also for weakly magnetized neutron stars with a solid crust. If the field strength is much weaker than $B_p\sim10^{15}$G, the magnetic fields may have only minor effects on the crustal toroidal modes (e.g., Lee 2007), and it will be justified to use frequency spectra of the crust modes theoretically obtained in the weak field limit to interpret observed QPOs. As briefly noted in the last section, it is not theoretically clear how the frequency spectra of the toroidal modes of magnetized stars should look like, and how the toroidal modes behave in the limit of $n\ge k\rightarrow\infty$. We may even speculate that the frequency spectra we obtained reflect the existence of continuous spectra (see, e.g., Goedbloed \& Poedts 2004, see also Levin 2007), but the detailed numerical analysis of continuous frequency spectra is beyond the scope of this paper. It will be worthwhile to examine the effects of an interior toroidal field on magnetic modes. It is also needed to extend the present analysis to a general relativistic formulation (e.g., Sotani et al 2006, 2007). | 7 | 10 | 0710.4986 |
0710 | 0710.1271_arXiv.txt | We present results of 3-neutrino flavor evolution simulations for the neutronization burst from an O-Ne-Mg core-collapse supernova. We find that nonlinear neutrino self-coupling engineers a single spectral feature of stepwise conversion in the inverted neutrino mass hierarchy case and in the normal mass hierarchy case, a superposition of two such features corresponding to the vacuum neutrino mass-squared differences associated with solar and atmospheric neutrino oscillations. These neutrino spectral features offer a unique potential probe of the conditions in the supernova environment and may allow us to distinguish between O-Ne-Mg and Fe core-collapse supernovae. | 7 | 10 | 0710.1271 |
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0710 | 0710.3104_arXiv.txt | The Milky Way is the only galaxy for which we can resolve individual stars at all evolutionary phases, from the Galactic center to the outskirt. The last decade, thanks to the advent of near IR detectors and 8 meter class telescopes, has seen a great progress in the understanding of the Milky Way central region: the bulge. Here we review the most recent results regarding the bulge structure, age, kinematics and chemical composition. These results have profound implications for the formation and evolution of the Milky Way and of galaxies in general. This paper provides a summary on our current understanding of the Milky Way bulge, intended mainly for workers on other fields. | How did the Milky Way form? Decades ago, the Galactic formation scenarios focused on the disk and the halo populations, because these were the components that astronomers knew something about. For example, in the classic works of Eggen et al. (1962) or Searle \& Zinn (1978) the bulge was not mentioned. All-sky optical maps (Fig.~\ref{maps}) did not show a clear, separate component in the inner Milky Way. Yet upon looking at the new DIRBE or 2MASS near-infrared maps of the whole sky, it is evident that the Milky Way is a spiral galaxy with a peanut-shape bulge. The simultaneous observation that bulge stars were mainly old made it clear that to answer this question the attention had to shift towards what seems to be the first massive component to be formed in the Milky Way. This is a Copernican-like revolution on Galactic scales, and it is happening now! We see our Galaxy as if it were an external galaxy for the first time, and are fortunate to have such new perspective, and also to be able to provide specific answers to the many questions regarding its formation. \begin{figure} \includegraphics[height=5.5in,width=5.3in,angle=00]{minniti_fig1.ps} \caption{ Top: Near-IR COBE-DIRBE map of the whole sky (Dwek et al. 1995). Bottom: Optical map for comparison (Copyright Axel Mallenhoff 2001). These sky maps illustrate why astronomers did not realize before the importance of the bulge.} \label{maps} \end{figure} Note that we have only one object to study: our Galaxy. This is a fundamental limitation because we have to get the answers right. The advantage in the case of our Galaxy, of course, is that we can study the subject in detail: in no other galaxies the fundamental problems of Astronomy can be surpassed. Basically, these fundamental problems are: (1) The distance problem: we can resolve the components of the Milky Way into stars, and study them in situ, obtaining 3-D positions and motions, and detailed chemical compositions of individual stars; and (2) The timescale problem: we see an instant snapshot, and must infer histories from that. In our galaxy we can separate and date the components, using the main-sequence turn-offs of different stellar populations and clusters. In spite of a small community working on the Galactic bulge, a revolution has occurred in this field in the last decade, with great progress in comparison with other fields of Astronomy. For the first time, we have the answers to important questions such as: {\bf 1.} Is the bulge a different Galactic component? {\bf 2.} When did the Galactic bulge form? {\bf 3.} How did it form? {\bf 4.} Is there a radial gradient in the bulge? {\bf 5.} Are there globular clusters associated with the bulge? {\bf 6.} Are there planets in the bulge? The 8m class telescopes were built last decade, aiming to answer these questions. The proposed options were endless (bulge formation from secular evolution of the disk, from the halo, in a single burst, in several episodes, by slow accretion of smaller sub-units, etc). The evidence and the answers described below serve as a basis for understanding more distant galaxies which cannot be studied (dissected) in similar detail. | For the first time, we have the answers to the following basic questions: {\bf 1.} Is the bulge a different component? Yes, based on all the evidence available, the bulge is a distinct Galactic component, with different kinematics and compositions from the thin disk, the thick disk and the halo. {\bf 2.} How did the Galactic bulge form? It formed on a short timescale ($\sim 1$ Gyr), as demonstrated by the $\alpha$-element enrichment. Despite the presence of the bar, models like a bulge formation via secular evolution of the disk can be firmly excluded. {\bf 3.} What is the age of the bulge? The bulk of the stellar population is $\sim 10$ Gyr old. However, there are traces of a small fractions of intermediate-age stars, and of metal-poor stars. The latter might well be the oldest population in the Galaxy. {\bf 4.} Is there a gradient in the bulge? Yes, there is a stellar population gradient as shown for example by the CMDs. Now it is found that this gradient is mostly due to metallicity, which decreases along the Galactic minor axis. {\bf 5.} Are there globular clusters associated with the bulge? Yes, there is a population of metal-rich globular clusters in the central regions that share the kinematics, spatial distribution, and composition of the bulge field stars. {\bf 6.} Are there planets in the bulge? Even though this question seems to belong to another field, another of the recent advances was the discovery of planets in the bulge by HST. The available data suggests that giant planets are as numerous in the bulge as they are in the Solar neighborhood. These revolutionary advances that impact the whole of extragalactic Astronomy cannot be attributed to the success of a single group, but to the combined contributions of different teams. Where controversy was present before, today similar answers are given. Progress has occurred! | 7 | 10 | 0710.3104 |
0710 | 0710.3618_arXiv.txt | In the course of carrying out a wide-field CCD imaging survey, two new methods for correlating the images to star catalogues have been developed, motivated by the need to efficiently handle the large number of stellar sources present on the images. Most previously published algorithms successfully cater for small lists ($\le$ 50 stars), but do not scale well to wide-fields containing $10^3$ or more stellar sources. The problem of matching coordinate lists of point sources is a necessary prerequisite for deriving an astrometric plate solution. The objective is to match a subset of stars found on an image to their corresponding entries in a stellar catalogue in order to determine the transformation between detector coordinates and sky coordinates. The algorithm must handle translation and rotation, and small changes in scale caused by temperature related changes in focal length. In addition, it must cope with additional and missing stars. That is, the two lists may only partially overlap. The efficiency of the algorithm is of paramount concern, since it is embodied within the closed-loop pointing system of the telescope and therefore affects the duty-cycle time, and ultimately constrains the number of images that can be acquired each night. Surveys that require very high photometric precision typically seek to accurately align their fields on the same detector pixels each night to overcome residual flat-fielding errors \citep{Everett}, and would benefit from the efficiency gains of a fast matching algorithm. Similarly, high cadence surveys, such as the Southern Sky Survey \citep{Keller} could improve precision and reduce its duty-cycle by utilizing a fast closed-loop pointing algorithm. Moreover, real-time attitude adjustments on spacecraft might be possible with the aid of an efficient matching algorithm to analyze on-board star camera images \citep[see for example][]{Fraser}. A number of algorithms have been proposed to solve this problem. \citet{Groth} describes an algorithm that matches geometrically similar shapes (triangles) in the two lists. By limiting the number of triangles constructed, and by only matching those triangles whose ratio of longest to shortest side are within a defined limit, his matching phase has a computational complexity of $O(n^{4.5})$ where $n$ is the number of stars in each list. \cite{Stetson} describes a very similar algorithm that he developed independently at around the same time. \citet{Murtagh} reviews a number of approaches and proposes his own, based upon characterization of a set of coordinates couples, with matching based on the proximity of feature vectors in the two lists. His method's matching phase has a computational complexity of $O(n^2)$. Nevertheless, Groth's algorithm appears to be the most widely accepted, with the methods applied across disciplines. For example, \citet{Arz} discuss its application to the problem of computer-aided identification of whale sharks, while \citet{Mars}, building upon the work of \citet{Groth}, describe an optimization to the voting phase of the algorithm, concluding that their method reduces the need for complicated filtering methods while successfully reducing the number of false matches. More recently, \citet{PB} describe another variation of triangle matching, optimized to handle large lists of objects extracted from wide field images. Large fields contain thousands of stars and pose a severe test for matching algorithms, requiring efficient methods to accommodate the large number of point sources. The following sections discuss two new methods for pattern matching that have a matching phase with a complexity that is nearly $O(1)$, at the cost of a slight loss in generality. They are collectively referred to as \emph{Optimistic Pattern Matching (OPM)} because they assume that (i) a good match is likely to be found, and (ii) the scale of the image is approximately known, thus permitting the use of an early exit strategy whereby only a small percentage of the candidate list is examined. By contrast, previous methods assumed an unknown scale which required the entire candidate list to be processed to determine the most likely match using a statistical approach. This required additional phases and complexity. In practice, an \emph{a priori} knowledge of an instrument's focal length is common place, and the use of a more general algorithm that assumes it is unknown mandates strategies that unnecessarily degrade performance. Section 2 describes the algorithms in detail. $OPM_A$ is based upon a new definition of triangle space, while $OPM_B$ uses an alternative shape characterization method. Section 3 tests their performance using a large sample of survey images and compares them to earlier methods. Conclusions are summarized in Section 4. | Two new techniques for matching two-dimensional coordinate lists in nearly constant time have been presented. The matching phase of $OPM_B$ is nearly $O(1)$, being independent of list size. These algorithms have a significant performance advantage over previous techniques, at a slight loss in generality, caused by the requirement that the approximate focal length of the optical system is known \emph{a priori}. This requirement permits the determination of the image scale from the physical dimensions of the detector, allowing $OPM$ algorithms to directly compare a subset of triangles (or shapes) to their counterparts derived from a reference catalogue, without having to process the entire set, as is the case when the scale is unknown. By employing early exit strategies, postponing work until absolutely necessary, testing candidates in the order most likely to yield success, and combining these with and an efficient mechanism for rejecting false positives, a highly efficient search, in nearly constant time is possible. Small uncertainties in the focal length, such as caused by temperature related changes, are accommodated by selecting an appropriate matching tolerance. The actual focal-length is determined and reported as part of the astrometric solution. The $OPM$ algorithms are particularly suited to processing large lists or in situations where pattern matching must be performed as quickly as possible. The performance of these algorithms makes it practical to search thousands of fields very quickly, if for example, the coordinates of the field center were unknown. Similarly, when only an approximate focal-length is known, it is perfectly reasonable to attempt to iteratively match the field using a range of focal-lengths. | 7 | 10 | 0710.3618 |
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0710 | 0710.2102_arXiv.txt | We use hydrodynamical simulations of disk galaxies to study relations between star formation and properties of the molecular interstellar medium (ISM). We implement a model for the ISM that includes low-temperature ($T<10^{4}$~K) cooling, directly ties the star formation rate to the molecular gas density, and accounts for the destruction of $\Hmol$ by an interstellar radiation field from young stars. We demonstrate that the ISM and star formation model simultaneously produces a spatially-resolved molecular-gas surface density Schmidt-Kennicutt relation of the form $\SigmaSFR \propto \SigmaHmol^{\nmol}$ with $\nmol\approx1.4$ independent of galaxy mass, and a total gas surface density -- star formation rate relation $\SigmaSFR \propto \Sigmagas^{\ntot}$ with a power-law index that steepens from $\ntot\sim2$ for large galaxies to $\ntot\gtrsim4$ for small dwarf galaxies. We show that deviations from the disk-averaged $\SigmaSFR \propto \Sigmagas^{1.4}$ correlation determined by \cite{kennicutt1998a} owe primarily to spatial trends in the molecular fraction $\fHmol$ and may explain observed deviations from the global Schmidt-Kennicutt relation. In our model, such deviations occur in regions of the ISM where the fraction of gas mass in molecular form is declining or significantly less than unity. Long gas consumption time scales in low-mass and low surface brightness galaxies may owe to their small fractions of molecular gas rather than mediation by strong supernovae-driven winds. Our simulations also reproduce the observed relations between ISM pressure and molecular fraction and between star formation rate, gas surface density, and disk angular frequency. We show that the Toomre criterion that accounts for both gas and stellar densities correctly predicts the onset of star formation in our simulated disks. We examine the density and temperature distributions of the ISM in simulated galaxies and show that the density probability distribution function (PDF) generally exhibits a complicated structure with multiple peaks corresponding to different temperature phases of the gas. The overall density PDF can be well-modeled as a sum of lognormal PDFs corresponding to individual, approximately isothermal phases. We also present a simple method to mitigate numerical Jeans fragmentation of dense, cold gas in Smoothed Particle Hydrodynamics codes through the adoption of a density-dependent pressure floor. | \label{section:introduction} Galaxy formation presents some of the most important and challenging problems in modern astrophysics. A basic paradigm for the dissipational formation of galaxies from primordial fluctuations in the density field has been developed \citep[e.g.][]{white1978a,blumenthal1984a,white1991a}, but many of the processes accompanying galaxy formation are still poorly understood. In particular, star formation shapes the observable properties of galaxies but involves a variety of complicated dynamical, thermal, radiative, and chemical processes on a wide range of scales \citep[see][for a review]{mckee2007a}. Observed galaxies exhibit large-scale correlations between their global star formation rate (SFR) surface density $\SigmaSFR$ and average gas surface density $\Sigmagas$ \citep{kennicutt1989a,kennicutt1998a}, and these global correlations serve as the basis for treatments of star formation in many models of galaxy formation. While such models have supplied important insights, detailed observations of galaxies have recently provided evidence that the molecular, rather than the total, gas surface density is the primary driver of global star formation in galaxies \citep[e.g.,][]{wong2002a,boissier2003a,heyer2004a,boissier2007a,calzetti2007a,kennicutt2007a}. In this study, we adopt an approach in which empirical and theoretical knowledge of the star formation efficiency (SFE) in dense, molecular gas is used as the basis for a star formation model in hydrodynamical simulations of disk galaxy evolution. This approach requires modeling processes that shape properties of the dense phase of the interstellar medium (ISM) in galaxies. The purpose of this paper is to present such a model. Stellar populations in galaxies exhibit salient trends of colors and metallicities with galaxy luminosity \citep[e.g.,][]{kauffmann2003c,blanton2005a,cooper2007a}. In the hierarchical structure formation scenario these trends should emerge through the processes of star formation and/or stellar feedback in the progenitors of present-day galaxies. Observationally, ample evidence suggests that the efficiency of the conversion of gas into stars depends strongly and non-monotonically on mass of the system. For example, the faint-end of the galaxy luminosity function has a shallow slope \citep[$\alphaL\approx1.0-1.3$, e.g.,][]{blanton2001a,blanton2003a} compared to the steeper mass function of dark matter halos \citep[$\alphaDM\approx2$, e.g.,][]{press1974a,sheth1999a}, indicating a decrease in SFE in low-mass galaxies. At the same time, the neutral hydrogen (HI) and baryonic mass functions may be steeper than the luminosity function \citep[$\alphaHI\approx1.3-1.5$, e.g.,][]{rosenberg2002a,zwaan2003a}. The baryonic \cite{tully1977a} relation is continuous down to extremely low-mass dwarf galaxies \citep[e.g.,][]{mcgaugh2005a,geha2006a}, indicating that the fractional baryonic content of galaxies of different mass is similar. Hence, low-mass galaxies that are unaffected by environmental processes are gas-rich, yet often form stars inefficiently. While feedback processes from supernovae and AGN \citep[e.g.,][]{brooks2007a,sijacki2007a}, or the efficiency of gas cooling and accretion \citep{dekel2006a,dekel2008a}, may account for part of these trends, the SFE as a function of galaxy mass may also owe to intrinsic ISM processes \citep[e.g.,][]{tassis2008a,kaufmann2007a}. To adequately explore the latter possibility, a realistic model for the conversion of gas into stars in galaxies is needed. Traditionally, star formation in numerical simulations of galaxy formation is based on the empirical Schmidt-Kennicutt (SK) relation \citep{schmidt1959a,kennicutt1989a,kennicutt1998a}, in which star formation rate is a {\it universal} power-law function of the total disk-averaged or global gas surface density: $\SigmaSFR\propto \Sigmagas^{\ntot}$ with $\ntot\approx 1.4$ describing the correlation for the entire population of normal and starburst galaxies. However, growing observational evidence indicates that this relation may not be universal on smaller scales within galaxies, especially at low surface densities. Estimates of the slope of the SK relation within individual galaxies exhibits significant variations. For example, while \citet{schuster2007a} and \citet{kennicutt2007a} find $\ntot\approx 1.4$ for the molecular-rich galaxy M51a, similar estimates in other large, nearby galaxies \citep[including the Milky Way (MW);][]{misiriotis2006a} range from $\ntot\approx1.2$ to $\ntot\approx 3.5$ \citep{wong2002a,boissier2003a} depending on dust-corrections and fitting methods. The disk-averaged total gas SK relation for normal (non-starburst) galaxies also has a comparably steep slope of $\ntot\approx2.4$, with significant scatter \citep{kennicutt1998a}. While the variations in the $\SigmaSFR-\Sigmagas$ correlation may indicate systematic uncertainties in observational measurements, intrinsic variations or trends in galaxy properties may also induce differences between the global relation determined by \cite{kennicutt1998a} and the $\SigmaSFR-\Sigmagas$ correlation in individual galaxies. Galaxies with low gas surface densities, like dwarfs or bulgeless spirals, display an even wider variation in their star formation relations. \citet{heyer2004a} and \citet{boissier2003a} show that in low-mass galaxies the SFR dependence on the total gas surface density exhibits a power-law slope $\ntot\approx 2-3$ that is considerably steeper than the global \citet{kennicutt1998a} relation slope of $\ntot\approx 1.4$. Further, star formation in the low-surface density outskirts of galaxies also may not be universal. Average SFRs appear to drop rapidly at gas surface densities of $\Sigmagas\lesssim 5-10{\rm\, \Msun\,pc^{-2}}$ \citep{hunter1998a,martin2001a}, indicating that star formation may be truncated or exhibit a steep dependence on the gas surface density. The existence of such threshold surface densities have been proposed on theoretical grounds \citep{kennicutt1998a,schaye2004a}, although recent GALEX results using a UV indicator of star formation suggest that star formation may continue at even lower surface densities \citep{boissier2007a}. Star formation rates probed by damped Lyman alpha absorption (DLA) systems also appear to lie below the \cite{kennicutt1998a} relation, by an order of magnitude \citep{wolfe2006a,wild2007a}, which may indicate that the relation between SFR and gas surface density in DLA systems differs from the local relation measured at high $\Sigmagas$. In contrast, observations generally show that star formation in galaxies correlates strongly with {\it molecular\/} gas, especially with the highly dense gas traced by HCN emission \citep{gao2004a,wu2005a}. The power-law index of the SK relation connecting the SFR to the surface density of molecular hydrogen consistently displays a value of $\nmol\approx1.4$ and exhibits considerably less galaxy-to-galaxy variation \citep{wong2002a,murgia2002a,boissier2003a,heyer2004a,matthews2005a,leroy2005,leroy2006,gardan2007a}. Molecular gas, in turn, is expected to form in the high-pressure regions of the ISM \citep{elmegreen1993a,elmegreen1994a}, as indicated by observations \citep{blitz2004a,blitz2006a,gardan2007a}. Analytical models and numerical simulations that tie star formation to the fraction of gas in the dense ISM are successful in reproducing many observational trends \citep[e.g.,][]{elmegreen2002a,kravtsov2003a,krumholz2005a,li2005a,li2006a,tasker2006a,tasker2007a,krumholz2007a,wada2007a,tassis2007a}. Recently, several studies have explored star formation recipes based on molecular hydrogen. \citet{pelupessy2006a} and \citet{booth2007a} implemented models for $\Hmol$ formation in gaseous disks and used them to study the molecular content and star formation in galaxies. However, these studies focused on the evolution of galaxies of a single mass and did not address the origin of the SK relation, its dependence on galaxy mass or structure, or its connection to trends in the local molecular fraction. Our study examines the SK relation critically, including its dependence on the structural and ISM properties of galaxies of different masses, to explain the observed deviations from the global SK relation, to investigate other connections between star formation and disk galaxy properties such as rotation or gravitational instability, and to explore how the temperature and density structure of the ISM pertains to the star formation attributes of galaxies. To these ends, we develop a model for the ISM and star formation whose key premise is that star formation on the scales of molecular clouds ($\sim 10$~pc) is a function of molecular hydrogen density with a universal SFE per free-fall time \citep[e.g.,][]{krumholz2007b}. Molecular hydrogen, which we assume to be a proxy for dense, star forming gas, is accounted for by calculating the local $\Hmol$ fraction of gas as a function of density, temperature, and metallicity using the photoionization code Cloudy \citep{ferland1998a} to incorporate $\Hmol$-destruction by the UV radiation of local young stellar populations. We devise a numerical implementation of the star formation and ISM model, and perform hydrodynamical simulations to study the role of molecular gas in shaping global star formation relations of self-consistent galaxy models over a representative mass range. The results of our study show that many of the observed global star formation correlations and trends can be understood in terms of the dependence of molecular hydrogen abundance on the local gas volume density. We show that the physics controlling the abundance of molecular hydrogen and its destruction by the interstellar radiation field (ISRF) play a key role in shaping these correlations, in agreement with earlier calculations based on more idealized models of the ISM \citep{elmegreen1993a,elmegreen1994a}. While our simulations focus on the connection between the molecular ISM phase and star formation on galactic scales, the formation of molecular hydrogen has also been recently studied in simulations of the ISM on smaller, subgalactic scales \citep{glover2007a,dobbs2007a}. These simulations are complementary to the calculations presented in our study and could be used as input to improve the molecular ISM model we present. The paper is organized as follows. The simulation methodology, including our numerical models for the ISM, interstellar radiation field, and simulated galaxies, is presented in \S \ref{section:methodology}. The results of the simulations are presented in \S \ref{section:results}, where the simulated star formation relations in galactic disks and correlations of the molecular fraction with the structure of the ISM are examined. We discuss our results in \S \ref{section:discussion} and conclude with a summary in \S \ref{section:summary}. Details of our tests of numerical fragmentation in disk simulations and calculations of the model scaling between star formation, gas density, and orbital frequency are presented in the Appendices. Throughout, we work in the context of a dark-energy dominated cold dark matter cosmology with a Hubble constant $H_{0} \approx 70\kms\Mpc^{-1}$. \begin{figure*} \figurenum{1} \epsscale{1} \plotone{fig1.eps} \caption{\label{fig:cooling} Cooling ($\Lambda$) and heating ($\Heating$) rates for interstellar and intergalactic gas as a function of gas density ($\nH$), temperature ($T$), metallicity ($Z$), and interstellar radiation field (ISRF) strength ($\Uisrf$), in units of the ISRF strength in the MW at the solar circle, as calculated by the code Cloudy \citep{ferland1998a}. Shown are the cooling (dotted lines), heating (dashed lines) and net cooling (solid) functions over the temperature range $T=10^{2}-10^{9}~\K$. Dense gas can efficiently cool via atomic and molecular coolants below $T=10^{4}\K$, depending on the gas density and the strength of the ISRF. A strong ISRF can enable the destruction of $\Hmol$ gas and thereby reduce the SFR. } \end{figure*} | \label{section:discussion} Results presented in the previous sections indicate that star formation prescriptions based on the local abundance of molecular hydrogen lead to interesting features of the global star formation relations. We show that the inclusion of an interstellar radiation field is critical to control the amount of diffuse $\Hmol$ at low gas densities. For instance, without a dissociating ISRF the low mass dwarf galaxy eventually becomes almost fully molecular, in stark contrast with observations. We also show that without the dissociating effect of the ISRF our model galaxies produce a much flatter relation between molecular fraction $\fHmol$ and pressure $\Pext$, as can be expected from the results of \cite{elmegreen1993a}. Including $\Hmol$-destruction by an ISRF results in a $\fHmol-\Pext$ relation in excellent agreement with the observations of \citet{wong2002a} and \citet{blitz2004a,blitz2006a}. Our model also predicts that the relation between $\Sigma_{\rm SFR}$ and $\Sigma_{\rm gas}$ should not be universal and can be considerably steeper than the canonical value of $n_{\rm tot}=1.4$, even if the three-dimensional Schmidt relation in molecular clouds is universal. The slope of the relation is controlled by the dependence of molecular fraction (i.e., fraction of star forming gas) on the total local gas surface density. This relation is non-trivial because the molecular fraction is controlled by pressure and ISRF strength, and can thus vary between different regions with the same total gas surface density. The relation can also be different in the regions where the disk scale-height changes rapidly (e.g., in flaring outer regions of disks), as can be seen from equation~\ref{equation:structural_schmidt_law} \citep[see also][]{schaye2007a}. We show that the effect of radial variations in the molecular fraction $\fHmol$ and gas scale heights ($\hSFR$ and $\hgas$) on the SFR can be accounted for in terms of a structural, SK-like correlation, $\SigmaSFR\propto\fHmol \hSFR \hgas^{-1.5} \Sigmagas^{1.5}$, that trivially relates the local SFR to the consumption of molecular gas with an efficiency that scales with the local dynamical time. A generic testable prediction of our model is that deviations from the SK $\SigmaSFR-\Sigmagas$ relation are expected in galaxies or regions of galaxies where the molecular fraction is declining or much below unity. As we discussed above, star formation in the molecule-poor ($\fHmol\sim0.1$) galaxy M33 supports this view as its total gas SK power-law index is $\ntot\approx 3.3$ \citep{heyer2004a}. While our simulations of an M33 analogue produce a steep SK power-law, the overall SFE in our model galaxies lies below that observed for M33. However, as recently emphasized by \cite{gardan2007a}, the highest surface density regions of M33 have unusually efficient star formation compared with the normalization of the \cite{kennicutt1998a} relation and so the discrepancy is not surprising. Given that dwarf galaxies generally have low surface densities and are poor in molecular gas, it will be interesting to examine SK relation in other small-mass galaxies. Another example of a low-molecular fraction galaxy close to home is M31, which has only a fraction $\fHmol\approx0.07$ of its gas in molecular form within 18 kpc of the galactic center \citep{nieten2006a}. Our model would predict that this galaxy should deviate from the total gas SK relation found for molecular-rich galaxies. Observationally, the SK relation of M31 measured by \citet{boissier2003a} is rather complicated and even has an increasing SFR with decreasing gas density over parts of the galaxy. Low molecular fractions can also be expected in the outskirts of normal galaxies and in the disks of low surface brightness galaxies. The latter have molecular fractions of only $\fHmol\lesssim 0.10$ \citep{matthews2005a,das2006a}, and we therefore predict that they will not follow the total gas SK relation obeyed by molecule-rich, higher surface density galaxies. At the same time, LSBs do lie on the same relation between $\Hmol$ mass and far infrared luminosity as higher surface brightness (HSB) galaxies \citep{matthews2005a}, which suggests that the dependence of star formation on molecular gas may be the same in both types of galaxies. An alternative formulation of the global star formation relation is based on the angular frequency of disk rotation: $\Sigma_{\rm SFR}\propto \Sigmagas\Omega$. That this relation works in real galaxies is not trivial, because star formation and dynamical time-scale depend on the local gas density, while $\Omega$ depends on the total mass distribution {\it within} a given radius. Although several models were proposed to explain such a correlation \citep[see, e.g.,][for reviews]{kennicutt1998a,elmegreen2002a}, we show in \S \ref{section:results:sfr_rotation} and the second Appendix that the star formation correlation with $\Omega$ can be understood as a fortuitous correlation of $\Omega$ with gas density of $\Omega\propto \rho_{\rm gas}^{\alpha}$, where $\alpha\approx 0.5$, for self-gravitating exponential disks or exponential disks embedded in realistic halo potentials. Moreover, we find that $\SigmaSFR\propto \Sigmagas\Omega$ breaks down at low values of $\Sigmagas\Omega$ where the molecular fraction declines, similarly to the steepening of the SK relation. Our models therefore predict that the $\SigmaSFR\propto\SigmaHmol\Omega$ relation is more robust than the $\SigmaSFR\propto\Sigmagas\Omega$ relation. An important issue related to the global star formation in galaxies is the possible existence of star formation thresholds \citep{kennicutt1998a,martin2001a}. Such thresholds are expected to exist on theoretical grounds, because the formation of dense, star-forming gas is thought to be facilitated by either dynamical instabilities \citep[see, e.g.,][for comprehensive reviews]{elmegreen2002a,mckee2007a} or gravithermal instabilities \citep{schaye2004a}. We find that the two-component Toomre instability threshold that accounts for both stars and gas, $Q_{\rm sg}<1$, works well in predicting the transition from atomic gas, inert to star formation, to the regions where molecular gas and star formation occur in our simulations. Our results are in general agreement with \citet{li2005a,li2005b,li2006a}, who used sink-particle simulations of the dense ISM in isolated galaxies to study the relation between star formation and the development of gravitational instabilities, and with observations of star formation in the LMC \citep{yang2007a}. Our simulations further demonstrate the importance of accounting for all mass components in the disk to predict correctly which regions galactic disks are gravitationally unstable. Given that our star formation prescription is based on molecular hydrogen, the fact that $Q_{\rm sg}$ is a good threshold indicator may imply that gravitational instabilities strongly influence the abundance of dense, molecular gas in the disk. Conversely, the gas at radii where the disk is stable remains at low density and has a low molecular fraction. We find that in our model galaxies, the shear instability criterion of \cite{elmegreen1993a} does not work as well as the Toomre $Q_{\rm sg}$-based criterion. Almost all of the star formation in our model galaxies occurs at surface densities $\Sigmagas\gtrsim 3{\rm\ M_{\odot}\,pc^{-2}}$, which is formally consistent with the \citet{schaye2004a} constant surface density criterion for gravithermal instability. However, as Figure~\ref{fig:kennicutt.gas} shows, we do not see a clear indication of threshold at a particular surface density and our GD-SF models that have an effectively isothermal ISM with $T\approx10^{4}\K$ (and hence do not have a gravithermal instability) still show a good correlation between regions where $\Qsg<1$ and regions where star formation operates. Our results have several interesting implications for interpretation of galaxy observations at different epochs. First, low molecular fractions in dwarf galaxies mean that only a small fraction of gas is participating in star formation at any given time. This connection between SFE and molecular hydrogen abundance may explain why dwarf galaxies are still gas rich today compared to larger mass galaxies \citep{geha2006a}, without relying on mediation of star formation or gas blowout by supernovae. Note that a similar reasoning may also explain why large LSBs at low redshifts are gas rich but anemic in their star formation. Understanding the star formation and evolution of dwarf galaxies is critical because they serve as the building blocks of larger galaxies at high redshifts. Such small-mass galaxies are also expected to be the first objects to form large masses of stars and should therefore play an important role in enrichment of primordial gas and the cosmic star formation rate at high redshifts \citep{hopkins2004a,hopkins2006am}. The star-forming disks at $z\sim2$ that may be progenitors of low-redshift spiral galaxies are observed to lack centrally-concentrated bulge components \citep{elmegreen2007a}. Given that galaxies are expected to undergo frequent mergers at $z>2$, bulges should have formed if a significant fraction of baryons are converted into stars during such mergers \citep[e.g.,][]{gnedin2000a}. The absence of the bulge may indicate that star formation in the gas rich progenitors of these $z\sim 2$ systems was too slow to convert a significant fraction of gas into stars. This low SFE can be understood if the high cosmic UV background, low-metallicities, and low dust content of high-$z$ gas disks keep their molecular fractions low \citep{pelupessy2006a}, thereby inhibiting star formation over most of gas mass and keeping the progenitors of the star-forming $z\sim 2$ disks mostly gaseous. Gas-rich progenitors may also help explain the prevalence of extended disks in low-redshift galaxies despite the violent early merger histories characteristic of $\Lambda$CDM universes, as gas-rich mergers can help build high-angular momentum disk galaxies \citep[][]{robertson2006c}. Mergers of mostly stellar disks, on the other hand, would form spheroidal systems. Our results may also provide insight into the interpretation of the results of \cite{wolfe2006a}, who find that the SFR associated with neutral atomic gas in damped Lyman alpha (DLA) systems is an order of magnitude lower than predicted by the local \citet{kennicutt1998a} relation. The DLAs in their study sample regions with column densities $N_{\mathrm{H}}\approx 2\times 10^{21}\rm\ cm^{-2}$, or surface gas densities of $\Sigma_{\rm DLA}\approx 20\rm\ M_{\odot}\,pc^{-2}$, assuming a gas disk with a thickness of $h\approx100$~pc. Suppose the local SK relation steepens from the local relation with the slope $n_0=1.4$ to a steeper slope $n_1$ below some surface density $\Sigma_{\rm b}>\Sigma_{\rm DLA}$. Then for $\Sigmagas<\Sigma_{\rm b}$, the SFR density will be lower than predicted by the local relation by a factor of $(\Sigmagas/\Sigma_{\rm b})^{n_1-n_0}$. For $n_0=1.4$ and $n_1=3$ the SFR will be suppressed by a factor of $>10$ for $\Sigmagas/\Sigma_{\rm b}<0.25$. Thus, the results \cite{wolfe2006a} can be explained if the total gas SK relation at $z\sim3$ steepens below $\Sigmagas\lesssim 100\rm\ M_{\odot}\, pc^{-2}$. We suggest that if the majority of the molecular hydrogen at these redshifts resides in rare, compact, and dense systems \citep[e.g.,][]{zwaan2006a}, then both the lack of star formation and the rarity of molecular hydrogen in damped Ly$\alpha$ absorbers may be explained simultaneously. Our results also indicate that the thermodynamics of the ISM can leave an important imprint on its density probability distribution. Each thermal phase in our model galaxies has its own log-normal density distribution. Our results thus imply that using a single lognormal PDF to build a model of global star formation in galaxies \citep[e.g.,][]{elmegreen2002a,wada2007a} is likely an oversimplification. Instead, the global star formation relation may vary depending on the dynamical and thermodynamical properties of the ISM. We can thus expect differences in the SFE between the low-density and low-metallicity environments of dwarf and high-redshift galaxies and the higher-metallicity, denser gas of many large nearby spirals. Note that many of the results and effects we discuss above may not be reproduced with a simple 3D density threshold for star formation, as commonly implemented in galaxy formation simulations. Such a threshold can reproduce the atomic-to-molecular transition only crudely and would not include effects of the local interstellar radiation field, metallicity and dust content, etc. A clear caveat for our work is that the simulation resolution limits the densities we can model correctly. At high densities, the gas in our simulations is over-pressurized to avoid numerical Jeans instability. The equilibrium density and temperature structure of the ISM and the molecular fraction are therefore not correct in detail. Note, however, that our pressurization prescription is designed to scale with the resolution, and should converge to the ``correct'' result as the resolution improves. In any event, the simulations likely do not include all the relevant physics shaping density and temperature PDFs of the ISM in real galaxies. The results may of course depend on other microphysics of the ISM as they influence both the temperature PDF and the fraction of gas in a high-density, molecular form. Future simulations of the molecular ISM may need to account for new microphysics as they resolve scales where such processes become important. Using hydrodynamical simulations of the ISM and star formation in cosmologically motivated disk galaxies over a range of representative masses, we examine the connection between molecular hydrogen abundance and destruction, observed star formation relations, and the thermodynamical structure of the interstellar medium. Our simulations provide a variety of new insights into the mass dependence of star formation efficiency in galaxies. A summary of our methodology and results follows. \begin{itemize} \item[1.] A model of heating and cooling processes in the interstellar medium (ISM), including low-temperature coolants, dust heating and cooling processes, and heating by the cosmic UV background, cosmic rays, and the local interstellar radiation field (ISRF), is calculated using the photoionization code Cloudy \citep{ferland1998a}. Calculating the molecular fraction of the ISM enables us to implement a prescription for the star formation rate (SFR) that ties the SFR directly to the molecular gas density. The ISM and star formation model is implemented in the SPH/N-body code GADGET2 \citep{springel2005c} and used to simulate the evolution of isolated disk galaxies. \item[2.] We study the correlations between gas surface density ($\Sigmagas$), molecular gas surface density ($\SigmaHmol$), and SFR surface density ($\SigmaSFR$). We find that in our most realistic model that includes heating and destruction of $\Hmol$ by the interstellar radiation field, the power law index of the SK relation, $\SigmaSFR\propto\Sigmagas^{\ntot}$, (measured in annuli) varies from $\ntot\sim2$ in massive galaxies to $\ntot\gtrsim4$ in small mass dwarfs. The corresponding slope of the $\SigmaSFR\propto\SigmaHmol^{\nmol}$ molecular-gas Schmidt-Kennicutt relation is approximately the same for all galaxies, with $\nmol\approx 1.3$. These results are consistent with observations of star formation in different galaxies \citep[e.g.,][]{kennicutt1998a,wong2002a,boissier2003a,heyer2004a,boissier2007a,kennicutt2007a}. \item[3.] In our models, the SFR density scales as $\SigmaSFR \propto \fHmol h_{\SFR} h_{\gas}^{-1.5} \Sigmagas^{1.5}$, where $h_{\gas}$ is the scale-height of the ISM and $h_{\SFR}$ is the scale-height of star-forming gas. The different $\SigmaSFR-\Sigmagas$ relations in galaxies of different mass and in regions of different surface density in our models therefore owe to the dependence of molecular fraction $\fHmol$ and scale height of gas on the gas surface density. \item[4.] We show that the $\SigmaSFR\propto\Sigmagas\Omega$ and $\SigmaSFR\propto\SigmaHmol\Omega$ correlations describe the simulations results well where the molecular gas and total gas densities are comparable, while the simulations deviate from $\SigmaSFR\propto\Sigmagas\Omega$ \citep[e.g.,][]{kennicutt1998a} at low $\Sigmag$ owing to a declining molecular fraction. We demonstrate that these relations may owe to the fact that the angular frequency and the disk-plane gas density are generally related as $\Omega\propto\sqrt{\rho}$ for exponential disks if the potential is dominated by either the disk, a \cite{navarro1996a} halo, \cite{hernquist1990a} halo, or an isothermal sphere. The correlation of $\SigmaSFR$ with $\Omega$ is thus a secondary correlation in the sense that $\Omega\propto\sqrt{\rho}$ is set during galaxy formation and $\Omega$ does not directly influence star formation. \item[5.] The role of critical surface densities for shear instabilities ($\SigmaA$) and \cite{toomre1964a} instabilities ($\SigmaQ$) in star formation \citep[e.g.,][]{martin2001a} is examined in the context of the presented simulations. We find that the two-component Toomre instability criterion $\Qsg<1$ is an accurate indicator of the star-forming regions of disks, and that gravitational instability and star formation are closely related in our simulations. Further, the $\Qsg$ criterion works even in simulations in which cooling is restricted to $T>10^4$~K where gravithermal instability cannot operate. \item[6.] Our simulations that include $\Hmol$-destruction by an ISRF naturally reproduce the observed scaling $\fHmol\propto\Pext^{0.9}$ between molecular fraction and external pressure \citep[e.g.,][]{wong2002a,blitz2004a,blitz2006a}, but we find that simulations without an ISRF have a weaker scaling $\fHmol\propto \Pext^{0.4}$. We calculate how the connection between the scalings of the gas surface density, the stellar surface density, and the ISRF strength influence the $\fHmol-\Pext$ relation in the ISM, and show how the simulated scalings reproduce the $\fHmol-\Pext$ relation even as the power-law index of the total gas Schmidt-Kennicutt relation varies dramatically from galaxy to galaxy. \item[7.] We present a method for mitigating numerical Jeans fragmentation in Smoothed Particle Hydrodynamics simulations that uses a density-dependent pressurization of gas on small scales to ensure that the Jeans mass is properly resolved, similar to techniques used in grid-based simulations \citep[e.g.,][]{truelove1997a,machacek2001a}. The gas internal energy $u$ at the Jeans scale is scaled as $u\propto\mJeans^{-2/3}$, where $\mJeans$ is the local Jeans mass, to ensure the Jeans mass is resolved by some $\NJeans$ number of SPH kernel masses $2\Nneigh m_{\mathrm{SPH}}$, where $\Nneigh$ is the number of SPH neighbor particles and $m_{\mathrm{SPH}}$ is the gas particle mass. For the simulations presented here, we find the \cite{bate1997a} criterion of $\NJeans=1$ to be insufficient to avoid numerical fragmentation and that $\NJeans\sim15$ provides sufficient stability against numerical fragmentation over the time evolution of our simulations. Other simulations may have more stringent resolution requirements \citep[e.g.,][]{commercon2008a}. We also demonstrate that isothermal galactic disks with temperatures of $T=10^{4}\K$ may be susceptible to numerical Jeans instabilities at resolutions common in cosmological simulations of disk galaxy formation, and connect this numerical effect to possible angular momentum deficiencies in cosmologically simulated disk galaxies. \end{itemize} The results of our study indicate that star formation may deviate significantly from the relations commonly assumed in models of galaxy formation in some regimes and that these deviations can be important for the overall galaxy evolution. Our findings provide strong motivation for exploring the consequences of such deviations and for developing further improvements in the treatment of star formation in galaxy formation simulations.\\[2mm] | 7 | 10 | 0710.2102 |
0710 | 0710.2428_arXiv.txt | { A recent analysis of cosmic-ray data from a space borne experiment by the AMS collaboration supports the observation of an excess in the cosmic-ray positron spectrum by previous balloon experiments. The combination of the various experimental data establishes a deviation from the expected background with a significance of more than four standard deviations. The observed change in the spectral index cannot be explained without introducing a new source of positrons. When interpreted within the MSSM a consistent description of the antiproton spectrum, the diffuse gamma-ray flux and the positron fraction is obtained which is compatible with all other experimental data, including recent WMAP data. \PACS{ {98.70.Sa}{Cosmic rays} \and {95.35.+d}{Dark Matter} \and {11.30.Pb}{Supersymmetry} } % } % | Among the cosmic-ray species, antiparticles and diffuse $\gamma$-rays are of particular interest because they are produced secondarily in hadronic interactions of protons and nuclei with the interstellar medium at low rates. Their small abundance makes them a sensitive probe for the existence of additional -- and possibly exotic -- cosmic-ray sources which would be visible as an excess of particles above conventional expectations. One of the most important unsolved questions in modern cosmology is the nature of dark matter. The most promising dark matter candidate is the weakly interacting lightest neutralino, $\chi_1^0$, predicted by supersymmetric extensions to the standard model of particle physics. The annihilation of neutralinos might constitute an additional primary source of particles with a unique spectral shape which would be determined by the parameters of supersymmetry, allowing to put constraints on new physics beyond the standard model. A recent reanalysis of the data from the \AMS{} spectrometer~\cite{aguilar07a} supports the observation of an excess of cosmic-ray positrons by the HEAT experiments~\cite{beatty04a}. In this work, we discuss the combined results on the cosmic-ray positron fraction $e^+ / (e^+ + e^-)$. Assuming that dark matter is largely constituted by neutralinos, we determine the cosmic-ray preferred parameter space of the minimal supersymmetric standard model (MSSM) from a simultaneous fit to the cosmic-ray positron, antiproton and diffuse $\gamma$-ray data. | \label{sec:conclusions} In this work, the combined recent experimental results on the cosmic-ray positron fraction have been presented. The data exhibit an excess of positrons above energies of 6\+\GeV{} which cannot be explained by purely secondary positron production alone and thus requires an additional primary source of positrons. In this work, we interpret this source to be the annihilation of supersymmetric neutralinos constituting dark matter. A simultaneous fit to the cosmic-ray positron, antiproton and $\gamma$-ray data shows that, for particular sets of the MSSM parameters, this hypothesis gives a fully consistent description of the cosmic-ray spectra which is compatible with all other experimental data. We find that the cosmic-ray data clearly prefer the focus point region of the MSSM parameter space but reveal almost no sensitivity to $\tan\beta$. | 7 | 10 | 0710.2428 |
0710 | 0710.5452_arXiv.txt | We use a series of ray-tracing experiments to determine the magnification distribution of high-redshift sources by gravitational lensing. We determine empirically the relation between magnification and redshift, for various cosmological models. We then use this relation to estimate the effect of lensing on the determination of the cosmological parameters from observations of high-$z$ supernovae. We found that, for supernovae at redshifts $z<1.8$, the effect of lensing is negligible compared to the intrinsic uncertainty in the measurements. Using mock data in the range $1.8<z<8$, we show that the effect of lensing can become significant. Hence, if a population of very-high-$z$ supernovae was ever discovered, it would be crucial to fully understand the effect of lensing, before these SNe could be used to constrain cosmological models. We show that the distance moduli $m-M$ for an open CDM universe and a $\Lambda$CDM universe are comparable at $z>2$. Therefore if supernovae up to these redshifts were ever discovered, it is still the ones in the range $0.3<z<1$ that would distinguish these two models. | High-redshift supernovae have become a major tool in modern cosmology. By measuring their apparent magnitudes, we can estimate their luminosity distances $d_L$ (see \citealt{tonryetal03,barrisetal04,riessetal04}, and references therein). Since the relationship between $d_L$ and the redshift $z$ depends on the cosmological parameters, observations of distant SNe can constrain the cosmological model. Prior to the announcement of the {\sl WMAP} results \citep{bennettetal03}, observations of high-$z$ SNe provided the most compelling evidence of the existence of a nonzero cosmological constant. Since then, they have been used in combination with the {\sl WMAP} data to refine the determination of the cosmological parameters. The luminosity distances $d_L$ are determined by combining the observed fluxes $F$ with estimates of the SNe luminosities $L$. Uncertainties in $d_L$ are caused by uncertainties in $L$, because SNe are not perfect standard candles. The flux $F$ is much easier to measure, but for distant sources the value of $F$ might be altered by gravitational lensing caused by the intervening distribution of matter. For instance, a magnification $\mu>1$ would result in a increase in $F$, and an underestimation of $d_L$. Estimating the effect of lensing on the statistics of high-$z$ supernovae is a complex problem. Using either an analytical model or ray-tracing simulations, we can estimate the effect of lensing of a large number of sources in a statistical sense. We would then need to redo the error analysis on the SNe data to include in a consistent way the effect of lensing. This would be a very complex task, and in this paper we have chosen a much simpler approach. {\it Our goal is not to obtain a precise estimate of the error introduced by lensing, but rather to assess the importance of this effect: is it dominant, important, or negligible, and for what range of redshift? and how does it affect the discrimination between different cosmological models?\/} To answer these questions, we take at face value the published results of Type~Ia SNe, including their error bars which account for every source of uncertainty but gravitational lensing. Then, we include {\it a posteriori\/} the effect of lensing, to estimate the change in the errors. This approach is not rigorous at all, and does not constitute a substitute for a rigorous treatment of the errors. But it has the great advantage of simplicity. We do not have to redo the detailed error analysis performed by the high-redshift SNe groups, and, more importantly, our conclusions will not be tied to any particular sample or particular data reduction and error analysis technique used by any particular group. We are seeking to make generic statements about the importance of lensing (or lack of) that are relevant to any current or future sample of high-$z$ SNe. The lensing of distant supernovae has been the focus of several recent studies. In an early study, \citet{wambsganssetal97} used ray-tracing experiments to estimate the effect of weak lensing on the determination of the deceleration parameter $q_0$. \citet{md05}, \citet{dv05}, and \citet{mv06} focused on SNe as a mean to study the nature of weak lensing. The issue of determining the cosmological parameters for distant SNe, and how this determination is affected by lensing, was addressed by \citet{wang05} who used semi-analytical models to determine the magnification distribution function, \citet{hl05} who used Monte Carlo ray-tracing simulations to study the effect of weak and strong lensing, and \citet{gunnarssonetal06} and \citet{jonssonetal06}, who estimated the effect on lensing along individual lines of sight by considering the properties of foreground galaxies in the same direction. These various studies concluded that the effect of lensing on current determinations of the cosmological parameters is small. \citet{alderingetal06} discussed the effect of gravitational lensing on a population of SNe at $z>1.7$. What distinguishes our approach is mostly its simplicity. Our calculations depend on very few assumptions, and this implies a certain amount of robustness to our results. Even though we rely on numerical simulations, this work should be regarded as a back-of-the-envelope calculation, whose purpose is to obtain a qualitative estimate of the effect of lensing on the determination of cosmological parameters by distant SNe. Using ray-tracing experiments, rather than a semi-analytical approach, enables us to extend our study to redshifts much higher than the ones considered by \citet{wang05} and \citet{hl05}. This paper is organized as follow: In \S2, we describe our calculation of the magnification distribution $P(\mu)$, and how to estimate that distribution at any redshift $z$. In \S3, we describe the real and mock samples of supernovae we use for our calculations. Results are presented in \S4. In \S5, we address various observational issues. Summary and conclusion are presented in \S6. | We have performed a series of ray-tracing experiments using a multiple lens-plane algorithm. We have determined the distributions of magnifications $P(\mu)$ for sources in the redshift range $0<z<8$, for three different cosmological models. We have used these distributions to estimate the effect of gravitational lensing on the determination of the cosmological parameters with high-redshift Type~Ia supernovae. We used a generic, {\it a posteriori} approach which is not tied to any particular sample. We found that errors introduced by lensing are unimportant for SNe with redshift $z<1.8$. These errors are negligible compared to the intrinsic errors already present in the supernovae data. Since those intrinsic errors do not prevent us from determining the cosmological parameters, the additional errors introduced by lensing have no consequences. A similar conclusion was reached by \citet{alderingetal06}. Using a mock catalog of high-$z$ SNe, extending to $z=8.1$, we showed that the effect of lensing on a hypothetical population of SNe at redshifts $z>2$ could be quite significant, and must be understood before such SNe could be used to constrain cosmological models. Furthermore, the open CDM and $\Lambda$CDM are difficult to distinguish at that redshift. We showed that, even if SNe at redshift $z\sim8$ were ever discovered, it is the SNe in the range $z=0.3-1$ that would still provide the best discriminant between these two models. The data at that redshift already exist, and they support the $\Lambda$CDM model. | 7 | 10 | 0710.5452 |
0710 | 0710.3369_arXiv.txt | It is thought that the decay of the magnetic field in \emph{magnetars} is the main source of its X-ray and $\gamma$-ray luminosity since these objects appear to radiate substantially more power than available from the energy rotational loss \cite{TD-96}. The spontaneous decay of the magnetic field could occur through \emph{ambipolar diffusion}, \emph{Hall drift} and \emph{ohmic diffusion}, which are non-ideal magnetohydrodynamics processes that occur in thousands of years, compared to dynamical sound and Alfv\'en time scales of milliseconds to seconds. Ambipolar diffusion promote a dissipative magnetic field advection, through the movement of a bulk of charged particles relative to the neutrons, the \emph{Hall drift} is a non-dissipative advection of the magnetic field caused by the electrical current associated with it, and the magnetic field ohmic diffusion is a dissipative process caused by the electrical resistivity. The time scales of these process were estimated in Ref.~2. For classical pulsar magnetic field strengths, these were found to be longer than the lifetimes of these stars, and therefore unlikely to be observationally important. For the case of magnetars, it was found that magnetic field decay by these processes might be occurring \cite{TD-96,ACT-04}. However, a full understanding of these processes, their interactions, and their effectiveness in neutron stars is still lacking. The work of \cite{GR-92} was analytical, and therefore useful to identify general processes and identify the relevant time scales, but not to address the action of the identified processes in their full nonlinear development, and their interaction with each other. The full evolution of the magnetic field can only be addressed by numerical simulations. Recent ideal three-dimensional single-fluid magnetohydrodynamics simulations showed that, in a stably stratified star, a complicated, random, initial field generally evolves on a short, Alfv\'en-like time scale to a relatively simple, roughly axisymmetric, large-scale configuration containing a toroidal and a poloidal component of comparable strength, both of which are required in order to stabilize each other \cite{BS-04}. It is interesting to study the effects of the non-ideal processes studied in Ref.~2 on the evolution of this configuration. As a first step in this direction, here we simulate the decay of a magnetic field, including the effects studied in Ref.~2, in a system where the magnetic field points in one Cartesian direction but varies only along an orthogonal direction. It is shown that the magnetic field evolves through different quasi-equilibrium states and we estimate the characteristic time scales where these quasi-equilibria occur. | Using numerical simulations in one dimension, we studied the effects of some non-ideal MHD processes on the magnetic field evolution in neutron stars. We found that the system evolves through succesive quasi-equilibria, and we estimated the characteristic time scales on which these quasi-equilibria occur. \begin{center} Acknowledgments \end{center} We acknowledge financial support through FONDECYT postdoctoral project 3060103, Gemini project 32070014, and regular FONDECYT projects 1060644 and 1070854. We also thank the ESO-Chile Mixed Committee. | 7 | 10 | 0710.3369 |
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